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+source,target
+ 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.
+ 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.
+ 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."
+ 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."
+ 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.
+ 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."
+ This last possibility is consistent with the dry merger hypothesis., This last possibility is consistent with the dry merger hypothesis.
+ 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."
+ 1990)., 1990).
+ 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}."
+.. 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."
+ 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}."
+ 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."
+ 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."
+ The JA’ color gradieut is defined as δουν)., The $J - K'$ color gradient is defined as $\Delta (J-K')/log(r)$.
+ According to this definition. ealaxies with blue cores have ACFA)flog(r)samples- 0.," According to this definition, galaxies with blue cores have $\Delta (J-K')/log(r)<0$ ."
+ 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.
+" 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.
+ The oulv carly-type, The only early-type
+"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 $255.01 \pm 0.83$ $\mathrm{rad/m^2}$.
+" 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""."
+ 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.
+ 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).
+ 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."
+" 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 $2.1\pm 0.4$ $\mathrm{rad/m^2}$."
+ Consequently. the interstellar RM is given. by RMin=255.01-2.1252.91€0.92 rad/im-.," Consequently, the interstellar RM is given by $\mathrm{RM_{int}}=255.01-2.1=252.91 \pm 0.92$ $\mathrm{rad/m^2}$."
+ 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}.
+ 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."
+ 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}."
+ The latter two quantities are depicted in figure 3.., The latter two quantities are depicted in figure \ref{fig:linpolfracs}.
+ 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.
+ 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}. ."
+ 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 $10\degr$."
+ They are depicted in figure 4.., They are depicted in figure \ref{fig:polangles}.
+ 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."
+ 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}.
+ 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."
+ 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}.
+ 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.
+ The GMRT data were corrected for this (??)..," The GMRT data were corrected for this \citep{Chandraa,Chandrab}."
+ 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.
+ 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}.
+ 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."
+" 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.
+ 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.
+ 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."
+ 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."
+ 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."
+ 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}."
+ 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}.
+ 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.
+ It 1s unknown how these variations translate to lower frequencies., It is unknown how these variations translate to lower frequencies.
+ 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.
+ 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."
+ 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}."
+ 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."
+ 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).
+ 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."
+ 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."
+ 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"
+Iu 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)."
+ 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."
+ 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."
+ 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.
+ 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 $(B-V)>1.2$."
+ 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.
+ 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."
+ 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 $(V-I)$."
+ 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.
+ 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."
+ 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."
+ 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 $<V<21$ )."
+ 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.
+ 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."
+ Further invesieatiou on the nature of these objects ποστς wortlavhile., Further investigation on the nature of these objects seems worthwhile.
+ 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.
+ 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."
+ 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."
+ 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.
+ 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.
+.. No corrections were lade to he Arnouts (1997) data., No corrections were made to the Arnouts (1997) data.
+ 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.
+ 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 $I\sim22.5$."
+ 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.
+ 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, $w(\theta)$."
+ 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)."
+ Departures from wuiformity should affect the correlation function especially at faint magnuitudes., Departures from uniformity should affect the correlation function especially at faint magnitudes.
+ 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"
+This 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 $2\Rsun$."
+ 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}."
+. 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."
+ 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.
+ 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."
+ This CALE has been previously been analvsed by several researchers., This CME has been previously been analysed by several researchers.
+ 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."
+ 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}."
+. 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}."
+ 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."
+ 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."
+ 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.
+the possible eruptiou survivor.,the possible eruption survivor.
+" 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."
+ 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."
+ 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. ("
+2001) 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.
+ 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."
+ 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)."
+ Stockdale et al., Stockdale et al.
+ 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. ("
+However. 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."
+ 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.
+" 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."
+ 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."
+ 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."
+ 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."
+ These bandpasses roughly correspond to V. 2. aud 7.," These bandpasses roughly correspond to $V$, $B$, and $I$."
+ Using the radio position of SN 1961V from Stockdale et al. (, Using the radio position of SN 1961V from Stockdale et al. (
+2001: it is coincident with the position measured by Cowan ct al.,2001; it is coincident with the position measured by Cowan et al.
+ 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."
+ 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."
+" 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."
+ 19995)., 1999b).
+" 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."
+ The source numbering is from Figure 1. as in F95.," The source numbering is from Figure 1, as in F95."
+ 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."
+ 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)."
+" 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 ."
+ 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.,"
+ automatically within one package. (, automatically within one package. (
+We 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.)
+ 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.
+ 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."
+ 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."
+ 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."
+ 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."
+ 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."
+ 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."
+ 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."
+ 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."
+ 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."
+ 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."
+ 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.
+ 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)."
+ 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.
+ As, As
+Using 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)."
+ 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 $B=0.78 \pm0.15$."
+"s The mean wean about he best fitting X moclelis στου£.)""VL."," The mean deviation about the best fitting BCES model is $\sigma({\rm log}\,P_{\rm Bondi})=0.11$."
+ 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.
+ 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 $P_{\rm Bondi}$."
+ 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."
+ 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 $P_{\rm jet}$.
+A power-law mocel provides a good description of the data., A power-law model provides a good description of the data.
+ 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."
+ 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.
+ However. PosiXdiUL 7. whereas Dy.(Oxdio.LA c," However, $P_{\rm Bondi}\propto d_{\rm L}^{-0.5}$ , whereas $P_{\rm jet}\propto d_{\rm L}^{1.5}$."
+pThus. the distance. dependences cannot lead to the observed. positive correlation.," Thus, the distance dependences cannot lead to the observed positive correlation."
+ 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.
+" 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."
+ 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 $P_{\rm
+jet}\propto T^{1.5}$, which cannot lead to the observed positive correlation."
+" 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 $P_{\rm jet}$."
+ 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. ("
+La detail. a mild anti-correlation is observed).,"In detail, a mild anti-correlation is observed)."
+ 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.
+" 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."
+ 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."
+ 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 $P_{\rm Bondi} =
+0.65(\pm0.16) + 0.77(\pm0.20)$ $\,P_{\rm jet}$, where $P_{\rm
+Bondi} = 0.1 {\dot M_{\rm Bondi}} c^2$ and $P_{\rm jet}$ is the power associated with inflating the cavities and providing the internal energy of the plasma that fills them."
+ 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."
+" 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 $P_{\rm
+jet}=10^{42}$ , to $3.7^{+3.3}_{-1.7}$ per cent at $P_{\rm
+jet}=10^{44}$ . ("
+Llere. 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"
+Vhe 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."
+ 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."
+ 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."
+ 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."
+ 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).
+ 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."
+ 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.
+ 19080. van raraclijs et al.," 1990, van Paradijs et al."
+ 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."
+ 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."
+ 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."
+ 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."
+ 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."
+ This, This
+backeround.,background.
+ 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 $\mf/\m12\ll1$."
+" 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."
+" 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}."
+. 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 -."
+pprorl —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"
+eiven bx the neutron decoupling line in the figure.,given by the neutron decoupling line in the figure.
+ 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."
+ These non-thermal reactions are nol described by the Kawano reaction network., These non-thermal reactions are not described by the Kawano reaction network.
+ 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.
+ 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."
+ 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."
+ 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.
+ 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."
+ 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."
+ Anv deuterons present when neutron decoupling occurs will be destroved., Any deuterons present when neutron decoupling occurs will be destroyed.
+ Likewise. anv a particles present will be broken apart.," Likewise, any $\alpha$ particles present will be broken apart."
+ 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 $\alpha+({\rm \sim
+GeV\,\,\,nucleon}) \rightarrow {^2}{\rm H}+X$ 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."
+ 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."
+" 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."
+ 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}."
+ 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 $\hh$."
+ 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.
+ 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."
+ 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 $\sim10\%$.
+ 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.
+ 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.
+ 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.. In
+1996).,.
+. The SBF technique is based ou (he spectral analvsis of an image., The SBF technique is based on the spectral analysis of an image.
+ 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. “"
+Weighted and cosmic-ray cleaned” images have been used instead.,Weighted and cosmic-ray cleaned” images have been used instead.
+ 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."
+ Only dark exposures have been selected., Only dark exposures have been selected.
+ 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.
+ 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}."
+ 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."
+ This technique was introduced. with the aim of measuring distances., This technique was introduced with the aim of measuring distances.
+ 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}."
+. 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."
+ 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}."
+. 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."
+ 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."
+ 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.
+ 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 $P_1$ )."
+ET) 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.
+ 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:"
+"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 $[\mathrm{Fe} / \mathrm{H}] \leq -1.5$."
+" 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."
+ 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)."
+" 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."
+ 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.
+" ? 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."
+" 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."
+" 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."
+ 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.
+ 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.
+ 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.
+" 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 $[\mathrm{Fe} / \mathrm{H}]$."
+ It is however detectable in giants down to [Fe/H]~—2.5., It is however detectable in giants down to $[\mathrm{Fe} / \mathrm{H}] \sim -2.5$.
+ 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.
+" 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."
+ ? 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}.
+ ? 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 $[\mathrm{Fe} / \mathrm{H}]$.
+" 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' $[\mathrm{S} / \mathrm{Fe}]$."
+ Obviously there is no consensus on the Galactic chemical evolution of sulphur., Obviously there is no consensus on the Galactic chemical evolution of sulphur.
+ 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?
+ 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.
+ 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.
+ 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.
+ 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.
+ 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."
+ 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}) ).
+ 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}.
+ 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.
+" 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."
+ They were carried out in service mode during October 2007 - February 2008., They were carried out in service mode during October 2007 - February 2008.
+ 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.
+ 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.
+ 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.
+" 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."
+" 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}."
+ 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.
+ 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}. .
+" 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}, ,"
+and obtain a very precise tomography of stellar interiors.,and obtain a very precise tomography of stellar interiors.
+further 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."
+" 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 $\delta v_{rms}={(\delta v_{N,ATM}^2+\delta
+v_{rms,N}^2)^{1/2}}$, 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 $\delta
+v_{rms}=\delta v_{rms,S}$, and the RV uncertainty is limited by stellar photon noise; 3, for an intermediate situation, if $\delta v_{rms,S}$ and $\delta
+v_{rms,N}$ are identical, then we apply the same weight on both RV measurements, $\delta v_{rms}={[\delta v_{rms,S}+{(\delta
+v_{N,ATM}^2+\delta v_{rms,N}^2)^{1/2}}]/2}$."
+ 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."
+ 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."
+ The study results are sumiuarized below., The study results are summarized below.
+ '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."
+ 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.
+ 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."
+ 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."
+ 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."
+ 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."
+ 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}. ."
+"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}) )."
+ 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}) )."
+ 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."
+ 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."
+ ? 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.
+ However. the stability of," However, the stability of"
+When 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."
+ First. it can induce motions in large amounts of material.," First, it can induce motions in large amounts of material."
+ 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."
+ 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."
+ 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."
+ The second strong οσο is a change of the ICM temperature., The second strong effect is a change of the ICM temperature.
+ 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."
+ 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)."
+ Clearly. the kinematics and the thermocdvnamics. are coupled.," Clearly, the kinematics and the thermodynamics are coupled."
+ 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."
+ 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."
+ 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."
+ 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.
+ 47., 4–7.
+ 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.
+ 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."
+ 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."
+ “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.
+ ‘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."
+ 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."
+ 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."
+" 1 of Reynolds. Ποια, Begelman 2002)."," 1 of Reynolds, Heinz Begelman 2002)."
+ 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."
+ 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."
+ 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."
+ 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."
+ 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."
+ 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."
+ 7)., 7).
+ 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.
+ 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.
+ 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.
+ 2005a.b).," 2005a,b)."
+ 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 $CZ=9000\kmps$ ).
+ 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 $0.080\pm
+0.021\eV$ $3.0\pm 0.8 m\AA$ ) and $0.059\pm 0.021\eV$ $2.2\pm 0.8
+m\AA$ )."
+ Nicastro οἱ al. (, Nicastro et al. (
+2005b) 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."
+ 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."
+ Fang et al. (, Fang et al. (
+2002) 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."
+ Melxernan et al. (, McKernan et al. (
+2003). 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 $cz=9900\kmps$ ).
+ 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.
+ 2005) and the racio-loucl quasar 11821|643 (Mathur. Weinberg Chen 2003)," 2005) and the radio-loud quasar H1821+643 (Mathur, Weinberg Chen 2003)."
+ While observational biases are evident. ib is striking," While observational biases are evident, it is striking"
+is then given by eliony=(3/4).lo.,is then given by $A_{10}n_1=(3/4)A_{10}$.
+ 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}) ).
+ 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 $^{-1}$ , $F=\int S(\nu)\,d\nu=\int
+S(\nu)\,(\nu/c)\,dv$ , 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."
+ 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}) )."
+ 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 $\theta$.
+ 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."
+ 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}) ),"
+" 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 $\sigma_v=(k_{\rm B}\,T/m_h)^{1/2}$ ."
+ 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.
+ 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.
+ 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."
+ We use these observers to compute error bars on the mock observations., We use these observers to compute error bars on the mock observations.
+ 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."
+ 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.
+ 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}) )."
+ 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.
+ 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."
+ At low galactic latitudes this smoothing angle is typically less than VY., At low galactic latitudes this smoothing angle is typically less than $1^\circ$.
+ 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."
+ 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 $T_B\sim
+100$ K) at high velocities, $|v_{\rm lsr}| > 100$ km $^{-1}$, with respect to the local standard of rest (LSR)."
+ 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."
+ 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."
+ 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.
+ 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."
+ 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."
+ 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}) ).
+ 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}. ."
+ 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."
+ 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."
+ 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, $|v_{{\rm
+lsr}}|>100 {\rm \,km\,s}^{-1}$ , is shown in Fig. \ref{fig:hist}."
+ 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."
+ 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.
+ Phe mock and real cata look remarkably similar., The mock and real data look remarkably similar.
+ 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."
+ 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."
+ 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.
+ 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}) )."
+ 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.
+index pairs containing at least one ( were obtained from CLTHfOBIVEVEReEB.MVUB de}.,"index pairs containing at least one $U$ were obtained from ${\rm CI}_i\vert_{i=1}^{i=10}=
+\{U-2B+V,U-V,U-R_C,U-I_C,B-V,B-R_C,B-I_C,V-R_C,V-I_C,R_C-I_C\}$ ."
+ 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."
+ 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."
+ 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}.
+ 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 $\Delta\log g_{\rm e}=\pm 0.04$.
+ These values are indicated by the dotted horizontal lines., These values are indicated by the dotted horizontal lines.
+ 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.
+ 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 $UBV(RI)_C$ ."
+The 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}."
+. 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 $\Delta
+T_{\rm e}(\varphi),\Delta\log g_{\rm e}(\varphi)$ resulted in $[M]=-1.60\pm 0.10$ , $E(B-V)=0.015\pm 0.01$."
+ 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."
+ 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."
+ 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 $\varphi=0.5$.
+ At y=0.98. the distribution is almost uniform.," At $\varphi=0.98$, the distribution is almost uniform."
+ 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 $\{T_{\rm e}^{(i)}(\log g_{\rm e},{\rm CI}_1,{\rm CI}_2,[M],E(B-V))\}_
+{i=1,2}$, i.e. 14 pairs of $\log g_{\rm e},T_{\rm e}$ were found instead of 30 pairs."
+ 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}."
+. 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."
+⊔≱∖↥∪∣⋈⋅↥⇂↥⋖⋅↓≻∐∥⊳∖∢⊾∪⇂⋅↿↓∐⋅↓⋯⊔↓↓≻∪⊔↿↓↕∢⋅↓⊀↓⋏∙≟↓∐≼∼⊔↓⋅∖⇁∢⊾∖∖⋎↓↕⋖⋅⊔ 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}.
+ 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.
+ ‘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.
+ 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."
+ 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 $\varphi \approx 0.93$."
+ To obtain ο. £). ete.," To obtain $v(r,t)$ , etc.,"
+ 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."
+ Fig. 3((, Fig. \ref{fig3}( (
+a) 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 $a^{\rm (dyn)}(R,t)=0$."
+" 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 $[\ast]$."
+At 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 $a^{\rm (dyn)}(R,t)/g_{\rm
+ s}(R,t)\approx -0.07$; therefore, this phase was included to obtain $d,{\cal M}$ of $[\ast]$ ."
+ 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 $a^{\rm
+ (dyn)}(R,\varphi=0.2,0.5)=5.8,-3.3\:\mbox{ms}^{-2}$ and $q=-0.52,38$ , respectively."
+ 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}( ("
+b)-(d).,b)-(d).
+ 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 $\varphi=0.35\mbox{-}0.45$ ."
+ 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}))."
+ |g)&O.L is expected. from Alogg.= 0.04.," $\vert q\vert\approx 0.1$ is expected from $\Delta \log g_{\rm e}=0.04$ ,"
+regions: in cach case the Lree-Lree emission is dominated byIli.,; in each case the free-free emission is dominated by.
+ 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."
+ 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}. ."
+ 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}."
+. 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."
+ 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}."
+". Summing over the aanel rregions.. qj,coll—54510wous1 M, "," Summing over the and regions, $q_\lya^{\mathrm
+coll}=3.4\times10^{47}~\se^{-1}\,\mathrm{M^{-1}_{\odot}}$ ."
+lor 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)."
+ 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 $1216~\ang$."
+" 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."
+ The free-free spectrum is given by eq. (19))., The free-free spectrum is given by eq. \ref{ffspeceq}) ).
+ 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.
+ Phat rate is where du;= Qgi/M. and quí;=Quo/Ale.," That rate is where $q_{\mathrm HI}\equiv Q_{\mathrm HI}/M_*$ and $q_{\mathrm
+HeII}\equiv Q_{\mathrm HeII}/M_*$."
+ 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."
+ 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 $z=15$."
+ 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."
+ section 3.3)), section \ref{allescsec}) ).
+ 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, $f_{\mathrm
+esc}=1$."
+ 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)."
+ 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 $z$."
+ 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 $n_{\mathrm IGM}(z)=
+1.7\times10^{-7}(1+z)^3~\cmt$ (which ignores the small fraction of baryons in collapsed haloes), $X=0.75$ and $Y=0.25$."
+ 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.
+ 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, $\tau \simeq 2\times10^6~\yr$ ."
+ 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."
+ 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"
+"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."
+ The total observing time was 4 hours., The total observing time was 4 hours.
+ 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."
+ 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."
+ 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."
+ 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."
+ 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."
+ 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."
+ 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.
+ 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.
+ The calibration was then transferred to the target source M82., The calibration was then transferred to the target source M82.
+ 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."
+ 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."
+ Each observation was taken with an exposure of about 18 ks., Each observation was taken with an exposure of about 18 ks.
+ 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."
+ 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}
+ 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."
+ 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."
+ 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."
+. 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.
+ 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 $^{-1}$."
+ 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.
+ 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 $^{-1}$ ."
+ The supernova was clearly detected in both epochs (see Fig. 1)., The supernova was clearly detected in both epochs (see Fig. \ref{fig:hsa}) ).
+ 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."
+ In the following eleven months the source expanded and faded significantly., In the following eleven months the source expanded and faded significantly.
+ 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 $\lambda$ .
+ This ensures that most ofthe, This ensures that most ofthe
+low 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."
+ 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."
+ 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.
+ 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."
+ 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)."
+ 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."
+ work (i.e. collisions between eas rich clumps)., work (i.e. collisions between gas rich clumps).
+ For simplicity we assume (hat all clusters are formed in this way., For simplicity we assume that all clusters are formed in this way.
+ 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."
+ 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.
+ 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)."
+ 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.
+ 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."
+ 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."
+ 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."
+ 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)."
+ 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.
+ 8 of Cote (1999) along wilh best fit Gaussians lor each group., 8 of Cote (1999) along with best fit Gaussians for each group.
+ 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."
+ In, In
+of 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."
+ 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."
+ 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}.
+ 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 $z > 2$."
+ 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."
+ 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 $\ln W$."
+ 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}."
+ 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."
+" 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."
+ 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 $\Omega_c$."
+ 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."
+ 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."
+ 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."
+ 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."
+of our model inadequacy criterion in the Appendix.,of our model inadequacy criterion in the Appendix.
+ 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."
+ 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.
+ 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."
+ 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.
+ 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}."
+. 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}."
+. 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."
+ 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."
+ 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."
+ The interpretation of the posterior probabilities in Eq. (1)), The interpretation of the posterior probabilities in Eq. \ref{model_probability}) )
+ 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.
+ 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."
+ 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.
+ 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."
+ 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."
+ With the marginal likelihoods available according to the Eq. 2..," With the marginal likelihoods available according to the Eq. \ref{marginal_likelihood},"
+" 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."
+ 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}) ).
+ 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."
+ The number + has an interpretation as a threshold value., The number $r$ has an interpretation as a threshold value.
+ 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)."
+ Therefore. if the best model according to Eq. (1))," Therefore, if the best model according to Eq. \ref{model_probability}) )"
+ satisfies Eq. (3)), satisfies Eq. \ref{MIC}) )
+ 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."
+ 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."
+ 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."
+ 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.
+ 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.
+ 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.
+ We describe the parameter probability densities using three numbers., We describe the parameter probability densities using three numbers.
+ 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}.
+. 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."
+ 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"
+ , 
+when 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.
+ 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}) ).
+" 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 et
+al. 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); $x_{yi}$: $0.3 \leq t \leq
+0.7$ Gyr), $x_{io}$: $1 \leq t \leq 2$ Gyr) and $x_{o}$: $5 \leq t \leq 13$ Gyr)."
+" In Fig. 5,,"," In Fig. \ref{pop1},"
+" 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."
+" 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 $\approx$ )."
+" 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."
+ 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}.
+". 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."
+ The mass-fraction of each population vector components is show in Fig., The mass-fraction of each population vector components is show in Fig.
+" 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: $m_o$ )."
+ 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.
+" 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."
+" 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 )."
+ 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.
+" 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}."
+. The map for 11157 is shown in Fig., The map for 1157 is shown in Fig.
+ 7 and presents values adevx5 well the observed underlying spectra., \ref{pop2} and presents values $adev\lesssim5$ well the observed underlying spectra.
+" 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)."
+" 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."
+" 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}.
+". 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}."
+ The comparison between both spectra is shown in Fig. 4.., The comparison between both spectra is shown in Fig. \ref{nifsxspex}.
+ It is clear that our NIFS spectra havea reliable flux calibration., It is clear that our NIFS spectra have a reliable flux calibration.
+ 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.
+" 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."
+" 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}. ."
+ 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.
+" 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}."
+" 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}.
+. 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 ).
+" 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 $\approx$ ${\rm km\,s^{-1}}$ )."
+" 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}. ."
+" 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."
+" Interestingly, very similarresults were found by Riffeletal.(2010b) for 11066."," Interestingly, very similarresults were found by \citet{rogemar10b} for 1066."
+" 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"
+as described in LOs.,as described in L08.
+ 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.
+ 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.
+ 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."
+ \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.
+ 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.
+ We used Subaru photometry in the V aud I bauds., We used Subaru photometry in the V and I bands.
+" 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."
+ 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."
+ 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."
+ 1)., 1).
+ 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."
+ 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."
+ 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."
+ 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.
+ 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.
+ 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."
+ 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 $z=0.183$ )."
+ 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.
+ 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."
+ 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."
+ The resulting fit is eood. \?/dof=19.2/16.," The resulting fit is good, $\chi^2/{\rm dof}=19.2/16$ ."
+ 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."
+ 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."
+ 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)."
+ 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 $p=0.74 \pm 0.30$ ."
+ 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.
+ 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)."
+ 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."
+ 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."
+ 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 space
+diagram}) )."
+ 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.
+ Defining membershipis especially important for. massive, Defining membershipis especially important for massive
+bv 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.
+ 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.
+ MB thanks Justin Nieusmia for technical assisance with performing the smunlatious., MR thanks Justin Nieusma for technical assistance with performing the simulations.
+ 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."
+ 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."
+ We iue erateful to this eroup for sluuiug some of thei results., We are grateful to this group for sharing some of their results.
+ SPO acknowledges support by NASA erant. NNCGOGCII9SC:. and NSF οτα 0908[80.," SPO acknowledges support by NASA grant NNG06GH95G, and NSF grant 0908480."
+. 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."
+Since 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."
+ 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).
+ 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.
+ 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."
+ 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)."
+ 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."
+ 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).
+ 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"
+formation. the dust grain growth: would always proceed in a hot environment.,"formation, the dust grain growth would always proceed in a hot environment."
+ Then the resulting planetesimals would be devoid of water. (, Then the resulting planetesimals would be devoid of water. (
+2) 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."
+ Dust particles grow mainly bv collisions., Dust particles grow mainly by collisions.
+ 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."
+ 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."
+ 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."
+ 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. ("
+3) 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."
+ 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."
+ 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.
+ 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.
+ Future work on these points is needed., Future work on these points is needed.
+ 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"
+Galaxy 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}.
+ 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}.
+ 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)."
+ 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."
+ 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}."
+ 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.
+ 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."
+ 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}."
+ 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."
+ 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."
+ 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}
+ 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."
+ 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 $L_{\rm
+ X,bol}=5.32\pm 0.15 \times10^{45}$ ergs/s \citep{Zhang2008}) ) and the host of a strong cool core \citep{Peterson2001}."
+ 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."
+ 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}."
+ 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."
+ 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.
+ 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.
+ We assume Hy= 70km s7!.Q4;=0.3.Q40.7.," We assume $H_0= 70$ km $^{-1}, \Omega_M =0.3,\Omega_\Lambda=0.7$."
+ 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, $\Delta_{\rm
+ virial}(z=0.25) =121$ \citep{Bryan1998}."
+ 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}. ."
+ The 90% completeness flux, The $90\%$ completeness flux
+We 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}.
+ 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.
+" 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 $\simeq +1.53$."
+ 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."
+ 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."
+ 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."
+ 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."
+ 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."
+ 11 aud 12 in €OOLCal.2010233)., 11 and 12 in \citealt{coo10}) ).
+ 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.
+ The normal {(Si.S ratios are consistent with the supernova scenario.," The normal [(Si,S)/Fe] ratios are consistent with the supernova scenario."
+ 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.
+ 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).
+ This is also supported by the observe Ni/Fo] |, This is also supported by the observed [Ni/Fe]. [
+Cr/Fo] is consistent with both superLOVa nodels.,Cr/Fe] is consistent with both supernova models.
+ 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."
+ 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."
+ Therefore. it is unucertain how eficieut the müxiug-fallback is.," Therefore, it is uncertain how efficient the mixing-fallback is."
+ 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."
+ 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."
+ 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}
+ the dotted line is for the nucleosynthesis yields of the PISN of a $170M_\odot$ star, which are taken from \citet{ume02}."
+". 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.
+ 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."
+ 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."
+The 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.
+ 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."
+ 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."
+. 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.
+ 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.
+ 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.
+ IZMC' was supported by NASA through the Chandra Fellowship Program., EMC was supported by NASA through the Chandra Fellowship Program.
+ AITW. PSR and. WSW acknowledge the United. States Ollice of Naval Rescarch.," MTW, PSR and KSW acknowledge the United States Office of Naval Research."
+ aancd acknowledge support from Chandra erant GOS-9045X., and acknowledge support from Chandra grant GO8-9045X.
+"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 $g^*$."
+ 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 $\epsilon = g$."
+" 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."
+ 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. $\kappa <<1$."
+ 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 $\mathbf{\theta}$."
+ 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."
+" 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."
+ 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."
+ 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)."
+ “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 $\approx 2\sigma$."
+ 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.
+ 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."
+ 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."
+" 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."
+" 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 $\mmill$."
+ 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."
+ 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."
+ 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)."
+ Thus. our analysis includes only thelocal environment around the clusters.," Thus, our analysis includes only the environment around the clusters."
+" 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."
+ 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}) )"
+ 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."
+ 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."
+ 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"
+broad emission roffesults)) aud on the interpretation of this emission rofDiscussion)).,broad emission \\ref{Results}) ) and on the interpretation of this emission \\ref{Discussion}) ).
+ 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."
+ The significance of these irieasuremoents are quantified in two wavs., The significance of these measurements are quantified in two ways.
+ 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."
+ 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."
+ 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."
+ 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."
+ 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.
+" 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 )."
+ 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."
+ 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."
+ These results are sununarized in Table , These results are summarized in Table \ref{tab:results}.
+We 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."
+ 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)."
+ 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."
+ 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."
+ 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 $P_{false} =$."
+"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."
+. 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."
+" 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."
+ iu prep)., in prep).
+ 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."
+ 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."
+ 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)."
+ 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."
+ 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."
+ 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."
+ 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}) )."
+ 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.
+ 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."
+ 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."
+ 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. (
+n prep).,in prep).
+ The low luninosity broad cussion seen in our SINS, The low luminosity broad emission seen in our SINS
+strength 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}."
+. 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."
+ 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."
+ 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}."
+.. 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.
+ 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."
+ (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."
+ 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}."
+ 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}."
+ 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}."
+. 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."
+ 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}.
+. 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."
+ 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}."
+. 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."
+ 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."
+ 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.
+ 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.
+ 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."
+ “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}.
+. 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."
+ 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."
+ 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."
+ We also use larger apertures for our calculations., We also use larger apertures for our calculations.
+ 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."
+ 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."
+ The percentage polarisation also remains stable (at ~ 3%)) across the ridge., The percentage polarisation also remains stable (at $\sim$ ) across the ridge.
+ 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."
+ 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."
+ 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."
+ 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."
+ 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."
+ 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."
+ found very Little polarisation. from. dust., found very little polarisation from dust.
+ 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.
+ 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."
+ 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.
+ 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."
+ 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."
+ 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.
+ 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"
+"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."
+ The gravitational potential is computed by solving Poisson's equation., The gravitational potential is computed by solving Poisson's equation.
+" 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 $\delta_{_{DM}}=\rho_{_{DM}}/\rho_{_{B}}-1$."
+" 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.
+ 1996)., 1996).
+ 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."
+ 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.
+ 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."
+ 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.
+" 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."
+" 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.
+ The basic idea is simple., The basic idea is simple.
+ 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.
+ 3.1)., 3.1).
+ 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."
+ 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.
+ 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."
+ 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."
+ 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).
+ 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.
+ 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."
+ 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.
+ 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)"
+With 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."
+ 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."
+ 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."
+ 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). (
+b) 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 $y_{\rm eff} \simeq 0.0045$.
+ 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."
+ 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."
+ 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}."
+. 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] $\sim -0.5$."
+ 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.
+ 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.
+ 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."
+ 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.
+on the plivsical nature of the hieh-velocity feature.,on the physical nature of the high-velocity feature.
+" 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?"
+ 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,
+Howell:99by}."
+ This analvsis will be presented elsewhere., This analysis will be presented elsewhere.
+ Here we will discuss some of the physical possibilities., Here we will discuss some of the physical possibilities.
+ 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."
+ 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."
+ 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}
+ and we will not repeat them here."
+ 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."
+ 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."
+ We will return to the possible effects of an accretion disk below., We will return to the possible effects of an accretion disk below.
+ 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.
+ 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."
+ 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."
+ 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."
+ 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.
+ 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."
+ 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.
+ This is, This is
+a 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.
+ 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."
+ 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."
+ 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."
+ 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}."
+. Sometimes. the bipolar pairs appear one behind another along the same filament.," Sometimes, the bipolar pairs appear one behind another along the same filament."
+ 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}."
+. 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."
+ 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."
+ 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.
+ 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}.
+. 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}. ."
+. 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}."
+. 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.
+ 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.
+ 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."
+ 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"
+With 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}."
+. 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."
+" 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.
+ 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."
+ Fig. 2..," Fig. \ref{fig:mloss},"
+" 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."
+ 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}.
+. 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.
+" 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."
+ 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.
+ 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}."
+. 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."
+ 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."
+" 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."
+" 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."
+ 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}
+ but a seismic test for mixing in"
+rest of the observation.,rest of the observation.
+ The best-fit parameters of both spectra are consistent with each other., The best-fit parameters of both spectra are consistent with each other.
+ 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.
+ 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}.
+ 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.
+ 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 $\chi^2_{\rm red}=4.4$ ).
+ 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. (
+1999) 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."
+ 2008a) are consistent. with the extrapolation to earlier times of this curve., 2008a) are consistent with the extrapolation to earlier times of this curve.
+ 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."
+ 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."
+ 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.
+ 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.
+ 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.
+ 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.
+ 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}) )."
+ 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.
+ 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.
+" 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."
+ 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-law
+covvariance.,variance.
+comparison 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."
+ Iu Fig., In Fig.
+ 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."
+ 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."
+ Why are there no solar analoes with intermediate Li abundances?, Why are there no solar analogs with intermediate Li abundances?
+ 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?
+ 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?
+ 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. (
+"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."
+ 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."
+ 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."
+ 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."
+ 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."
+ 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."
+ 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."
+. 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."
+ 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."
+ 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.
+. 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."
+ Iudeed. as shown by iu their Fig.," Indeed, as shown by in their Fig."
+ 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."
+ 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.
+ 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.
+ 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."
+ 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."
+ 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."
+ 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."
+. 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.
+ 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."
+ 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."
+ 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."
+ Tutercstinely. this," Interestingly, this"
+"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}."
+ugita&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 $\langle \rho_{\rm b} \rangle$."
+ 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}."
+. 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."
+" 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."
+ 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 ${\cal M}_{\ell} = v_{\ell} / c_{\rm s,\,IGM}$."
+" 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}. ."
+" 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 ${\cal M}_{\ell} \sim (20 - 120)$."
+ 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."
+ For this outer cocoon. taking again," For this outer cocoon, taking again"
+Because 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."
+ These red objects preferentially populate clusters., These red objects preferentially populate clusters.
+ 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}) ).
+ We acquired spectra for the objects with the IHLIectospec NAM el al., We acquired spectra for the objects with the Hectospec (Fabricant et al.
+ 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."
+ The Hectospec observation planning software (Roll et al., The Hectospec observation planning software (Roll et al.
+ 1998) efficient acquisition of a pre-selected sample of galaxies., 1998) enables efficient acquisition of a pre-selected sample of galaxies.
+ The software enables assignment of priorities as a function of galaxy. properties., The software enables assignment of priorities as a function of galaxy properties.
+ 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."
+ Exposure times ranged n 0.75 — 1.5 hours., Exposure times ranged from 0.75 — 1.5 hours.
+ We reduced the data with the standard IHectospec pipeline (Mink et al., We reduced the data with the standard Hectospec pipeline (Mink et al.
+" 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.
+ 2005)., 2005).
+ 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.
+ 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 $^{-1}$. ("
+Geller et al.,Geller et al.
+ 2010: Fabricant et al., 2010; Fabricant et al.
+ 2005)., 2005).
+ 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.
+ 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."
+ 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.
+ Table 1 lists all of the objects and their redshilts., Table 1 lists all of the objects and their redshifts.
+" 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)."
+ 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.
+ 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."
+ The overlapping objects are brighter than the tvpical galaxy in our sample., The overlapping objects are brighter than the typical galaxy in our sample.
+ 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 $^{-1}$ .
+ The dispersion around the mean is 36.8 kis f., The dispersion around the mean is 36.8 km $^{-1}$.
+ 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 $g-r > 1.0$."
+ 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.
+ Our target selection vields 3992 objects., Our target selection yields 3992 objects.
+ In (his subsample. we obtained spectra for 3021 galaxies ancl 149 stars.," In this subsample, we obtained spectra for 3021 galaxies and 149 stars."
+ In(he original selection. we include," Inthe original selection, we include"
+observed 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."
+ 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).
+ 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.
+ 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.
+ 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."
+ 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?"
+ 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.
+ 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."
+ 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)."
+ ∐∪∖∖⊽≼↲∖⇁≼↲↕⋅⋅⊔∐↲⋡∖↽∐≀↧↴∖↽≼↲↕↽≻∪∪↕⋅∫∖��−∣↽≻≀↧↴∐≼⇂≼↲∐∐⊳∖⊽⊳∖⊽↥∪∐≼↲∐∎↓≺∢↕≼↲∐≺∢↕≼↲⊳∖⊽↙↽⊋⋮⋔⋝∖ <]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."
+ 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)."
+ 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)."
+ 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."
+ 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).
+ 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)."
+ Aluminum-caleium-silicates may. be sustained to somewhat hieher temperatures., Aluminum-calcium-silicates may be sustained to somewhat higher temperatures.
+ 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.
+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 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."
+ “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 $H(z)$.
+ Inevitably the underlving. model the cosmologist has is incorrect in some wav., Inevitably the underlying model the cosmologist has is incorrect in some way.
+ 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}."
+The 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."
+ 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}."
+ 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."
+ 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}."
+It 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.
+ 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.
+ 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.
+ 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.
+ 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}) )."
+ This characteristic clustering scale. which emerges from physical processes in," This characteristic clustering scale, which emerges from physical processes in"
+changing from ~5% for SNe =5 to <2% for SNe >50.,changing from $\sim 5\%$ for SNe =5 to $< 2\%$ for SNe $>50$.
+ 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."
+ 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.
+ 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.,"
+ in preparation)., in preparation).
+ 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."
+ 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."
+ 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 $c_{\rm vir} = 6$."
+ 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."
+ 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."
+ 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 $^{-2}$ .
+ 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."
+ 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 $\bar{\kappa} (\theta_2, \theta_{\rm out}) \sim 0$."
+ 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 $M \ge 10^{14} M_\odot$."
+ 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."
+ 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."
+ 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}."
+ 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.
+ 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."
+ 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."
+ 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."
+ In the following. we compare the results obtained by different methods.," In the following, we compare the results obtained by different methods."
+ 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)."
+ Finally. we included in our analysis the photometric redshifts.," Finally, we included in our analysis the photometric redshifts."
+ 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}."
+" 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."
+ This quantity should be computed for each lensed galaxy., This quantity should be computed for each lensed galaxy.
+ 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 $ ."
+ 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."
+ 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."
+ 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 $M_{\rm 2D}$ )."
+ In the case, In the case
+especially 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."
+ 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 $r$."
+ 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."
+" 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."
+ 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.
+ 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).
+ 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.
+ 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).
+ 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."
+ 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."
+ 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."
+ 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."
+" 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 $\Gamma (r)$."
+ 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)."
+ The BP9S $S5M D and Α profiles. respectively. ave shown in Figs.," The BP98 SSM $\Gamma$ and $N$ profiles, respectively, are shown in Figs."
+ 2(a.b).," 2(a,b)."
+ 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 $\Gamma (r)$.
+ 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."
+ 3., 3.
+" 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."
+ 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."
+ 3., 3.
+for gamma-ray burst sources (Paezviisski 1991. 2001: Haensel et al.,"for gamma-ray burst sources (Paczyńsski 1991, 2001; Haensel et al."
+ 1991)., 1991).
+ ‘This is reason enough to study such coalescences., This is reason enough to study such coalescences.
+ 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."
+ 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.
+ 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.
+ 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)."
+ 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 $2M_\odot$."
+" 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."
+ For recent reviews see Cheng. Dai and Lu (1998) ancl Madsen (1999).," For recent reviews see Cheng, Dai and Lu (1998) and Madsen (1999)."
+ 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.
+ 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."
+ 1998)., 1998).
+ 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."
+ 1999. Stergioulas et al.," 1999, Stergioulas et al."
+ 1999. Gondek-Itosisska et al.," 1999, Gondek-Rosińsska et al."
+ 2001. Bombaci et a.," 2001, Bombaci et al."
+ 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."
+ SGRs (Alcock. Farhi and Olinto 1986b: Llorvath et al.," SGRs (Alcock, Farhi and Olinto 1986b; Horvath et al."
+ 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."
+ 1999)., 1999).
+ his reference list is far (rom exhaustive., This reference list is far from exhaustive.
+ llowever. there is a shadow over all this activity.," However, there is a shadow over all this activity."
+" The Galaxy must not be contaminated with ""seeds? of quark matter.", The Galaxy must not be contaminated with “seeds” of quark matter.
+ 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."
+ 1986a), 1986a).
+ Jul neutrons are easily absorbed, But neutrons are easily absorbed
+Some 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}.
+. 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}.
+. 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.
+ 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."
+ 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."
+ 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."
+ 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."
+ 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.
+ 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}.
+. 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 $d_\mathrm{min}$.
+ 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}) )
+ 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.
+ In Sec., In Sec.
+" 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."
+ 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.
+ 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."
+ 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.
+ 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.
+ 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 $\Phi \approx 10^{-8}\,(R_{GC}/30\, 
+\mathrm{Gpc}^{-3}\mathrm{yr}^{-1})$ per globular cluster per year (see Sec. \ref{sec:mindist}) )"
+ 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 $\Phi \times 10^7\, \mathrm{years} \approx 0.1\,(R_{GC}/30\, 
+\mathrm{Gpc}^{-3}\mathrm{yr}^{-1})$."
+ 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."
+ 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.
+ 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.
+ 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}."
+. These objects likely exhibit also a higher GRB rate (Grindlayetal..2006:Lee2010).," These objects likely exhibit also a higher GRB rate \citep{grindlay2006,lee2010}."
+ 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."
+ 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."
+ 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.
+" 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.
+ 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.
+ In Sec., In Sec.
+ 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.
+ 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.
+ 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."
+ 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.
+ 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.
+ 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}."
+ 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.
+ 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.
+ With these data it will be possible to, With these data it will be possible to
+The 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}) )
+on VLT/UT2 (ESO. Chile).," on VLT/UT2 (ESO, Chile)."
+ 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."
+ 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 )."
+ 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."
+ 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}."
+ 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."
+ 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).
+ 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."
+ 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)."
+ 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."
+ 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."
+" 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?"
+" for ""doubtful member"")."," for “doubtful member"")."
+ 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})"
+ 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."
+ 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.
+ 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}:"
+ 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."
+ Star 1240 has also been observed by Yong et al. (, Star 1240 has also been observed by Yong et al. (
+their star 8. as seen from their table 2).,"their star 8, as seen from their table 2)."
+ 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."
+ 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."
+ 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."
+ 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.
+ 1994) for the two member stars., 1994) for the two member stars.
+ 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. (
+2008). 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."
+ and Kassis et al., and Kassis et al.
+ systems. respectively.," systems, respectively."
+ 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. ("
+2005) and Carraro et al. (,2005) and Carraro et al. (
+2004: our star 398. their 801).,"2004: our star 398, their 801)."
+ 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 $RV$."
+ The average RV is 24.6640.37 km s!, The average $RV$ is $\pm$ 0.37 km $^{-1}$.
+ 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}) ),"
+ which is also in perfect agreement with Kaluzny (1994))., which is also in perfect agreement with Kaluzny \cite{kalBe29}) ).
+ 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 $^{-1}$ ).
+ 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."
+ 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}) )."
+ Recent studies (e.g. Carretta et al. 2005)), Recent studies (e.g. Carretta et al. \cite{carretta05}) )
+ 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."
+ 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).
+ 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."
+ Star 1614 rotates very rapidly and was discarded from further analysis., Star 1614 rotates very rapidly and was discarded from further analysis.
+ 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}) ).
+ 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."
+ 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."
+ For the Sun we obtain logn(Fe )=7.49+0.04 (standard deviation. or rms) and logz(Fe m=7.54+0.03. adopting the following effective temperature. surface gravity and microturbulence velocity:," For the Sun we obtain $\log n(\rm{Fe}$ $\pm$ 0.04 (standard deviation, or rms) and $\log n(\rm{Fe}$ $\pm$ 0.03, adopting the following effective temperature, surface gravity and microturbulence velocity:"
+between star formation rate and coronal soft X-ray luminosity.,between star formation rate and coronal soft X-ray luminosity.
+ In some particularly spectacular cases. such as M82. the X-ray emission exhibits a biconical morphology suggestive of outflows (2).. and this could also be the case for our own Milky Way €??).," In some particularly spectacular cases, such as M82, the X-ray emission exhibits a biconical morphology suggestive of outflows , and this could also be the case for our own Milky Way ."
+. Since the existence of hot gaseous coronae around sufficiently massive galaxies is a quintessential component of the WF9] model (and thus of the many semi-analytic models based on it). failure to detect the predicted X-ray emission poses an interesting problem for the canonical view of galaxy formation in a CDMuniverse!.," Since the existence of hot gaseous coronae around sufficiently massive galaxies is a quintessential component of the WF91 model (and thus of the many semi-analytic models based on it), failure to detect the predicted X-ray emission poses an interesting problem for the canonical view of galaxy formation in a CDM."
+. A possible solution was recently proposed by CIO., A possible solution was recently proposed by .
+. They showed that disc. star-forming galaxies in the cosmological hydrodynamic simulations of the do. in fact. develop the Kind of gaseous coronae predicted by WF9]. but with an associated X-ray emission that is one to two orders of magnitude fainter than the ΝΕΟΙ prediction.," They showed that disc, star-forming galaxies in the cosmological hydrodynamic simulations of the do, in fact, develop the kind of gaseous coronae predicted by WF91, but with an associated X-ray emission that is one to two orders of magnitude fainter than the WF91 prediction."
+ The main reason for this is that. contrary to the assumption of ΝΕΟΙ. the density profile of the hot gas does not follow the density profile of the dark matter.," The main reason for this is that, contrary to the assumption of WF91, the density profile of the hot gas does not follow the density profile of the dark matter."
+ Instead. it is much less centrally concentrated as a result of energy injection from SNe at the peak of the star formation activity. 2~I—3.," Instead, it is much less centrally concentrated as a result of energy injection from SNe at the peak of the star formation activity, $z\sim1-3$."
+ This raises the entropy of the gas. both by driving outflows that shock the gas to high temperatures and by ejecting low-entropy gas from the progenitor haloes.," This raises the entropy of the gas, both by driving outflows that shock the gas to high temperatures and by ejecting low-entropy gas from the progenitor haloes."
+ CIO showed that although the X-ray emission around dise galaxies in ccomes predominatly from an extended quasi-hydrostatie corona heated by shocks and gravitational contraction. the galaxies reproduce the observed scaling of the soft X-ray luminosity with éy-band luminosity. dise rotation velocity and star formation rate.," C10 showed that although the X-ray emission around disc galaxies in comes predominatly from an extended quasi-hydrostatic corona heated by shocks and gravitational contraction, the galaxies reproduce the observed scaling of the soft X-ray luminosity with $K$ -band luminosity, disc rotation velocity and star formation rate."
+ The presence of the latter correlation — which arises because both the X-ray luminosity and the star formation rate scale with halo mass — contradicts the main source of evidence in favour of the view that the X-ray emission observed around real dise galaxies must be directly associated with star formation (perhaps through a galactic wind or fountain)., The presence of the latter correlation – which arises because both the X-ray luminosity and the star formation rate scale with halo mass – contradicts the main source of evidence in favour of the view that the X-ray emission observed around real disc galaxies must be directly associated with star formation (perhaps through a galactic wind or fountain).
+ The low surface brightness of coronal gas renders it difficult to detect with παπάXAMM-Newton.. making it an ideal target for future facilities such or the UXO)).," The low surface brightness of coronal gas renders it difficult to detect with and, making it an ideal target for future facilities such or the )."
+ Sensitive. high-resolution spectroscopy with these facilities has the potential to verify directly whether coronal gas is quasi-hydrostatic or outflowing.," Sensitive, high-resolution spectroscopy with these facilities has the potential to verify directly whether coronal gas is quasi-hydrostatic or outflowing."
+ In the meantime. it is interesting to explore another consequence of the ΝΕΟΙ framework. namely that the properties of the hot coronal gas should be broadly independent of the morphological characteristics of the galaxy.," In the meantime, it is interesting to explore another consequence of the WF91 framework, namely that the properties of the hot coronal gas should be broadly independent of the morphological characteristics of the galaxy."
+ Thus. just as bright dise galaxies. bright elliptical galaxies are also expected to have extended coronae of X-ray emitting hot gas.," Thus, just as bright disc galaxies, bright elliptical galaxies are also expected to have extended coronae of X-ray emitting hot gas."
+ Furthermore. since ellipticals are dominated by old stellar populations and have little ongoing star formation. the view that most of the X-ray emission comes from outflowing winds triggered by SNe-II ean be readily ruled out2).," Furthermore, since ellipticals are dominated by old stellar populations and have little ongoing star formation, the view that most of the X-ray emission comes from outflowing winds triggered by SNe-II can be readily ruled out."
+". Active galactic nuclei (AGN) are not thought to be important in the production of X-ray emission from L, ellipticals(2).", Active galactic nuclei (AGN) are not thought to be important in the production of X-ray emission from $L_\star$ ellipticals.
+". In this paper we explore the properties of diffuse X-ray emission in L, galaxies of different optical morphologies.", In this paper we explore the properties of diffuse X-ray emission in $L_\star$ galaxies of different optical morphologies.
+" Recent data from aand pprovide relatively large samples of X-ray luminosities and temperatures (or upper limits) for normal nnon-interacting. relatively isolated) L, dise and elliptical galaxies."," Recent data from and provide relatively large samples of X-ray luminosities and temperatures (or upper limits) for normal non-interacting, relatively isolated) $L_\star$ disc and elliptical galaxies."
+ The observational samples that we analyse in this paper are introduced ins d, The observational samples that we analyse in this paper are introduced in \\ref{sec:xray_data}.
+ nS refsce:scalings. wepresentcomparisonsof derived rayscalings fordisesandellipticals," In \\ref{sec:scalings}, we present comparisons of derived X-ray scalings for discs and ellipticals."
+ Ln scetinterpretation... weproposcthatthehoteoronalgasindiscandellipl ical acerebionontoashock healedaquasi hydrostalicholcorona enderploretheconsequencesofthisproposa," In \\ref{sec:interpretation}, we propose that the hot coronal gas in disc and elliptical galaxies has a common origin -- accretion onto a shock-heated quasi-hydrostatic hot corona – and explore the consequences of this proposal."
+t Meine faands NeLlarcuntikelylobethesourceestheholcoronalgasingalaricso fan ," We include a short appendix, in which we show that ejecta from SNe-Ia and SNe-II are unlikely to be the source of the hot coronal gas in galaxies of any morphological type."
+We first introduce the sample of galaxies that we use in this study. which are taken from previously published X-ray studies.," We first introduce the sample of galaxies that we use in this study, which are taken from previously published X-ray studies."
+ As the X-ray flux from point sources can be comparable to (or even dominate) that from the hot gas of normal galaxies. high quality X- observations are required to determine the emission that truly originates from the coronal gas.," As the X-ray flux from point sources can be comparable to (or even dominate) that from the hot gas of normal galaxies, high quality X-ray observations are required to determine the emission that truly originates from the coronal gas."
+ For this reason. we have limited our compilation of galaxies to those that have aand/or oobservations reported in the literature.," For this reason, we have limited our compilation of galaxies to those that have and/or observations reported in the literature."
+ We describe the sample below. with a brief discussion of how the X-ray luminosity of the hot gas was measured in each case.," We describe the sample below, with a brief discussion of how the X-ray luminosity of the hot gas was measured in each case."
+ Our sample of elliptical galaxies with X-ray luminosities and temperatures is taken from two studies: andMJIO., Our sample of elliptical galaxies with X-ray luminosities and temperatures is taken from two studies: and.
+". The D06 sample is comprised of 18 low optical luminosity (Li;<3.101""Ly. ) field elliptical galaxies observed withChandra. drawn from a larger sample of early type galaxies with aarchival data (C. Jones. in prep.)."," The D06 sample is comprised of 18 low optical luminosity $L_{\rm B} \la 3\times10^{10} L_{\rm B,\odot}$ ) field elliptical galaxies observed with, drawn from a larger sample of early type galaxies with archival data (C. Jones, in prep.)."
+ The non-thermal X-ray emission Tom low-mass X-ray binaries (LMXBs) was accounted for in DOG by spatial excision of bright sources and spectral modelling (using a power-law component) of unresolved sources., The non-thermal X-ray emission from low-mass X-ray binaries (LMXBs) was accounted for in D06 by spatial excision of bright sources and spectral modelling (using a power-law component) of unresolved sources.
+ MJIO selected heir sample of 23 nearby field early type galaxies from previously oublished. X-ray catalogues2)., MJ10 selected their sample of 23 nearby field early type galaxies from previously published X-ray catalogues.
+ Approximately half of their sample was observed with aand the other half withXAZAM-Newton.. with several galaxies wving data from both satellites.," Approximately half of their sample was observed with and the other half with, with several galaxies having data from both satellites."
+ The MJIO sample nicely complements that of DO6. as it is comprised of relatively bright systems CL=10Lg.," The MJ10 sample nicely complements that of D06, as it is comprised of relatively bright systems $L_{\rm K} \ga 10^{11} L_{\rm
+K,\odot}$ )."
+ In similar fashion to DOG. MILO accounted for non-thermal emission from LMXBs through the inclusion of a power-law component in their spectral modelling.," In similar fashion to D06, MJ10 accounted for non-thermal emission from LMXBs through the inclusion of a power-law component in their spectral modelling."
+ While it is possible to remove the contribution from unresolved non-thermal sources on the basis of their spectra (which can be can be differentiated from thermal spectra if there are sufficient photons). removal of X-ray emission from a potential point source population is more difficult.," While it is possible to remove the contribution from unresolved non-thermal sources on the basis of their spectra (which can be can be differentiated from thermal spectra if there are sufficient photons), removal of X-ray emission from a potential point source population is more difficult."
+ An good example is the emission originating from the so-called Galactic Ridge., An good example is the emission originating from the so-called Galactic Ridge.
+ This emission had previously been believed to come from hot gas. since its spectrum is consistent with an optically-thin plasma (with 7 Ky and even displays a prominent iron K line.," This emission had previously been believed to come from hot gas, since its spectrum is consistent with an optically-thin plasma (with $T\sim10^{7-8}$ K) and even displays a prominent iron K line."
+ However. pointed out that the X-ray surface brightness traces almost perfectly," However, pointed out that the X-ray surface brightness traces almost perfectly"
+this method.,this method.
+ The inclination (first guess following Verheijen Saucisi 1993). centre. aud the systemic velocity were obtained by miuimnizing the residual velocity field. which was obtained by subtracting the model (rom the observed velocity field.," The inclination (first guess following Verheijen Sancisi 1993), centre, and the systemic velocity were obtained by minimizing the residual velocity field, which was obtained by subtracting the model from the observed velocity field."
+ We can observe that several peculiarities of the velocity field are well reproduced by the model. showing the robustuess of the method.," We can observe that several peculiarities of the velocity field are well reproduced by the model, showing the robustness of the method."
+ Tn this re;»ort we have explained why tlie effect of warps cannot be neglectος iu the procedure to derive rotation curVern. alid have stressed why this might affect auy further kineuatical or dvnamical analysis.," In this report we have explained why the effect of warps cannot be neglected in the procedure to derive rotation curves, and have stressed why this might affect any further kinematical or dynamical analysis."
+ We have oroposed a «lifferent approach ιο extract the rotation curve ln wa‘pec. highly inclined galaxies. which can now be determinecl accurately.," We have proposed a different approach to extract the rotation curve in warped, highly inclined galaxies, which can now be determined accurately."
+ Based on this 1ew approach. we are abe lo accurately trace the frequency o. peculiarities iu our salnp eo. TBB ealaxies.," Based on this new approach, we are able to accurately trace the frequency of peculiarities in our sample of TBB galaxies."
+ We report au iuc‘easing treud of kinematical opsidediess. [rom spleroical bulge galixles 1ΟΝΕs TBB galaxies.," We report an increasing trend of kinematical lopsidedness, from spheroidal bulge galaxies towards TBB galaxies."
+" Concerning the qtestiou whether in""actious coutribute siguilicauty to bar formation aud subsequen b/p structu'e. basec on our sample we cii1u the theoretical preclictiou that TBBs lave Inore preonninen peculiaritjes. aud ineractions seem to ve tl1e likely formation mechanism of the TBB galaxies."," Concerning the question whether interactions contribute significantly to bar formation and subsequent b/p structure, based on our sample we confirm the theoretical prediction that TBBs have more prominent peculiarities, and interactions seem to be the likely formation mechanism of the TBB galaxies."
+ Uποηlately. wih our limited nuber of objecs we are not able to extend tils conclusion to the Class of b/p spi‘al galaxies. and this res| mieht be valil ouly for the ocal univese.," Unfortunately, with our limited number of objects we are not able to extend this conclusion to the class of b/p spiral galaxies, and this result might be valid only for the local universe."
+ However. this tren exists and fwtler analysis might be colucted ou a arger sample of «)bjects.," However, this trend exists and further analysis might be conducted on a larger sample of objects."
+ However. fo: models of galaxy evoltLOL and to iivestigate the role assumed ον interaction events il reshaping galaxy components (bars. bulges aud disks) it is essential that the assyinajetries: uic lopsideduess be well-deteriiied over as large a ναιjca range aie| recishifts as possible.," However, for models of galaxy evolution and to investigate the role assumed by interaction events in reshaping galaxy components (bars, bulges and disks) it is essential that the asymmetries and lopsidedness be well-determined over as large a dynamical range and redshifts as possible."
+ Institments like the Atacama Large Millimeter Array (ALMA) are needed for good salistics of warping phenomena at cliffereut redshifts 1o investigate 1ie. still uuknowu origin of warped galaxies. whether they are sha»ed under the influence oL a decoupled dark aud luminous matter potential or whether they are dominated by their euviroument.," Instruments like the Atacama Large Millimeter Array (ALMA) are needed for good statistics of warping phenomena at different redshifts to investigate the still unknown origin of warped galaxies, whether they are shaped under the influence of a decoupled dark and luminous matter potential or whether they are dominated by their environment."
+ O.Lem We wish to acknowledge Tom Oosterloo and Baffaella Morgauti for helpful discussious., 0.1cm We wish to acknowledge Tom Oosterloo and Raffaella Morganti for helpful discussions.
+ DV is grateful for financial support of the uuder uunber GRIN 115 “The Magellanic System. Galaxy Interaction and the Evolution of Dwarf Galaxies.," DV is grateful for financial support of the under number GRK 118 'The Magellanic System, Galaxy Interaction and the Evolution of Dwarf Galaxies'."
+ We thauk the Observatoire de Bordeaux for the warm hospitality during the workshop., We thank the Observatoire de Bordeaux for the warm hospitality during the workshop.
+ Q.5cm, 0.5cm
+in eveles 20993. 29904 ancl 29995.,"in cycles 29993, 29994 and 29995."
+ his is 1 sume as that found by Larrop-Allin (1999b) of AO=0.07790.0002. and Schwope ((1997) of No=0.0779.," This is the same as that found by Harrop-Allin (1999b) of $\Delta\phi=0.0779\pm 0.0002$, and Schwope (1997) of $\Delta\phi=0.0779$."
+ The lack of variation in the curation of the eclipse is remarkable anc implies that the accretion spot is at the same Latitude for these observations., The lack of variation in the duration of the eclipse is remarkable and implies that the accretion spot is at the same latitude for these observations.
+ On the other hand. this is shorter by 6 s than the 0.0782.0.0792 modelled by Schwope ((2001) from the soft X-ray eclipses. where the ingress was uncertain due to the low countrate and absorption.," On the other hand, this is shorter by 6 s than the $0.0782-0.0792$ modelled by Schwope (2001) from the soft X-ray eclipses, where the ingress was uncertain due to the low countrate and absorption."
+ For a particular mass ratio. the width of the eclipse determines the inclination of the svstem.," For a particular mass ratio, the width of the eclipse determines the inclination of the system."
+" Using qo=AlfAL,=0.25 (Schwope 22001). the inclination is 85.0"" (assuming M,=O.9AL. and p=40°. €=5 where AJ, is the mass of the white chvarl ancl 2.6 are the magnetic longitude and colatitucle of the accretion region."," Using $q = M_2 / M_1 = 0.25$ (Schwope 2001), the inclination is $85.0^\circ$ (assuming $M_1=0.9 M_{\odot}$ and $\beta=40^\circ$, $\zeta=50^\circ$, where $M_1$ is the mass of the white dwarf and $\beta, \zeta$ are the magnetic longitude and colatitude of the accretion region."
+ These correct for the location of the aeeretion spot on the surface of the primary: see below)., These correct for the location of the accretion spot on the surface of the primary; see below).
+ Complete ingress of the stream occurs ato ὁ=0.982.0.070.0.977 and 0.982 for eclipses 20982. 20003. 20904 and 29995) respectively.," Complete ingress of the stream occurs at $\phi=0.982, 0.979, 0.977$ and $0.982$ for eclipses 29982, 29993, 29994 and 29995 respectively."
+ During the period of complete eclipse the secondary is the only contributor to the light curve and provides a constant contribution which is greatest in the red. band. as expected. with no contribution in the blue band.," During the period of complete eclipse the secondary is the only contributor to the light curve and provides a constant contribution which is greatest in the red band, as expected, with no contribution in the blue band."
+ Figure 5. shows more clearly that the observations reach total stream. eclipse at dilferent. phases even though the spot. ingresses/egresses both occur at the same time., Figure \ref{fig:inoutsmooth} shows more clearly that the observations reach total stream eclipse at different phases even though the spot ingresses/egresses both occur at the same time.
+ This rapid change in the accretion magnetic Ποιά interaction between successive eclipses was first observed by Glenn (1994) who found a cdillerence in the time it took for the gaream to be completely. eclipsed of more than a minute »etween their two successive eclipses., This rapid change in the accretion $\--$ magnetic field interaction between successive eclipses was first observed by Glenn (1994) who found a difference in the time it took for the stream to be completely eclipsed of more than a minute between their two successive eclipses.
+ The shape of the aceretion stream eclipse profile in the observed light curves gives a qualitative prediction of where xiehtness enhancements can be expected in the accretion stream., The shape of the accretion stream eclipse profile in the observed light curves gives a qualitative prediction of where brightness enhancements can be expected in the accretion stream.
+ When the stream is brighter at later phases in the stream eclipse. as in eveles 90053 and 29995in (Figure in 5)). iere must be more emission in the threading region. as js is the only part of the stream still visible.," When the stream is brighter at later phases in the stream eclipse, as in cycles 29982 and 29995 (Figure \ref{fig:inoutsmooth}) ), there must be more emission in the threading region, as this is the only part of the stream still visible."
+ For the same reason. i£ the final stage of the stream eclipse takes ace at later phases. the threacing region must have moved: jus in ονο]ος 29982 and. 29995.r either the accretion stream is brought further out [rom the line of centres during the rallistic part of the trajectory. or the magnetically channelec stream rises further out of the orbital plane (sve address this in Section 5. below).," For the same reason, if the final stage of the stream eclipse takes place at later phases, the threading region must have moved: thus in cycles 29982 and 29995, either the accretion stream is brought further out from the line of centres during the ballistic part of the trajectory, or the magnetically channeled stream rises further out of the orbital plane (we address this in Section \ref{sec:mapoverview} below)."
+ This could be caused either by a small change in the location of the accretion spot on the primary surface. moving the accretion spot further from the line of centres. or perhaps a change in the velocity of the stream as it leaves the Ly point.," This could be caused either by a small change in the location of the accretion spot on the primary surface, moving the accretion spot further from the line of centres, or perhaps a change in the velocity of the stream as it leaves the $_{1}$ point."
+ Likewise. if the accretion stream. is fainter near the white dwarl. the svstem brightness will be fainter immediately after acerction region ingress.," Likewise, if the accretion stream is fainter near the white dwarf, the system brightness will be fainter immediately after accretion region ingress."
+ This may be the case for evcle 29995 (Figure 5))., This may be the case for cycle 29995 (Figure \ref{fig:inoutsmooth}) ).
+ For evele 29995 there is a rise in the vellowfred ratio during the stream ingress (Figure 2)). so that successively more red light is blocked from the stream as the eclipse progresses.," For cycle 29995 there is a rise in the yellow/red ratio during the stream ingress (Figure \ref{fig:huaqr04colours}) ), so that successively more red light is blocked from the stream as the eclipse progresses."
+ This implies that the last part of the stream to be eclipsed. near the threading region. emits relatively less at longer wavelengths than the rest of the stream.," This implies that the last part of the stream to be eclipsed, near the threading region, emits relatively less at longer wavelengths than the rest of the stream."
+ This rise is not evident in evcle 29993. but maybe. present at a reduced level in evele 29994 (Figure 3)).," This rise is not evident in cycle 29993, but maybe present at a reduced level in cycle 29994 (Figure \ref{fig:huaqr03colours}) )."
+ This trend. is consistent with a threading region which is inereasinely hotter in the three successive. eclipses., This trend is consistent with a threading region which is increasingly hotter in the three successive eclipses.
+ At the same time as the large vellowτοῦ. ratio change in evele 29995 there is no change in the blue/vellow ratio., At the same time as the large yellow/red ratio change in cycle 29995 there is no change in the blue/yellow ratio.
+ After eclipse when the stream close to the primary is uncovered. first. the cdilferent vellow/red ratios in eveles 29993 and 29995 are consistent with a hotter stream near the white dwarf in eclipse 29993 than in 29995.," After eclipse when the stream close to the primary is uncovered first, the different yellow/red ratios in cycles 29993 and 29995 are consistent with a hotter stream near the white dwarf in eclipse 29993 than in 29995."
+ Xn increase in the blue/zvellow: ratio after eclipse is seen in all three eveles (29993. 29994 and 29995) and also implies that the maenctically confined stream material is becoming hotter towards the threacling region.," An increase in the blue/yellow ratio after eclipse is seen in all three cycles (29993, 29994 and 29995) and also implies that the magnetically confined stream material is becoming hotter towards the threading region."
+ The pre-eclipse dip around ὁzO.ST in evele 29994 (ligure 3)) is also sliehthy bluer than at earlier or later phases., The pre-eclipse dip around $\phi\approx 0.87$ in cycle 29994 (Figure \ref{fig:huaqr03colours}) ) is also slightly bluer than at earlier or later phases.
+ This was also the case in the light curves of Larrop- ((1999b)., This was also the case in the light curves of Harrop-Allin (1999b).
+ We have inferred from a qualitative analysis of the light curves that the trajectory and/or intensity of the stream varies over successive orbital periods. ancl we can identify brightness enhancements along the stream length.," We have inferred from a qualitative analysis of the light curves that the trajectory and/or intensity of the stream varies over successive orbital periods, and we can identify brightness enhancements along the stream length."
+ However. it is more cdillicult to infer the location of the brightness," However, it is more difficult to infer the location of the brightness"
+llish-mass stars plav an important role in dvuaniucally agitating and ionizing the interstellar medium in galaxies and strongly influence galactic evolution.,High-mass stars play an important role in dynamically agitating and ionizing the interstellar medium in galaxies and strongly influence galactic evolution.
+ It is therefore crucial to understaud. (heir formation mechanisms., It is therefore crucial to understand their formation mechanisms.
+ Despite numerous (theoretical and observational studies attempting to address this issue in the last few decades. the detailed," Despite numerous theoretical and observational studies attempting to address this issue in the last few decades, the detailed"
+The processes menutioued above. except for Coulomb scattering. are tuplemented through the Peuclope physics package (Salvat.Feruaudez-Varea&Seipan2006).. which is valid above ~ 250 eV. Broensstrahluug iu Penelope includes clectron-clectron breimsstraliluung. which can be siguificant above several hundred keV. in flares 2007).. usine cross sections calculated after Seltzer&Berecr(1985).,"The processes mentioned above, except for Coulomb scattering, are implemented through the Penelope physics package \citep{salvat06}, which is valid above $\sim$ 250 eV. Bremsstrahlung in Penelope includes electron-electron bremsstrahlung, which can be significant above several hundred keV in flares \citep{kontar07}, using cross sections calculated after \citet{seltzer85}."
+". Coulomb scattering needs special consideration because its 10621 free path is much smaller than that of anv other process and a full treatment would take too mach time,", Coulomb scattering needs special consideration because its mean free path is much smaller than that of any other process and a full treatment would take too much time.
+ GEANT simulates the effect of iiultiple Coulomb interactions after a οἴνοι step as a statistical expectation. instead of treatingE them oue by one.," GEANT4 simulates the effect of multiple Coulomb interactions after a given step as a statistical expectation, instead of treating them one by one."
+ Chin&Spyroui(2009) sugeested that because the different GIEANT limodels for proton inelastic collisions are not consistent with cach other iu the euergy auge of tens of MeV. these models ave not ready for problem solving: our initial stations using these packages coufixii this conclusion for our application.," \citet{chin09} suggested that because the different GEANT4 models for proton inelastic collisions are not consistent with each other in the energy range of tens of MeV, these models are not ready for problem solving; our initial simulations using these packages confirm this conclusion for our application."
+ We found that the proton inclastic process as simulated iu the IEulrouPlhivsicsQCGSP.BBERTITIP plivsies 11:odule produces a coutimmun-like photon spectrin extending to  20 MeV. This couflicts with most observatious of Gamma-ray flares (Share&Murphy1995.¢.¢.).. iu which this compoucut falls off dramatically at 8 MeV. above the complex of lines from the de-excitation of nitrogen. carbon and oxveen.," We found that the proton inelastic process as simulated in the HP physics module produces a continuum-like photon spectrum extending to $\sim$ 20 MeV. This conflicts with most observations of gamma-ray flares \citep[e.g.]{share95}, in which this component falls off dramatically at 8 MeV, above the complex of lines from the de-excitation of nitrogen, carbon and oxygen."
+ The observed spectrin is also more dominated by individual lines aud looks less like a σολπι., The observed spectrum is also more dominated by individual lines and looks less like a continuum.
+ Another package. IEulrouPhiysicsQCGHIIP. shows a iore realistic cutoff! around 8 MeV. but still produces oulv a continuume-like shape aud not the observed couples. of de-excitation lines.," Another package, HP, shows a more realistic cutoff around 8 MeV but still produces only a continuum-like shape and not the observed complex of de-excitation lines."
+ In Figure 1 we show the differeuce between these two plysics modules for the proton melastie scattering component aud the eutire proton-derived flux. including the higlh-euergey compoucut frou piou production.," In Figure \ref{fig:modcompare}
+ we show the difference between these two physics modules for the proton inelastic scattering component and the entire proton-derived flux, including the high-energy component from pion production."
+ For this comparison we injected mono-energetie protons of 1 GeV downward iuto the mocel Sun described iun section 2.2 and recorded all the photons comuneout at an anele ) from the solar normal such that cos(.7)20.8., For this comparison we injected mono-energetic protons of 1 GeV downward into the model Sun described in section \ref{sec:solarmod} and recorded all the photons comingout at an angle $\beta$ from the solar normal such that $\cos(\beta)>0.8$.
+ Pion production and decay. ou the other hand. are sinipler processes aud there is agreement among nultiple imodols for their simulation.," Pion production and decay, on the other hand, are simpler processes and there is agreement among multiple models for their simulation."
+ As can be seeu iu Figure L.. the two CEANT plivsies modules agree well as to this photou colmponent (dominant above LO MeV).," As can be seen in Figure \ref{fig:modcompare}, the two GEANT4 physics modules agree well as to this photon component (dominant above 10 MeV)."
+ We also compared these resultsIL from GEANTI with the output of the code developed by Reuven Ramaty. Ronald Murpliy. aud others. which has been successfully applied to large flares 1987).," We also compared these results from GEANT4 with the output of the code developed by Reuven Ramaty, Ronald Murphy, and others, which has been successfully applied to large flares \citep{murphy87}."
+.. For this comparison. we put mono-cherectic protons of 1 GeV into a atimosphere of hydroseu and Lelimm with ΠοΤΠ = 0.1. and recorded all the photons produced.," For this comparison, we put mono-energetic protons of 1 GeV into a atmosphere of hydrogen and helium with He/H = 0.1, and recorded all the photons produced."
+ As shown in Fieure 2.. the results are cousisteut with GEANTL: in the enerev range δ15 MeV. the difference is less than," As shown in Figure \ref{fig:pioncompare}, the results are consistent with GEANT4; in the energy range 8–15 MeV, the difference is less than."
+ Since we believe we can therefore trust GEANT for piou processes as well as clectromaguctic cascades. we will ignore the inelastic process for this work and concentrate| on the δ15 MeV range. which should be dominated by bremesstrahhme both from primary electrons aud from secondary electrous ancl positrous from pion decay.," Since we believe we can therefore trust GEANT4 for pion processes as well as electromagnetic cascades, we will ignore the inelastic process for this work and concentrate on the 8–15 MeV range, which should be dominated by bremsstrahlung both from primary electrons and from secondary electrons and positrons from pion decay."
+ Tu the present simulation. we treat the solar atmosphere as parallel lavers.," In the present simulation, we treat the solar atmosphere as parallel layers."
+ This is a reasonable approximation because the size of the cussion region is always uch smaller than the solar radius., This is a reasonable approximation because the size of the emission region is always much smaller than the solar radius.
+ We use the analytical approximation of Ἱνοίοκαotal.(2007) to the Tarvard-Sinithsouian refereuce atmosphere (Gingerichetal.1971). as the solar mass-censity profile. which is Tere ids the height measured from plhotosphere and  is the scale height. which is ~LOO kan for 2 <0 aud ~110 an Or 2 20.," We use the analytical approximation of \citet{kotoku07} to the Harvard-Smithsonian reference atmosphere \citep{gingerich71} as the solar mass-density profile, which is Here $z$ is the height measured from photosphere and $h$ is the scale height, which is $\sim 400$ km for $z<$ 0 and $\sim 110$ km for $z>$ 0."
+ The specific vertical structure 110del will not affect the simulation result. however.," The specific vertical structure model will not affect the simulation result, however."
+ All the processes we care about except decays depend ouly ou the column deusity along the path., All the processes we care about except decays depend only on the column density along the path.
+ Because our moclel solar atmosphere is 11ade up of xwallel slabs. we cau always transform the paraimecter + into cohuun depth without changing our results.," Because our model solar atmosphere is made up of parallel slabs, we can always transform the parameter $z$ into column depth without changing our results."
+" The oulv exception is decay processes, which depend on time."," The only exception is decay processes, which depend on time."
+" However. the longest lifetime we need to consider is that of z*. which is 2.6\10 ὃν, Diving this short time. the pion travels less than 1 kan even if it has an energy of up to «1041 MeV, which is much shorter than the leneth scale of the svsteii."," However, the longest lifetime we need to consider is that of $\pi^{\pm}$, which is $\times 10^{-8}$ s. During this short time, the pion travels less than 1 km even if it has an energy of up to $\times10^{4}$ MeV, which is much shorter than the length scale of the system."
+ So the decay always happens approximately where the short ived particles are produced aud the result will not change with different vertical structures., So the decay always happens approximately where the short lived particles are produced and the result will not change with different vertical structures.
+ Another consideration is hat when the chromosphere is bombarded by the flare accelerated particles. it will evaporate aud fill up the magnetic oop.," Another consideration is that when the chromosphere is bombarded by the flare accelerated particles, it will evaporate and fill up the magnetic loop."
+ This will change the vertical structure., This will change the vertical structure.
+ If the particles were to interact lieh in a narrow coluun. the pattern of photon escape as a function of solar normal augle would be very differcut. with taugeutial escape much easier.," If the particles were to interact high in a narrow column, the pattern of photon escape as a function of solar normal angle would be very different, with tangential escape much easier."
+ However. Asclavanden&Beuz(1997) find that the deusity of evaporation upflow is around. 1019 cur7.," However, \citet{aschwanden97} find that the density of evaporation upflow is around $^{10}$ $^{-3}$."
+ Take the longest flare oop with a leneth of around 104+ci. the column deusitv change will be less than 0.01 οσα7.," Take the longest flare loop with a length of around $^{11}$cm, the column density change will be less than 0.01 $^{-3}$."
+" As we will show in §3.1.. all the processes we are interested iu happen at a colui density ercater that 1 οὐο,"," As we will show in \ref{sec:simu}, all the processes we are interested in happen at a column density greater that 1 $^{-3}$."
+ Therefore. evaporation VA1ould not stronely iufiueuce our results.," Therefore, evaporation should not strongly influence our results."
+ The elemental abundance of our model Sum is taken frou Crevesse.Asplund&Sauval(2007)... and we assume +iat the abundances of the photosplere and corona are the same.," The elemental abundance of our model Sun is taken from \citet{grevesse07}, and we assume that the abundances of the photosphere and corona are the same."
+ In future work we will iuplemenut more realistic shotospheric and coronal abundances. includiug the cuhancement of low first jonization poteutial (FIP) elements in the corona.," In future work we will implement more realistic photospheric and coronal abundances, including the enhancement of low first ionization potential (FIP) elements in the corona."
+ Ixotokuetal.(2007) used a pure hydrogen atinosphliere. which uuderestimates brenisstrahluug cficiency. since soblat rises as approximately the square of atomic umber Z. aud also unclerestimates pair production. the cross-section for which iucreases even dramatically with Z.," \citet{kotoku07} used a pure hydrogen atmosphere, which underestimates bremsstrahlung efficiency, since that rises as approximately the square of atomic number $Z$ , and also underestimates pair production, the cross-section for which increases even dramatically with $Z$ ."
+ We fiud that 10 MeV photons. for example. produce more positrons oe 1a realistic atmosphere than in a lbydrogcu atinospliere.," We find that 10 MeV photons, for example, produce more positrons in a realistic atmosphere than in a hydrogen atmosphere."
+2007).,.
+.. The problemi arose because all NMAM-rewton lighteurves start at the arrival time of first photon by default: heuce. source aud vackeround liehteurves. and lieghtcwrves from three EPIC’ detectors ALOSL. MOS2. and nn are asvuchronous bv default.," The problem arose because all XMM-newton lightcurves start at the arrival time of the first photon by default; hence, source and background lightcurves, and lightcurves from the three EPIC detectors— MOS1, MOS2, and pn— are asynchronous by default."
+ Coubinine τοσο lighteurves (6.8. combining iustrünients. or vackeround subtraction) often resulted in PDS with artificial broken power law shapes.," Combining these lightcurves (e.g. combining instruments, or background subtraction) often resulted in PDS with artificial broken power law shapes."
+ Therefore we make no use of the PDS and instead rely on our well established method of using low state oenüssion spectra (powerlawlard2006)/ at conspicuously high buuinosities to ideutifv black hole caudidates (Barnardetal.2008:Barnard&I&olb2009:etal. 2011)..," Therefore we make no use of the PDS and instead rely on our well established method of using low state emission spectra \citep[power law emission with photon index 1.4--1.7 and little to no thermal emission, ][]{mr06} at conspicuously high luminosities to identify black hole candidates \citep{barnard08, barnard09,barnard11b}."
+ We preseut the most cetailed justification of our selection criteria in Barnardetal.C2011).., We present the most detailed justification of our selection criteria in \citet{barnard11b}.
+ Iu this paper we present our analysis of δ Chandra ACTS observations of J0012.3|1115 over ~12 vears and our sereudipitous OST observation. along with our re-analysis of the 60 ks 2002 NATNewton observation.," In this paper we present our analysis of 84 Chandra ACIS observations of J0042.3+4115 over $\sim$ 12 years, and our serendipitous HST observation, along with our re-analysis of the 60 ks 2002 XMM-Newton observation."
+ We use the IST data to place RNJ0012.31£115. in ALI. and argue for a low mass donor.," We use the HST data to place J0042.3+4115 in M31, and argue for a low mass donor."
+ We use loue-ter aud short-term variability. aud also emission spectra. to reiustate J0012.3|1115 as a black hole caudidate.," We use long-term and short-term variability, and also emission spectra, to reinstate J0042.3+4115 as a black hole candidate."
+ We discuss the observations and data analysis in the next section. followed by our results in Section 3. aud a discussion iu Section 1.," We discuss the observations and data analysis in the next section, followed by our results in Section 3, and a discussion in Section 4."
+ We analyzed 81 Claudra ACIS observations of the central region of M31. spaced over ~12 years. using CIAO version L3.," We analyzed 84 Chandra ACIS observations of the central region of M31, spaced over $\sim$ 12 years, using CIAO version 4.3."
+" For cach observation we extracted 0.37.0. keV. source and backerouud spectra| from circular regious with 10"" radius he background region was close to the source region. aud source free,"," For each observation we extracted 0.3–7.0 keV source and background spectra from circular regions with $''$ radius; the background region was close to the source region, and source free."
+ Corresponding response natrices and ancillary respouse files were also uade., Corresponding response matrices and ancillary response files were also made.
+ We obtained 0.510 keV. laminosities frou cach observation using° NSPEC version 12.6.0., We obtained 0.3–10 keV luminosities from each observation using XSPEC version 12.6.0.
+ Observations with 2200 net source counts were recle fitted with absorbed power law models: were erouped to eive. at least .20 counts ov eebiu., Observations with $>$ 200 net source counts were freely fitted with absorbed power law models; spectra were grouped to give at least 20 counts per bin.
+ For observatious with «200 net source counts we assumed an absorbed power law model with Ny —104107 |atom ? |and D = 1.5. and found the 0.310 keV huninosity equivaleut o 1 count |. then nultiplied this conversion Actor bv the iufensitv: we chose this meocel )ecause it approximates the best fit to our deepest Chandra observation of J0012.3|L115 in its ow state.," For observations with $<$ 200 net source counts we assumed an absorbed power law model with $N_{\rm H}$ $\times$ $^{21}$ atom $^{-2}$ and $\Gamma$ = 1.5, and found the 0.3–10 keV luminosity equivalent to 1 count $^{-1}$, then multiplied this conversion factor by the intensity; we chose this model because it approximates the best fit to our deepest Chandra observation of J0042.3+4115 in its low state."
+ Luminositv uucertaimties for freely fitted spectra are estimated by NSPEC by calculating a range of fluxes obtained by varvine the ciission xuwanmeters: the uncertainties for the faint spectra are derived directly from iutensity uncertainties., Luminosity uncertainties for freely fitted spectra are estimated by XSPEC by calculating a range of fluxes obtained by varying the emission parameters; the uncertainties for the faint spectra are derived directly from intensity uncertainties.
+ Additionally. we analyzed the 60 ks 2002 January NMMN-Neswtou observation of M31. (Rev 281) with SAS version 10.0.0.," Additionally, we analyzed the 60 ks 2002 January XMM-Newton observation of M31 (Rev 381) with SAS version 10.0.0."
+" We extracted 0.310 keV EPIC-pu liehteurves aud spectra from circular source and backerouud reeious with 15"" radius.", We extracted 0.3--10 keV EPIC-pn lightcurves and spectra from circular source and background regions with $''$ radius.
+ The backeround region was on the same CCD as the source region. and at a simular offaxis anele.," The background region was on the same CCD as the source region, and at a similar off-axis angle."
+ The spectra were grouped to eusure a minimi of 50 counts per bin., The spectra were grouped to ensure a minimum of 50 counts per bin.
+ JOO12.3UST|L115 was screudipitously observed in one of our ACS/WTC observations of M31 transents, J0042.3+4115 was serendipitously observed in one of our HST ACS/WFC observations of M31 transients.
+ Observation j9udlT0l10 was made ou 2009 August 25. with the E135W filter for 1360 s. Wo registered a combined Chandra 0.3.7.0 keV nuage and the IIST image to the same D band inage provided by the Local Group Galaxy Survey (Masseyetal.2006) using the IRAF taskΜΑΗ: we used X-raybright elobular clusters to register the Chandra image. and bright. unsaturated stars to register the IST inage.," Observation j9ud17010 was made on 2009 August 25, with the F435W filter for 4360 s. We registered a combined Chandra 0.3–7.0 keV image and the HST image to the same B band image provided by the Local Group Galaxy Survey \citep{massey06} using the IRAF task; we used X-ray bright globular clusters to register the Chandra image, and bright, unsaturated stars to register the HST image."
+ We determined the best N-ray positionof J0012.3|1115 with the IRAF taskIMCENTROID., We determined the best X-ray position of J0042.3+4115 with the IRAF task.
+ We used the IRAF package DAOPIIOTII to perform photometry on the stars within 3e of the position of J0012.5|1115., We used the IRAF package DAOPHOTII to perform photometry on the stars within $\sigma$ of the position of J0042.3+4115.
+" The centroid. of. the ταν- cussion⋅⋅ frou. JOO12.3|1115 was located at RA = 00;[2:22.95|. Dec= 11:15:35.23. with lo uncertainties of 0.009” oe LRA and 0.007"" in Dec. Combining this with the Pus."," The centroid of the X-ray emission from J0042.3+4115 was located at RA = 00:42:22.954, Dec = 41:15:35.23, with $\sigma$ uncertainties of $''$ in RA and $''$ in Dec. Combining this with the r.m.s."
+" uncertainties iu registration vields N-rav positional uncertainties of 0.09"" in. RA and 0.19” in Dec. Fiewe. .1 shows a detail. our TST PMimage. y.uperposed with au ellipse represeutius the 26 ucertaiuties iu the position of J0012.3|L115."," uncertainties in registration yields X-ray positional uncertainties of $''$ in RA and $''$ in Dec. Figure \ref{opt} shows a detail our HST image, superposed with an ellipse representing the $\sigma$ uncertainties in the position of J0042.3+4115."
+ There are several stars within the ellipse. the brightest⋅ of. which. has a οσα- D inaeuitude. of," There are several stars within the ellipse, the brightest of which has a Vega B magnitude of"
+ There are several stars within the ellipse. the brightest⋅ of. which. has a οσα- D inaeuitude. of.," There are several stars within the ellipse, the brightest of which has a Vega B magnitude of"
+ (Cliaueetal.2001).. (Watsonetal.2007).., \citep{chiang2001}. \citep{2007prpl.conf..523W}.
+ Tauris while hotter ones are called Herbie Ac/Be stars., Tauris while hotter ones are called Herbig Ae/Be stars.
+ Typically. T Tauri stars are interpreted to be PAIS stars of roughly solar mass or less. while Herbig Ac/Be stars are voung stars of intermediate mass. DAL.<MSM," Typically, T Tauri stars are interpreted to be PMS stars of roughly solar mass or less, while Herbig Ae/Be stars are young stars of intermediate mass, $2 M_\odot < M < 8
+M_\odot$."
+ But such a one-dimensional classification cannot fully represent the range of PAIS: stellar properties. which depend on both mass and age.," But such a one-dimensional classification cannot fully represent the range of PMS stellar properties, which depend on both mass and age."
+ This is particularly true or dntermeciate-mass stars. which are stable agaist convection and consequcutly evolve horizontally across he [ertzsprune-Russell diagram. towards the main sequence. rsme m surface temperature aud chaugius in spectral type as they eo.," This is particularly true for intermediate-mass stars, which are stable against convection and consequently evolve horizontally across the Hertzsprung-Russell diagram towards the main sequence, rising in surface temperature and changing in spectral type as they go."
+ As very voune intermediatc-nass stars contract toward becoming Herbie Ac/Be stars. thev first pass through a stage in which they are classified as Τ Tauri stars. but ones with vastly Heher bhuuinositv than typical.," As very young intermediate-mass stars contract toward becoming Herbig Ae/Be stars, they first pass through a stage in which they are classified as T Tauri stars, but ones with vastly higher luminosity than typical."
+ These “imtermieciate-nass TO Tawi stars” (IMTTS: Calvet et al., These ``intermediate-mass T Tauri stars” (IMTTS; Calvet et al.
+ 2001) are the vounecst imtermediate-niass stars available for observation at visible or near-IR waveleueths., 2004) are the youngest intermediate-mass stars available for observation at visible or near-IR wavelengths.
+ Due to the slope of the stellar initial mass function aud the rapid speed at which stars evolve through this carly stage. IMITTSs are relatively rare aud lave been the subject of fow detailed studies.," Due to the slope of the stellar initial mass function and the rapid speed at which stars evolve through this early stage, IMTTSs are relatively rare and have been the subject of few detailed studies."
+ Yetthese objects potentially offer, Yetthese objects potentially offer
+presents the rrelation for nearby with reverberation-based BH mass measurements (Figure 1;; a).,presents the relation for nearby with reverberation-based BH mass measurements (Figure \ref{MBH}; ).
+ The bulge luminosities are based on GALFIT modeling of or F547M images., The bulge luminosities are based on GALFIT modeling of /F550M or F547M images.
+" Since our virial BH masses are HST//F550Mbased on a radius-luminosity relation derived from these same objects, the BH masses are directly comparable, hopefully minimizing systematic errors."," Since our virial BH masses are based on a radius-luminosity relation derived from these same objects, the BH masses are directly comparable, hopefully minimizing systematic errors."
+" However, we note that their BH masses range from —10?Mo,10"" larger than the BH masses in our sample by about 2 orders of magnitude."," However, we note that their BH masses range from $10^7-10^9\msun$, larger than the BH masses in our sample by about 2 orders of magnitude."
+ This comparison can tell us whether the rrelationMi Lpuigethat is observed in massive AGNs still exists at the low masses probed here., This comparison can tell us whether the relation that is observed in massive AGNs still exists at the low masses probed here.
+" In order to make this comparison, we adopt the f factor fromdex)."," In order to make this comparison, we adopt the $f$ factor from."
+". We (2004,fit the sample with the following log-linear relation where Mo is the median magnitude measured from the fitted sample, which is held constant during the fitting."," We fit the sample with the following log-linear relation where $M_0$ is the median magnitude measured from the fitted sample, which is held constant during the fitting."
+" For the following fits, we do not include galaxies with only upper limits on their bulge magnitudes4."," For the following fits, we do not include galaxies with only upper limits on their bulge magnitudes."
+". Following and(2002),, we include the intrinsic scatter eo in the definition of x? (Equation 1 of 2006)) so that the best fit gives a X?=1."," Following and, we include the intrinsic scatter $\epsilon_0$ in the definition of $\chi^2$ (Equation 1 of ) so that the best fit gives a $\chi^2=1$."
+" Thus, we get an upper limit on the intrinsic scatter."," Thus, we get an upper limit on the intrinsic scatter."
+ Compared to the Bentz et al., Compared to the Bentz et al.
+" fit, we find a significantly steeper slope and lower zeropoint when our sample is included."," fit, we find a significantly steeper slope and lower zeropoint when our sample is included."
+ The difference between the two samples can be quantified by calculating the predicted BH masses based on our bulge luminosities and the best-fitting Mpreaictrelation for the Bentz et al., The difference between the two samples can be quantified by calculating the predicted BH masses $M_{\text{predict}}$ based on our bulge luminosities and the best-fitting relation for the Bentz et al.
+ sample., sample.
+" We calculate a mean difference of =1.150.03, where the latter is the (log(errorMpredict/Mpu)) on the mean."," We calculate a mean difference of $\langle \log(M_{\text{predict}}/\mbh) \rangle =
+1.15\pm 0.03$, where the latter is the error on the mean."
+" When we fit our combined samples, we find an intrinsic scatter of ερ= for an assumed scatter in virial BH masses of 0.48 to 0.42006)."," When we fit our combined samples, we find an intrinsic scatter of $\epsilon_0=0.41-0.45$ for an assumed scatter in virial BH masses of 0.48 to 0.4."
+" Both calculations 2010)show that(e.g., the galaxies in our sample have systematically smaller BH masses at a fixed bulge mass compared with the sample of Bentz et al.", Both calculations show that the galaxies in our sample have systematically smaller BH masses at a fixed bulge mass compared with the sample of Bentz et al.
+" Differences in stellar populations, which we estimate would dim our bulges by ~1 mag (Paper I), are not sufficient to explain these discrepancies."," Differences in stellar populations, which we estimate would dim our bulges by $\sim 1$ mag (Paper I), are not sufficient to explain these discrepancies."
+" To investigate further whether these differences represent a real physical difference (e.g., as a function of mass), we also compare to inactive BHs with dynamical BH masses from(2009)."," To investigate further whether these differences represent a real physical difference (e.g., as a function of mass), we also compare to inactive BHs with dynamical BH masses from."
+". As explained in Paper I, we adopt a color difference of V—J=1.34 mag to shift the Gülltekin et al."," As explained in Paper I, we adopt a color difference of $V-I=1.34$ mag to shift the Gülltekin et al."
+ sample from the V-band to the J-band., sample from the $V$ -band to the $I$ -band.
+ The average difference between the predicted and observed BH masses based on the Gülltekin et al., The average difference between the predicted and observed BH masses based on the Gülltekin et al.
+ fit is (log(Mpredict/Mpu))=0.85+0.04., fit is $\langle\log(M_{\text{predict}}/\mbh) \rangle=0.85\pm0.04$.
+ If we fit our sample combined with that of Gülltekin et al.," If we fit our sample combined with that of Gülltekin et al.,"
+ the intrinsic scatter increases significantly (see Table 1))., the intrinsic scatter increases significantly (see Table \ref{fittingresult}) ).
+ This scatter is driven by a systematic shift between our sample and the Gülltekin et al., This scatter is driven by a systematic shift between our sample and the Gülltekin et al.
+ sample., sample.
+" For a given BH mass, our galaxies have larger bulge luminosities."," For a given BH mass, our galaxies have larger bulge luminosities."
+" This shift may be real, or it may be explained by younger stars, and"," This shift may be real, or it may be explained by younger stars, and"
+negative. representing that the universe evolves from deceleration o acceleration.,"negative, representing that the universe evolves from deceleration to acceleration."
+ Different choice of 2 shows different evolution ustory., Different choice of $\beta$ shows different evolution history.
+ At any given time. the larger value of .? harmonizes the arger Hubble parameter and. the larger deceleration parameter after a certain time.," At any given time, the larger value of $\beta$ harmonizes the larger Hubble parameter and, the larger deceleration parameter after a certain time."
+ It is emphasized that larger value of <7 corresponds larger acceleration for our accelerated expanding universe at present time., It is emphasized that larger value of $\beta$ corresponds larger acceleration for our accelerated expanding universe at present time.
+ Furthermore. with ? increasing. the time for the transition from thegravity-dominated era to the dominated era is gradually late.," Furthermore, with $\beta$ increasing, the time for the transition from thegravity-dominated era to the DE-dominated era is gradually late."
+ In this section. we consider the constraints arising from the strong lensing observations on our gravity model.," In this section, we consider the constraints arising from the strong lensing observations on our gravity model."
+ We have calculated the number density of lensed quasars of the present model using the SQLS sample., We have calculated the number density of lensed quasars of the present model using the SQLS sample.
+" Specitically. from the differential lensing probability P(A@,1,). which is multiplying likelihoodwith respect to two variables. image separation A@ and 72, (determined by the redshift 2. of the source quasars). if we know the number distribution Αα), then the expected number density is given by and the expect lens number with image separations between ελ and 405 is computed by which can be used to compare with 11 lenses in the sample to constrain the parameters of the /(/7) model."," Specifically, from the differential lensing probability $P(\Delta\theta, R_s)$, which is multiplying likelihoodwith respect to two variables, image separation $\Delta\theta$ and $R_s$ (determined by the redshift $z_s$ of the source quasars), if we know the number distribution $N(z_s)$, then the expected number density is given by and the expect lens number with image separations between $\Delta\theta_1''$ and $\Delta\theta_2''$ is computed by which can be used to compare with 11 lenses in the sample to constrain the parameters of the $f(R)$ model."
+" Furthermore. it is important to point out that. according to the integral mean value theorem. we can also rewrite the expected lens number with image separation ranging from 4X6,"" to 40s"" as follows: where A@ is one exist value between AG; and NO»."," Furthermore, it is important to point out that, according to the integral mean value theorem, we can also rewrite the expected lens number with image separation ranging from $\Delta\theta_1{''}$ to $\Delta\theta_2{''}$ as follows: where $\Delta\theta$ is one exist value between $\Delta\theta_1$ and $\Delta\theta_2$."
+ The sample we used has 22683 quasars. and the redshifts range from 0.6 to 2.2.," The sample we used has 22683 quasars, and the redshifts range from 0.6 to 2.2."
+" To get Αίας}. we count the source quasar number for each redshift bin. with bin-size Ac,= 0.05."," To get $N(z_s)$ , we count the source quasar number for each redshift bin, with bin-size $\Delta z_s=0.05$ ."
+ We, We
+"where Figure 6 shows measurements of this quantity from the simulations, as a function of the smoothing scale adopted.","where Figure \ref{nd_combined_fig} shows measurements of this quantity from the simulations, as a function of the smoothing scale adopted."
+" The left panel is real units, the right panel is in errorbar units as in Figure 5.."," The left panel is real units, the right panel is in errorbar units as in Figure \ref{nd_sigma_distances_fig}."
+" By combining the two optimal levels the sensitivity slightly improves, but still remains at the lo level for fyi=100."," By combining the two optimal levels the sensitivity slightly improves, but still remains at the $\sigma$ level for $f_{\rm NL}=100$."
+ This is because the associated errorbars at different thresholds are correlated., This is because the associated errorbars at different thresholds are correlated.
+" Next, we consider the clustering strength of the excursion set regions."," Next, we consider the clustering strength of the excursion set regions."
+" We analyze in detail the case of v=2.00, where according to our previous findings the sensitivity to a local [ur type non-Gaussianity is maximized."," We analyze in detail the case of $\nu=2.00$, where according to our previous findings the sensitivity to a local $f_{\rm NL}$ type non-Gaussianity is maximized."
+ Figure 7 shows measurements of the hot and cold pixel correlations above/below threshold from the simulations., Figure \ref{clustering_fig} shows measurements of the hot and cold pixel correlations above/below threshold from the simulations.
+ ACDM (22)..," $\Lambda{\rm CDM}$ \citep{dunkley2009,komatsu2008}."
+ x photon diffusion aud the observed power is expected o be dominated by sources of foreground enission aud he interaction of the CMD with intervening structure., by photon diffusion and the observed power is expected to be dominated by sources of foreground emission and the interaction of the CMB with intervening structure.
+ Tn particular. the inverse Compton scattering of CMD yhotous off hot plasina bound to clusters of galaxies (7?) eives rise to a spectral distortion of the CAIB known as the Suuvaev-Zeldovich effect (SZE).," In particular, the inverse Compton scattering of CMB photons off hot plasma bound to clusters of galaxies \citep{sunyaev72} gives rise to a spectral distortion of the CMB known as the Sunyaev-Zel'dovich effect (SZE)."
+ At frequemcics ess than ~220 GIIz. the SZE produces a decrement iu he CMD intensity iu the direction of à galaxy. cluster.," At frequencies less than $\sim 220\,$ GHz, the SZE produces a decrement in the CMB intensity in the direction of a galaxy cluster."
+ Galaxy clusters produce (secondary) auisotropy power on arcnmünute scales correspouding to them augular size., Galaxy clusters produce (secondary) anisotropy power on arcminute scales corresponding to their angular size.
+ The observed SZE power depends seusitively ou he abundance of clusters aud the history of structure ormatiou., The observed SZE power depends sensitively on the abundance of clusters and the history of structure formation.
+ Iu particular. the auslitucde of the SZE power scales as of. where oy is the sius fluctuation of matter on scales of Sh+ Mpc. aud serves as an independeut xobe of the amplitude of density perturbations (?)..," In particular, the amplitude of the SZE power scales as $\sigma_8^7$, where $\sigma_8$ is the rms fluctuation of matter on scales of $8{\rm h}^{-1}\,$ Mpc, and serves as an independent probe of the amplitude of density perturbations \citep{komatsu2002}."
+ Evidence for small scale power bevoud that expectec roni the prunary CAIB has been reported bw the Berkeley Iinois Marvland Association (DIM) an Cosmic Backerouncd huager (CBI) interferometers operating at 30CITz.," Evidence for small scale power beyond that expected from the primary CMB has been reported by the Berkeley Illinois Maryland Association (BIMA) and Cosmic Background Imager (CBI) interferometers operating at $30\,$ GHz."
+" These mcasurcmeuts fud a leve of SZE anisotropy power consistent with a value of Jy, somewhat ereater than those preferred by other contemporary measurements. which favor cd,~0δ (2?).."," These measurements find a level of SZE anisotropy power consistent with a value of $\sigma_8$ somewhat greater than those preferred by other contemporary measurements, which favor $\sigma_8 \sim 0.8$ \citep{vikhlinin2009,komatsu2008}."
+ DIM observations at 00 (11 covering a tota of O.Ldee? of sky produced a nearly 26 detection of excess power in a flat baud centered at a multipole (=523; (7).," BIMA observations at $30\,$ GHz covering a total of $0.1\,{\rm deg}^2$ of sky produced a nearly $2\sigma$ detection of excess power in a flat band centered at a multipole $\ell=5237$ \citep{dawson2006}."
+ Due to the non-Gaussian distribution of the SZE on the skv aud the low significance of the detection. this resulted iu only weak coustraiuts ou the," Due to the non-Gaussian distribution of the SZE on the sky and the low significance of the detection, this resulted in only weak constraints on the"
+Some works have pointed out that the cucrey fiux associated to other wave modes nav be added up to balauce the chromospheric radiative losses.,Some works have pointed out that the energy flux associated to other wave modes may be added up to balance the chromospheric radiative losses.
+ Strausctal.(2008) detected low-frequcucy gravity waves in the Suns atmosphere. and claimed that their amount of energv is conrparable to the radiative losses of the chromosphere.," \citet{Straus+etal2008} detected low-frequency gravity waves in the Sun's atmosphere, and claimed that their amount of energy is comparable to the radiative losses of the chromosphere."
+ Ou the other haud. the simulated acoustic flux obtained by them is a factor of 10 smaller. consistent with Fossa&Carlsson (2005).," On the other hand, the simulated acoustic flux obtained by them is a factor of 10 smaller, consistent with \citet{Fossum+Carlsson2005}."
+.. Reearding the Alfvénn waves. dePouticuetal.(2007) claim to detect them in spicules in the upper chromosphere aud estimate the energy fiux carried bv these waves to be sufficient to accelerate the solar wind aud possibly to heat the quiet corona.," Regarding the Alfvénn waves, \citet{dePontieu+etal2007} claim to detect them in spicules in the upper chromosphere and estimate the energy flux carried by these waves to be sufficient to accelerate the solar wind and possibly to heat the quiet corona."
+ Most of the works that address the problem of upper atmosphere heating by waves refer to quiet Sun regions., Most of the works that address the problem of upper atmosphere heating by waves refer to quiet Sun regions.
+ Iu this paper we aim to study the cuerev balance iu a sunspot unubra., In this paper we aim to study the energy balance in a sunspot umbra.
+" Au earlier work iu this topic is the one by Ikueeretal.(1981).. who analize the time lags aud rns velocities between the lines acd aand found that the waves supply to the chromosphere an enoev flix of about 5ς10! ere 7 st. which is much lower than the chromospheric radiative losses of about 109 ere cur2 1,"," An earlier work in this topic is the one by \citet{Kneer+etal1981}, who analize the time lags and rms velocities between the lines and and found that the waves supply to the chromosphere an energy flux of about $5\times 10^4$ erg $^{-2}$ $^{-1}$, which is much lower than the chromospheric radiative losses of about $^6$ erg $^{-2}$ $^{-1}$."
+ Our approach consists in using 3D umuerical smmlatious to reproduce the oscillatory pattern from the photosphere to the chromosphere Oaned frou temporal series of spectropolarimetric data., Our approach consists in using 3D numerical simulations to reproduce the oscillatory pattern from the photosphere to the chromosphere obtained from temporal series of spectropolarimetric data.
+ The vesemblance between the simulated and observed waves allows us to extract conclusions from the munerical calculations provided that they correspond to a realistic phenomenon., The resemblance between the simulated and observed waves allows us to extract conclusions from the numerical calculations provided that they correspond to a realistic phenomenon.
+ In Section 2. we describe the observational data and the numerical method euiployed., In Section \ref{sect:procedures} we describe the observational data and the numerical method employed.
+ Sections 3. and Lb discuss the construction aud properties of the iiagnetolivdrostatie (MITIS) model of suuspot aud the procedure used to iutroduce the driver. respectively.," Sections \ref{sect:mhs_model} and \ref{sect:driver} discuss the construction and properties of the magnetohydrostatic (MHS) model of sunspot and the procedure used to introduce the driver, respectively."
+ The analysis of the simauulatious aud their comparison with the observations is performed iu Section 5.. while in Section 6 woe use these results to evaluate the energv balance.," The analysis of the simulations and their comparison with the observations is performed in Section \ref{sect:analysis}, while in Section \ref{sect:energy} we use these results to evaluate the energy balance."
+ Finally. the discussion aud. conclusions are preseuted iun Section 7..," Finally, the discussion and conclusions are presented in Section \ref{sect:conclusions}."
+" The observational dataconsists of a temporal series of co-spatial aud simultaneous data obtained with two different instruments. the POlarimetric Littrow Spectrograph (POLIS.Becketal.2005b) απ the Tenerife Iutrared Polavimeter ΤΠ.Colladosetal. 2007).. both attached to the Cerman Vacuuu Tower telescope at the Observatorio del Teide at Tenerife ou August 27°"" 2007,"," The observational dataconsists of a temporal series of co-spatial and simultaneous data obtained with two different instruments, the POlarimetric LIttrow Spectrograph \citep[POLIS,][]{Beck+etal2005b} and the Tenerife Infrared Polarimeter II \citep[TIP-II,][]{Collados+etal2007}, both attached to the German Vacuum Tower telescope at the Observatorio del Teide at Tenerife on August $^{th}$ 2007."
+ TIP-IL provides Stokes spectra of the 10830 iregion. including the À L0827 aan À 10830 spectral lines. with a spatial sampling of 0718 per pixel.," TIP-II provides Stokes spectra of the 10830 region, including the $\lambda$ 10827 and $\lambda$ 10830 spectral lines, with a spatial sampling of $0''18$ per pixel."
+ The blue channel of POLIS vields iuteusitv spectra of the line aud some plotospheric lline bleuds in its wines. with a spatial sampling of1 0729 per pixel.," The blue channel of POLIS yields intensity spectra of the line and some photospheric line blends in its wings, with a spatial sampling of $0''29$ per pixel."
+ The slit was placed over the center of a sunspot located near the ceuter of the Sun., The slit was placed over the center of a sunspot located near the center of the Sun.
+ The thorough analysis of this dataset was presented in Felipeetal.(2010b)., The thorough analysis of this dataset was presented in \citet{Felipe+etal2010b}.
+. We have performed 3D numerical simulations to reproduce the observed wave pattern from the photosphere to the chromosphere., We have performed 3D numerical simulations to reproduce the observed wave pattern from the photosphere to the chromosphere.
+ The umuerical wethod is described iu detail iu Felipeetal.(2010a)., The numerical method is described in detail in \citet{Felipe+etal2010a}.
+. The code solves the uoulinear ANID equatious for perturbations., The code solves the nonlinear MHD equations for perturbations.
+ A MIS ιούς. of suuspot is perturbed with a driver force iu the momentum equation., A MHS model of sunspot is perturbed with a driver force in the momentum equation.
+ The properties of the AUIS model aud the driver used iu the current work. are discussed in the following sections.," The properties of the MHS model and the driver used in the current work, are discussed in the following sections."
+ A perfectly matched laver (PAIL) boundary condition. (Bereneer1991) is applied to all boundaries in order to avoid wave reflections., A perfectly matched layer (PML) boundary condition \citep{Berenger1994} is applied to all boundaries in order to avoid wave reflections.
+" The radiative cucrey losses are implemented following Newtons cooling law: where Ti is the perturbation in the temperature. rg ds the radiative relaxation time. aud e, is the specific heat at constant volume."," The radiative energy losses are implemented following Newton's cooling law: where $T_1$ is the perturbation in the temperature, $\tau_R$ is the radiative relaxation time, and $c_v$ is the specific heat at constant volume."
+ Following Spiegel(1957). 7j Is given by where x is the mean absortion coefficient aud ap is the Stefan-Boltzimann constant.," Following \citet{Spiegel1957}, $\tau_R$ is given by where $\chi$ is the mean absortion coefficient and $\sigma_R$ is the Stefan-Boltzmann constant."
+ This expression is valid at photospheric heights. but not at the chromosphere as it was derived by Spicecl in the approximation of local thermodvuamic equilibrium.," This expression is valid at photospheric heights, but not at the chromosphere as it was derived by Spiegel in the approximation of local thermodynamic equilibrium."
+ As the values of τη given by the Spiegel formula are not correct at cliromosphleric heights. we took the freedom to modifv them in order to niic the low value of rg=10 s obtained fromm the observational analysis of the wave propagation at chromospheric heights iu Felipeetal.(2010b).," As the values of $\tau_R$ given by the Spiegel formula are not correct at chromospheric heights, we took the freedom to modify them in order to mimic the low value of $\tau_R=10$ s obtained from the observational analysis of the wave propagation at chromospheric heights in \citet{Felipe+etal2010b}."
+. The MITIS model is coustructed following the method bv Ishomenko&Collados(2008)., The MHS model is constructed following the method by \citet{Khomenko+Collados2008}.
+ These authors developed a method. to calculate a thick sunspot structure in maenetostatic equilibrimm with distributed currents;7c... showing a contiuuous variation of field streneth and sas pressure across the spot. from sub-photospleric to chromospheric lavers.," These authors developed a method to calculate a thick sunspot structure in magnetostatic equilibrium with distributed currents, showing a continuous variation of field strength and gas pressure across the spot, from sub-photospheric to chromospheric layers."
+ Iu curreut-distributed models the feld decreases from its value at the axis of the sunspot to almost zero at larec radial distances.," In current-distributed models, the field decreases from its value at the axis of the sunspot to almost zero at large radial distances."
+ These ieodels are coustructed by combining the advantages of two different methods: iu he deep lavers. the maguetic topology is set and the hermocdvunamic variables are forced to match with this structure (Schlüter&Temesvary1958:Low1975.1980).. constructing the so-called vselfsimilar” models: at shotospheric heights aud above. the pressure distribution is prescribed as boundary condition at the axis of the sunspot aud iu the distant quiet Sun atinosphliere and in )etwoeen the pressure and magnetic field are iteratively changed until the svstem reaches an equilibrium state (Pizzo1986).," These models are constructed by combining the advantages of two different methods: in the deep layers, the magnetic topology is set and the thermodynamic variables are forced to match with this structure \citep{Schluter+Temesvary1958, Low1975, Low1980}, constructing the so-called “self-similar” models; at photospheric heights and above, the pressure distribution is prescribed as boundary condition at the axis of the sunspot and in the distant quiet Sun atmosphere and in between the pressure and magnetic field are iteratively changed until the system reaches an equilibrium state \citep{Pizzo1986}."
+. The stratification of thermodvuamuc variables in the photosphere (needed for the boundary. couditious of the above methods) was retrieved from the inversion of the SStokes spectra iu the wmbra aud the quiet Sun atmosphere. averaged in time and space. using SIR (RuizCobo&delToroIniesta. 1992)..," The stratification of thermodynamic variables in the photosphere (needed for the boundary conditions of the above methods) was retrieved from the inversion of the Stokes spectra in the umbra and the quiet Sun atmosphere, averaged in time and space, using SIR \citep{RuizCobo+delToroIniesta1992}. ."
+ From the inversion of the SStokes profiles we can retrieve the stratification of magnetic field and thermodynamic variables up to, From the inversion of the Stokes profiles we can retrieve the stratification of magnetic field and thermodynamic variables up to
+estimates in a relatively narrow redshift range. vet one wide enough to include a large enough sample to craw meaningful conclusions.,"estimates in a relatively narrow redshift range, yet one wide enough to include a large enough sample to draw meaningful conclusions."
+" At 0.5«z<0.6 (Lable 3)). our LEG masses (?) average 0.047 dex higher than (7). for 1161. ""standard quasars (EPWILM. <1000. Kms). but 0.393 dex higher for 661 ANLs (1000 km/s)."," At $0.5 < z < 0.6$ (Table \ref{table:ourmasses}) ), our $\beta$ masses \citep{Vestergaard2006} average $0.047$ dex higher than \citep{Shen2011} for 1161 “standard” quasars (FWHM $< 1000$ km/s), but $0.393$ dex higher for 661 ANLs $> 1000$ km/s)."
+ More generally. the discrepancy between our unconstrained LL? masses and the ?δισ masses. which are based upon a line that does not change in ANLs compared to other objects. grows with increasing OUL width CFable 3)).," More generally, the discrepancy between our unconstrained $\beta$ masses and the \citet{Shen2008} masses, which are based upon a line that does not change in ANLs compared to other objects, grows with increasing ] width (Table \ref{table:ourmasses}) )."
+ For comparison. 111 and virial masses in the ? catalog show eood agreement for all quasars.," For comparison, $\beta$ and virial masses in the \citet{Shen2008} catalog show good agreement for all quasars."
+ Lf the uncertainty in comparing Le? ancl masses. is uncorrelated between cillerent quasars with a scatter of 0.22 dex (7). the svstematic error produced by restricting the narrow Lh? component to 1200. km/s would be the dominant source of error in ANL masses.," If the uncertainty in comparing $\beta$ and masses is uncorrelated between different quasars with a scatter of $\sim 0.22$ dex \citep{Shen2008}, the systematic error produced by restricting the narrow $\beta$ component to 1200 km/s would be the dominant source of error in ANL masses."
+ Alternative estimates for the statistical uncertainty in 1L2 alone include 0.4 dex (?) and 0.21 dex (?).. with the latter in addition to 0.15 dex scatter inMegll.," Alternative estimates for the statistical uncertainty in $\beta$ alone include 0.4 dex \citep{Vestergaard2006} and 0.21 dex \citep{Steinhardt2010c}, with the latter in addition to 0.15 dex scatter in."
+ Neither of these estimates investigated whether errors in LL? and masses are correlated. as they will be for errors emanating from the empirical relationship between monochromatic continuum luminosity ancl inferred. radius to the broacd-line region.," Neither of these estimates investigated whether errors in $\beta$ and masses are correlated, as they will be for errors emanating from the empirical relationship between monochromatic continuum luminosity and inferred radius to the broad-line region."
+ Even if the uncertainty were 0.4 dex and entirely statistical. the diserepaney. between 1.2 ancl mass estimates for ANLs would be at least comparable.," Even if the uncertainty were 0.4 dex and entirely statistical, the discrepancy between $\beta$ and mass estimates for ANLs would be at least comparable."
+ Thus. this systematic disagreement cannot be neglected for individual objects. Let alone in the analysis of large catalogs where the sample size allows a substantial reduction in the statistical error.," Thus, this systematic disagreement cannot be neglected for individual objects, let alone in the analysis of large catalogs where the sample size allows a substantial reduction in the statistical error."
+ We also note that the ? masses show much better agreement. between Lh? for ANLs. despite requiring as part of the fitting routine that the narrow component of LL? along with the Ol] doublet have FWIAL <1200 kmi/s. As above. an artilicially-low broad LL3 width: will xoducean artificiallv low virial mass MuxFWIIM.," We also note that the \citet{Shen2011} masses show much better agreement between $\beta$ for ANLs, despite requiring as part of the fitting routine that the narrow component of $\beta$ along with the ] doublet have FWHM $\leq 1200$ km/s. As above, an artificially-low broad $\beta$ width will producean artificially low virial mass $M_{\textrm{vir}} \propto \textrm{FWHM}^2$."
+ temarkablv. this lower LL3 mass turns out to be empirically almost exactly the amount needed to bring 111 ancMel mack into agreement: ? 123 masses average 0.094 dex higher or 5417 welloecasurecd “stancare™ quasars (ENIM «1000 kms). but 0.075. dex lower for 4716 well-measured ANLSs 1000 km/s).," Remarkably, this lower $\beta$ mass turns out to be empirically almost exactly the amount needed to bring $\beta$ and back into agreement: \citet{Shen2011} $\beta$ masses average $0.094$ dex higher for 5417 well-measured “standard” quasars (FWHM $< 1000$ km/s), but $0.075$ dex lower for 4716 well-measured ANLs $> 1000$ km/s)."
+" Similarly. at O5<2.0.6. ? 11 masses average 0.111 dex higher for 101. ""standard quasars (FWHAL <1000 km/s). but 0.065 dex lower for 661 ANLs 1000 km/s)."," Similarly, at $0.5 < z < 0.6$, \citet{Shen2011} $\beta$ masses average $0.111$ dex higher for 1161 “standard” quasars (FWHM $< 1000$ km/s), but $0.065$ dex lower for 661 ANLs $> 1000$ km/s)."
+ It should be noted that although Ll? ancl masses appear close to agreement for both standard and ANL quasars. the tvpical olfset has changed by ~0.17 dex between the two samples.," It should be noted that although $\beta$ and masses appear close to agreement for both standard and ANL quasars, the typical offset has changed by $\sim 0.17$ dex between the two samples."
+ Regardless. these are not based: upon the full ENIM of the LL? line.," Regardless, these are not based upon the full FWHM of the $\beta$ line."
+ As demonstrated by 2.. without a theoretical basis for choosing a mass estimator. cillercnt empirical calibration techniques will produce a wide variety of possible masses due to the very small number of reverberation mapping svstems available.," As demonstrated by \citet{Rafiee2011}, without a theoretical basis for choosing a mass estimator, different empirical calibration techniques will produce a wide variety of possible masses due to the very small number of reverberation mapping systems available."
+ Phere is a theoretical basis for using the entire line FWHIAL ancl assuming virial motion in order to determine the virial component of the broad-line region velocity as CyXEPWHAP., There is a theoretical basis for using the entire line FWHM and assuming virial motion in order to determine the virial component of the broad-line region velocity as $v_{\textrm{vir}} \propto \textrm{FWHM}^2$.
+ However. this relation now seems difficult to reconcile with the existence of ANLs even for 1. let alone at a smaller radius from the central black hole.," However, this relation now seems difficult to reconcile with the existence of ANLs even for $\beta$, let alone at a smaller radius from the central black hole."
+ It may be possible to develop a theoretical basis for reducing the measured. ENCIIM to a purely virial INLLM if the broadening can be well-mocdoeled., It may be possible to develop a theoretical basis for reducing the measured FWHM to a purely virial FWHM if the broadening can be well-modeled.
+ However. we currently lack that understanding. and there is no basis for using the 7. prescription for reducing the FWHILAL over any other that could be produced. from calibrating 11.1 masses againstMell for ANLs.," However, we currently lack that understanding, and there is no basis for using the \citet{Shen2011} prescription for reducing the FWHM over any other that could be produced from calibrating $\beta$ masses against for ANLs."
+ Virial mass estimation vieles. systematically cilferent results for ANLs when using Ll? and.Mgll. even though the effect was previously. masked. by constraints in the ? line fitting routine.," Virial mass estimation yields systematically different results for ANLs when using $\beta$ and, even though the effect was previously masked by constraints in the \citet{Shen2008} line fitting routine."
+ At least one of the two must be wrong. cither because of substantial non-virial motion in the broad line or because the empirical relationship between the radius to the broad-line region and continuum luminosity breaks down for ANLs.," At least one of the two must be wrong, either because of substantial non-virial motion in the broad line or because the empirical relationship between the radius to the broad-line region and continuum luminosity breaks down for ANLs."
+ Decause the Megll line does not appear to change with increasing OLLI]: while the broad. component of LL? is strongly. correlated. we might suspect that. the Il? masses are in error. perhaps due to a strong outllow.," Because the line does not appear to change with increasing ] while the broad component of $\beta$ is strongly correlated, we might suspect that the $\beta$ masses are in error, perhaps due to a strong outflow."
+we are not likely to know better than we know the distance to the centre of AISI.,we are not likely to know better than we know the distance to the centre of M31.
+ ὃν measuring the Galactic parallax to cach point along a stream. the streams three-dimensional trajectory could be determined. without. any knowledge of the Galaxys eravitational potential.," By measuring the Galactic parallax to each point along a stream, the stream's three-dimensional trajectory could be determined without any knowledge of the Galaxy's gravitational potential."
+ Phe requirement that this trajectory be an orbit must constrain the potential rather tightly., The requirement that this trajectory be an orbit must constrain the potential rather tightly.
+ In the previous section we have chosen to sacrifice some of the diagnostic power of equation by eliminating t on the grounds that it will be hard to measure., In the previous section we have chosen to sacrifice some of the diagnostic power of equation by eliminating $\vecthat{t}$ on the grounds that it will be hard to measure.
+ With t eliminated. the distance can only be recovered. by adopting a trial gravitational potential and searching over the fiducial distance sy.," With $\vecthat{t}$ eliminated, the distance can only be recovered by adopting a trial gravitational potential and searching over the fiducial distance $s_0$."
+ For sullicientlv small s. the radical in equation becomes imaginary. so there is a lower bound on the values ob sy that should be considered: below this bound the values of s obtained by solving equations. will somewhere approach the value at which the radical in equation becomes imaginary.," For sufficiently small $s$, the radical in equation becomes imaginary, so there is a lower bound on the values of $s_0$ that should be considered: below this bound the values of $s$ obtained by solving equations will somewhere approach the value at which the radical in equation becomes imaginary."
+ This event signals that the measured value of µ is too small to be consistent with the rellex of the Sun's velocity at the proposed distance., This event signals that the measured value of $\mu$ is too small to be consistent with the reflex of the Sun's velocity at the proposed distance.
+ Phus the kinematics of the problem. imposes a lower bound. on sy., Thus the kinematics of the problem imposes a lower bound on $s_0$.
+ Phere is no similar kinematic upper bound on sy because the radical in equation is real for all large s., There is no similar kinematic upper bound on $s_0$ because the radical in equation is real for all large $s$.
+ We have tested the ability of the algorithm to reconstruct orbits in the same Mivamoto-Nagai (1975) potential that was used in Paper LL namely with q=0.2.," We have tested the ability of the algorithm to reconstruct orbits in the same Miyamoto-Nagai (1975) potential that was used in Paper I, namely with $q=0.2$."
+ Phe numerical procedures were essentially the same as those described. in Paper Lo except. for one significant item: the fiducial point was placed in the micelle of the stream rather than at one of its ends. and equations were integrated in both directions away from this point.," The numerical procedures were essentially the same as those described in Paper I except for one significant item: the fiducial point was placed in the middle of the stream rather than at one of its ends, and equations were integrated in both directions away from this point."
+ This modification is advantageous as the data constrain derivatives of the observables much more tightly. at the middle of the stream than at its ends., This modification is advantageous as the data constrain derivatives of the observables much more tightly at the middle of the stream than at its ends.
+ The right panel of shows the track over the sky of an orbit that starts at. galactic [latitude b and a distance sy=Ίσα as viewed from the Sun. which is on a circular orbit at Ro=δα.," The right panel of shows the track over the sky of an orbit that starts at galactic latitude $b=-5^\circ$ and a distance $s_0=15a$ as viewed from the Sun, which is on a circular orbit at $R_0=8a$."
+ The left panel shows the magnitude of the proper motion along the orbit., The left panel shows the magnitude of the proper motion along the orbit.
+ The points in show the rms variation in the energy of orbits that are reconstructed from the 42 data points shown in for various assumed. distances s to the fiducial point and three trial potentials: the true. potential. which has scale-length ratio q=0.2 and two less Dlattened. potentials.," The points in show the rms variation in the energy of orbits that are reconstructed from the 42 data points shown in for various assumed distances $s$ to the fiducial point and three trial potentials: the true potential, which has scale-length ratio $q=0.2$ and two less flattened potentials."
+ The rms variation ind has a sharp minimum at the true distance when the true potential is used. and higher minima when the wrong potential is used.," The rms variation in $E$ has a sharp minimum at the true distance when the true potential is used, and higher minima when the wrong potential is used."
+ The exquisite precision with which the distance can be determined from this plot is remarkable: better than three parts in 107., The exquisite precision with which the distance can be determined from this plot is remarkable: better than three parts in $10^4$.
+ In the smallest. value reached by the ris energy variation is significantly higher than the corresponding ligure for the same orbit when line-of-sight velocities are used. (Paper D., In the smallest value reached by the rms energy variation is significantly higher than the corresponding figure for the same orbit when line-of-sight velocities are used (Paper I).
+ However. for other orbits smaller values of the ris variation in energy are obtained [rom proper-motion data: not surprisingly the data with the greatest diagnostic power varies with the nature of the orbit.," However, for other orbits smaller values of the rms variation in energy are obtained from proper-motion data; not surprisingly the data with the greatest diagnostic power varies with the nature of the orbit."
+ We have complemented the work of Paper L by. showing that when proper motions can be measured along a section of a single orbit. the full phase-space coordinates for the orbit can be reconstructed. as readily. as is the case when line-of-sight velocities have been measured.," We have complemented the work of Paper I by showing that when proper motions can be measured along a section of a single orbit, the full phase-space coordinates for the orbit can be reconstructed as readily as is the case when line-of-sight velocities have been measured."
+ However. we have also recovered a more powerful result: if the direction as well as the magnitude of the proper motions can be accurately measured. the three-dimensional geometry of the orbit can be recovered without assuming anvthing about he CGalaxv's gravitational potential.," However, we have also recovered a more powerful result: if the direction as well as the magnitude of the proper motions can be accurately measured, the three-dimensional geometry of the orbit can be recovered without assuming anything about the Galaxy's gravitational potential."
+ This reconstruction is possible because the Sun's velocity must. be responsible or the motion of stars perpendicular to. the orbit. so rom proper motion perpendicular to an orbit distances can be inferred. that are as fundamental as conventional rigonometric parallaxes.," This reconstruction is possible because the Sun's velocity must be responsible for the motion of stars perpendicular to the orbit, so from proper motion perpendicular to an orbit distances can be inferred that are as fundamental as conventional trigonometric parallaxes."
+ The extension of the work of Paper EL to proper motions is useful because proper motions can be readily measured for he main-sequence sars that define the stream., The extension of the work of Paper I to proper motions is useful because proper motions can be readily measured for the main-sequence stars that define the stream.
+ These stars are typically so faint zm:20) that it is hard to obtain line- velocities of the requisite accuracy for large numbers of them., These stars are typically so faint $I\gta20$ ) that it is hard to obtain line-of-sight velocities of the requisite accuracy for large numbers of them.
+ In Paper LE we show that even when linc-ol-sight velocities are available. a combination of observational errors and the fact that a tical stream coes not strictly follow an orbit make the task of determining whether a given potential is! compatible with an observed. stream complex.," In Paper II we show that even when line-of-sight velocities are available, a combination of observational errors and the fact that a tidal stream does not strictly follow an orbit make the task of determining whether a given potential is compatible with an observed stream complex."
+ Le proves necessary to search a multi-dimensional space of phase-space tracks that are compatible with the data for ones that are dynamically consistent orbits., It proves necessary to search a multi-dimensional space of phase-space tracks that are compatible with the data for ones that are dynamically consistent orbits.
+ A comparable search. would be required. when the data included. proper motions rather than line-of-sight velocities., A comparable search would be required when the data included proper motions rather than line-of-sight velocities.
+ In a future paper we will extend to proper motions the techniques for handling this problem. that are developed for line-of-sight velocities in Paper LL., In a future paper we will extend to proper motions the techniques for handling this problem that are developed for line-of-sight velocities in Paper II.
+ Another interesting avenue is to determine the solar motion by treating ib as an unknown when reconstructing orbits for several streams. using either line-of-sight velocities or proper motions. and choosing the value which allows a consistent interpretation of these streams.," Another interesting avenue is to determine the solar motion by treating it as an unknown when reconstructing orbits for several streams, using either line-of-sight velocities or proper motions, and choosing the value which allows a consistent interpretation of these streams."
+ Recently Ixoposovetal.(2009) used a combination of linc-of-sight. velocities ancl proper motions for the GD-1 stream to constrain Vo., Recently \cite{Koposov} used a combination of line-of-sight velocities and proper motions for the GD-1 stream to constrain $\vect{v}_0$.
+ Stronger constraints should be attainable by modelling several streams simultaneously., Stronger constraints should be attainable by modelling several streams simultaneously.
+ We thank the anonymous referee for his/her suggestions., We thank the anonymous referee for his/her suggestions.
+ AE acknowledges the support of PPARC/STEC during preparation of this work., AE acknowledges the support of PPARC/STFC during preparation of this work.
+their unique location in color-color space (?)..,their unique location in color-color space \citep{Steidel96}.
+ Morphological analvses of these Lyman Break Galaxies (LBGs) in the rest-frame UV have revealed (hat most of these svstems are disturbed and disk-like (1.e.. with exponential profiles). and onlv ~30% have proliles consistent. with ealactic spheroids (e.g..?2??7)..," Morphological analyses of these Lyman Break Galaxies (LBGs) in the rest-frame UV have revealed that most of these systems are disturbed and disk-like (i.e., with exponential profiles), and only $\sim 30\%$ have profiles consistent with galactic spheroids \citep[e.g.,][]{Ferguson04,Lotz06,Ravindranath06}."
+ Moreover. ? have found that the distribution of ellipticities derived from GALFIT are skewed towards higher values. and that this asymmetry becomes more pronounced al higher redshift.," Moreover, \citet{Ravindranath06} have found that the distribution of ellipticities derived from GALFIT are skewed towards higher values, and that this asymmetry becomes more pronounced at higher redshift."
+" This dominance of ""elongated"" morphologies seen in LBGs has been interpreted as evidence for galaxy formation within filamentary structure or the presence of proto-bars Chat span the entire visible disk of the galaxy.", This dominance of “elongated” morphologies seen in LBGs has been interpreted as evidence for galaxy formation within filamentary structure or the presence of proto-bars that span the entire visible disk of the galaxy.
+ Lya Emitters(LAEs) are a class of high-redshift galaxies which have not been studied as extensively as LBGs., $\alpha$ Emitters(LAEs) are a class of high-redshift galaxies which have not been studied as extensively as LBGs.
+ Like their Lyman break counterparts. LAEs between are actively star-forming objects (e.g...?).. and at the bright end. the sample of LAEs overlaps (hat of LBGs.," Like their Lyman break counterparts, LAEs between $2 < z < 4$ are actively star-forming objects \citep[e.g.,][]{CowieHu}, and at the bright end, the sample of LAEs overlaps that of LBGs."
+" However. a tvpical Lya emitting galaxy has a lower stellar mass (~10?AZ. ). higher mass-specilic star-Iormation rate (LOSAL. vyLU,F and lower dust content (E(D—V)9 0.1) (?2??) than any galaxy discovered via the Lyinan-break technique."," However, a typical $\alpha$ emitting galaxy has a lower stellar mass $\sim 10^9 M_{\sun}$ ), higher mass-specific star-formation rate $\sim 10^{-8} M_{\sun}$ $^{-1} M_{\sun}^{-1}$, and lower dust content $E(B-V) \sim 0.1$ ) \citep{Venemans05,Gawiser07,lai}
+ than any galaxy discovered via the Lyman-break technique."
+ Aloreover. LAEs are less clustered than LDGs. suggesting a lower dark-matter halo mass. and their bias factor 5b~ 1.7) consistent with Chat expected from the progenitors of present-dav L* galaxies (??)..," Moreover, LAEs are less clustered than LBGs, suggesting a lower dark-matter halo mass, and their bias factor $b \sim 1.7$ ) consistent with that expected from the progenitors of present-day $L^*$ galaxies \citep{Gawiser07,guaita}."
+ To date. relatively few LAEs have been studied morphologically.," To date, relatively few LAEs have been studied morphologically."
+ In the rest-frame UV. most 3S2€6 LAEs are small (<1 kpe). compact (C> 2.5). and barely resolved at resolution (?22??)..," In the rest-frame UV, most $3 \lesssim z \lesssim 6$ LAEs are small $\lesssim 1$ kpc), compact $C>2.5$ ), and barely resolved at resolution \citep{Venemans05,Pirzkal07,Overzier08,taniguchi}."
+ Those LAEs that are not compact (~20— 457)) exhibit chimpy or iregular morphology. with diffuse components extending to several kiloparsecs.," Those LAEs that are not compact $\sim 20 - 45$ ) exhibit clumpy or irregular morphology, with diffuse components extending to several kiloparsecs."
+ Due to their small sizes and low luminosities. however. categorizing (hese objects within the context ol existing classification schemes has been difficult.," Due to their small sizes and low luminosities, however, categorizing these objects within the context of existing classification schemes has been difficult."
+ lere we present a study. of the vest-lrame ultraviolet morphology of a sample of z=3.1 LAEs (?) present on archival V-band images from the Galaxy Evolution from Morphology and SEDS (GEMS.?).. Great Observatories Origins Deep Survev (GOODS.?).. and Hubble Uliradeep Field GIUDF.?) surveys.," Here we present a study of the rest-frame ultraviolet morphology of a statistically-complete sample of $z=3.1$ LAEs \citep{GronwallLAE} present on archival $V$ -band images from the Galaxy Evolution from Morphology and SEDS \citep[GEMS,][]{GEMS}, Great Observatories Origins Deep Survey \citep[GOODS,][]{GOODS}, and Hubble Ultradeep Field \citep[HUDF,][]{HUDF} surveys."
+ In Paper I (?).. we investigated the size distribution of LAEs bv developingan analvsis pipeline which identifies each component of an LAE. and measures its size and brightness.," In Paper I \citep{Bond09}, we investigated the size distribution of LAEs by developingan analysis pipeline which identifies each component of an LAE, and measures its size and brightness."
+ We found that most LAEs are extremely compact. with hall-light radii. r.«2 kpe. and that a S/N>30 is required [or a robust measure of morphology.," We found that most LAEs are extremely compact, with half-light radii, $r_e < 2$ kpc, and that a $S/N > 30$ is required for a robust measure of morphology."
+ Here. we extend this study. ancl provide more detailed morphological information on each object.," Here, we extend this study, and provide more detailed morphological information on each object."
+ In 2 and 3.. we describe (he data ancl detail the algorithms used in our morphology-fitàng pipeline analvsis.," In \ref{sec:data} and \ref{sec:method}, we describe the data and detail the algorithms used in our morphology-fitting pipeline analysis."
+ In 4.. we present the best-fitting morphological parameters for each galaxy. and explore how (hese parameters vary with image depth.," In \ref{sec:results}, we present the best-fitting morphological parameters for each galaxy, and explore how these parameters vary with image depth."
+ Finally. in 5.. we discuss the implications of our findings.," Finally, in \ref{sec:discussion}, , we discuss the implications of our findings."
+" Throughout (his paper. we will assume a concordance cosmology with 4/4,=T1 kms ! |. Q,,= 0.27."," Throughout this paper, we will assume a concordance cosmology with $H_0=71$ km $^{-1}$ $^{-1}$ , $\Omega_m=0.27$ ,"
+"Spectropolarimetry is a. ""photon-hungry"" observational technique. requiring a particular sequence of observations to completely and independently determine the polarization properties of astronomical objects.","Spectropolarimetry is a “photon-hungry” observational technique, requiring a particular sequence of observations to completely and independently determine the polarization properties of astronomical objects."
+ In order to measure particularly low levels of polarization (< 0.1%). high signal-to-noise ratio (S/N) observations are required at each step in the observing For dual-beam spectropolarimeters. such as the European Southern Observatory (ESO) Very Large Telescope (VLT) FOcal Reducer and low dispersion Spectrograph. (FORSI: Appenzelleretal. 1998)). the determination. of the linear polarization Stokes parameters. Q and U. requires N=4 separate observations.," In order to measure particularly low levels of polarization $\lesssim 0.1\%$ ), high signal-to-noise ratio (S/N) observations are required at each step in the observing For dual-beam spectropolarimeters, such as the European Southern Observatory (ESO) Very Large Telescope (VLT) FOcal Reducer and low dispersion Spectrograph (FORS1; \citealt{1998Msngr..94....1A}) ), the determination of the linear polarization Stokes parameters, $Q$ and $U$, requires N=4 separate observations."
+ This observing sequence introduces redundancies. but permits the complete determination and correction for instrumental effects that would. if uncorrected for N=2 observations. lead to spurious observed Target of Opportunity spectropolarimetry (where observations are conducted at specific epochs to study transient events which occur without prior warning) of time-variable phenomena is at a disadvantage. since at each epoch the technique requires a factor ~8 more time on target than for pure spectroscopy to achieve the same S/N 11 the individual spectra.," This observing sequence introduces redundancies, but permits the complete determination and correction for instrumental effects that would, if uncorrected for N=2 observations, lead to spurious observed Target of Opportunity spectropolarimetry (where observations are conducted at specific epochs to study transient events which occur without prior warning) of time-variable phenomena is at a disadvantage, since at each epoch the technique requires a factor $\sim 8$ more time on target than for pure spectroscopy to achieve the same S/N in the individual spectra."
+ For objects such as SNe. the position of the target on the sky may not permit sufficient time on target to conduct all of the necessary observations. leaving the dataset acquired. incomplete.," For objects such as SNe, the position of the target on the sky may not permit sufficient time on target to conduct all of the necessary observations, leaving the dataset acquired incomplete."
+ If one observation. for one retarder plate angle. is absent. only one Stokes parameter can be independently measured.," If one observation, for one retarder plate angle, is absent, only one Stokes parameter can be independently measured."
+ There is. however. potential redundancy between the observations for the completely determined Stokes parameter and the partially determined parameter such that instrumental effects can be removed. allowing for complete determination of both parameters. but at a higher degree of Leonardetal.(2001) present a technique where. by assuming the same ratio of gains for the ordinary and extraordinary rays. instrumental effects can be removed to calculate the second Stokes parameter for incomplete datasets.," There is, however, potential redundancy between the observations for the completely determined Stokes parameter and the partially determined parameter such that instrumental effects can be removed, allowing for complete determination of both parameters, but at a higher degree of \citet{2001ApJ...553..861L} present a technique where, by assuming the same ratio of gains for the ordinary and extraordinary rays, instrumental effects can be removed to calculate the second Stokes parameter for incomplete datasets."
+ When “normalized flux differences” (e.g.Jehinetal.2005:Patat&Romaniello2006) are used to calculate the Stokes parameter. the consideration is of a correction for the instrumental signature. Which quantifies the deviation of data from the ideal case of Here we present a technique to determine the instrumental signature corrections for observed data. such that. 1n the event of an incomplete observation. both Stokes parameters can be determined.," When “normalized flux differences” \citep[e.g.][]{forsman,2006PASP..118..146P} are used to calculate the Stokes parameter, the consideration is of a correction for the instrumental signature, which quantifies the deviation of data from the ideal case of Here we present a technique to determine the instrumental signature corrections for observed data, such that, in the event of an incomplete observation, both Stokes parameters can be determined."
+ The concept of spectropolarimetry with dual beam spectropolarimeters 1s. briefly. outlined along with discussion of correction for instrumental effects in refsection:basicconcept..," The concept of spectropolarimetry with dual beam spectropolarimeters is, briefly, outlined along with discussion of correction for instrumental effects in \\ref{section:basicconcept}."
+ In refsection:models.. analytical and Monte-Carlo models of a dual beam spectropolarimeter are presented. and the effects of various factors on the final measured polarization for complete and incomplete datasets are presented.," In \\ref{section:models}, analytical and Monte-Carlo models of a dual beam spectropolarimeter are presented, and the effects of various factors on the final measured polarization for complete and incomplete datasets are presented."
+ Real observations of linearly polarized point sources. acquired using VLT FORS1. and the instrumental signature corrections are presented in refsection:observations..," Real observations of linearly polarized point sources, acquired using VLT FORS1, and the instrumental signature corrections are presented in \\ref{section:observations}."
+ Dual-beam spectropolarimeters. such as the FORSI instrument (Appenzelleretal.1998).. use a sequence of a retarder plate and a Wollaston prism to measure the polarization components.," Dual-beam spectropolarimeters, such as the FORS1 instrument \citep{1998Msngr..94....1A}, use a sequence of a retarder plate and a Wollaston prism to measure the polarization components."
+ Rotation of the retarder plate varies the angle at which the orthogonal polarization components are sampled., Rotation of the retarder plate varies the angle at which the orthogonal polarization components are sampled.
+ The Wollaston prism separates these two components spatially into the ordinary (0) and extraordinary rays (ο)., The Wollaston prism separates these two components spatially into the ordinary ${o}$ ) and extraordinary rays ${e}$ ).
+" At a general retarder plate position ;. themeasured normalized flux difference F7 is defined (e.g.Jehinetal.2005) as such that the total intensity 7=f,+fo."," At a general retarder plate position $i$, the normalized flux difference $F^{m}_{i}$ is defined \citep[e.g.][]{forsman} as such that the total intensity $I=f_{o}+f_{e}$."
+" F7? is normalized by the total flux intensity and is independent of varying sky transparency and exposureThe optimum observing sequence is for a half-wavelength retarder plate to be positioned at N=4 position angles 6;: 00:00. 8,222255. 8,245700 and 063267755."," $F^{m}_{i}$ is normalized by the total flux intensity and is independent of varying sky transparency and exposureThe optimum observing sequence is for a half-wavelength retarder plate to be positioned at N=4 position angles $\theta_{i}$: $\theta_{0}$ 0, $\theta_{1}$ 5, $\theta_{2}$ 0 and $\theta_{3}$ 5."
+ At @. the orthogonal polarization components are 1n the horizontal and vertical directions.," At $\theta_{0}$, the orthogonal polarization components are in the horizontal and vertical directions."
+ For 8» the same components are observed. but the beams are swapped such that fi= fi.," For $\theta_{2}$ the same components are observed, but the beams are swapped such that $f^{i}_{o}=f^{i+2}_{e}$ ."
+" Similarly. for 4, and 43 the diagonal polarization components are observed. and the polarization components observed as the"," Similarly, for $\theta_{1}$ and $\theta_{3}$ the diagonal polarization components are observed, and the polarization components observed as the"
+or more outbursts.,or more outbursts.
+ According to Poisson statistics. if the outburst rate (ποσο we take + to have units such that the expected umber of outbursts ina time Tis £T). the chance that any eiven source is detected will be Given au eusenible of a sources. the probability that m sources will be detected (and a0i not detected) is This is simply a binondal distribution. which cau be inverted using Baves’ theorems to obtain τονeT) in the same wav as Equation 9..," According to Poisson statistics, if the outburst rate $r$ (where we take $r$ to have units such that the expected number of outbursts ina time $T$ is $rT$ ), the chance that any given source is detected will be Given an ensemble of $n$ sources, the probability that $m$ sources will be detected (and $n-m$ not detected) is This is simply a binomial distribution, which can be inverted using Bayes' theorem to obtain ${\rm prob}(n|m,rT)$ in the same way as Equation \ref{eq:bayes}."
+ We can receive some euidauce about the rate at which SCRs are active by notius that. out of 25 vears of monitoring. SCR 180620 has been active for z7 wears. SCR 1900|L1 for = vears. SCR 052666 (n the LAIC) for z2 vears. aud SCRLl 1627.11 for +2 mouths (Aptelaral.2001:Woods&Thompson 2006).," We can receive some guidance about the rate at which SGRs are active by noting that, out of 25 years of monitoring, SGR 1806–20 has been active for $\approx$ 7 years, SGR 1900+14 for $\approx$ 4 years, SGR 0526–66 (in the LMC) for $\approx$ 2 years, and SGR 1627–41 for $\approx$ 2 months \citep{apt01,wt06}."
+. The bursts tend to be associated with periods of activity lasting months to vears., The bursts tend to be associated with periods of activity lasting months to years.
+" It is clear that the activity level is variable frou source to source. so We consider a range of rates,"," It is clear that the activity level is variable from source to source, so we consider a range of rates."
+ Cüveu m=3 and T—25 vears. in Figure 9 we plot for a range of r the most-likely value and coufidence lit for ».," Given $m$ =3 and $T$ =25 years, in Figure \ref{fig:numsgrs} we plot for a range of $r$ the most-likely value and confidence limit for $n$."
+ If on average one outburst is expected every LO vears. which is near the mean of the sample that has already been ideutified. then with confideuce we should have already. detected all of the Calactic SCRs.," If on average one outburst is expected every 10 years, which is near the mean of the sample that has already been identified, then with confidence we should have already detected all of the Galactic SGRs."
+ However. if an SCR ouly becomes active once per 100 years. there could be 13! in total in the Galaxy," However, if an SGR only becomes active once per 100 years, there could be $13^{+18}_{-7}$ in total in the Galaxy."
+ Our best estimate is therefore ISthat the number of Galactic SGRs is much siualler than that of standard ANPs. uuless SCRs are active less frequently than once per several lauded vears.," Our best estimate is therefore that the number of Galactic SGRs is much smaller than that of standard AXPs, unless SGRs are active less frequently than once per several hundred years."
+ Transicut ANPs would be the most difficult to ideutifv. and could comprise the majority of the maguetars.," Transient AXPs would be the most difficult to identify, and could comprise the majority of the magnetars."
+ We define an ANP as trausient if it has a luminosity of C107? ((0.510 keV) for most of its current life. but iucreases in huninosity to ~10° ffor durations of on order a vear (e.g.Cotthelfetal.2001:Tamctal.2006).," We define an AXP as transient if it has a luminosity of $\la$$10^{33}$ (0.5–10 keV) for most of its current life, but increases in luminosity to $\sim$$10^{35}$ for durations of on order a year \citep[e.g.,][]{got04,tam06}."
+. We define the voutburst™ as the vear-loug period when the ANP is luminous., We define the “outburst” as the year-long period when the AXP is luminous.
+ Three transient AXPs are known (XTE Jis10197. ΝΟ J161710.2.155216. and LE 1517.05108). alone with one candidate (AN J1815.0.0258).," Three transient AXPs are known (XTE J1810–197, CXOU J164710.2–455216, and 1E 1547.0–5408), along with one candidate (AX J1845.0–0258)."
+ Qur archival aand ssurvey provides the most scusitive search for trausieut ANPs in quiescence., Our archival and survey provides the most sensitive search for transient AXPs in quiescence.
+ Nonetheless. when looking for a pulsar with Z«—3 «10 ((0.510 keV) aud «10∙≺∏∐⋅↴∖↴↿∐⋅↖⇁↸∖⋅↖⇁↕↴∖↴∪↕∐∙↖↽≝⋊∣↖⋰⋰⊒∶↙ coniplete for 0. of the GalacticL vouug stellar population.," Nonetheless, when looking for a pulsar with $L_{\rm X}$ $3\times10^{33}$ (0.5–10 keV) and $A_{\rm rms}$, our survey is only complete for $\approx$ of the Galactic young stellar population."
+ The simple calculation (Eq. 9)), The simple calculation (Eq. \ref{eq:bayes}) )
+ assumniug no new barely-detectable laenetars n f=0.005 of the implies that with confidence there could be up toHL 510 waiting to be discovered., assuming no new barely-detectable magnetars in $f$ =0.005 of the implies that with confidence there could be up to 540 waiting to be discovered.
+" Uufortunatelv. this estimate is not a strict upper bound. because quiescent ANPs could be less ""umous or have lower pulsed fractious than we have assunied."," Unfortunately, this estimate is not a strict upper bound, because quiescent AXPs could be less luminous or have lower pulsed fractions than we have assumed."
+ Au ANP in outburst could be detected over a larger raction of the Galaxy. but it still requires an observation with a pointed XN-rav observatory.," An AXP in outburst could be detected over a larger fraction of the Galaxy, but it still requires an observation with a pointed X-ray observatory."
+ To coustrain their nuubers based on ANPs in outburst. we must account or both the fraction of the vouug stars in the Galaxy hat can be surveyed by a given observatory (Eq. 8)).," To constrain their numbers based on AXPs in outburst, we must account for both the fraction of the young stars in the Galaxy that can be surveyed by a given observatory (Eq. \ref{eq:binomial}) ),"
+ and the chance that AXPs in the reeion will produce an outburst (Eq 11))., and the chance that AXPs in the region will produce an outburst (Eq \ref{eq:poissprob}) ).
+ We assume that there is a population of N inaenetars in the Galaxy with an outburst rate r. and that we searched for them with a survey of duration T that covered a fraction f of the magnetar birth places.," We assume that there is a population of $N$ magnetars in the Galaxy with an outburst rate $r$, and that we searched for them with a survey of duration $T$ that covered a fraction $f$ of the magnetar birth places."
+ Then. the probability that i maguetars will be found is the joiut probability that η out of No maguetars will lic in the survey region. and that am out of à will produce outbursta: ↽∕∏∐↴∖↴↸⊳⋜⋯↴⋝↸∖↕∐↖↽↸∖↥⋅↑↸∖≺∏↴∖↴↕∐∶↴∙⊾↕≧⋜↧⋅↖↽↸∖↴∖↴⋅↑↕∐∖∪↥⋅↸∖⋯∙↕∐∪↥⋅≼∐∖↥⋅↑∪ ≼∐∖↥⋅↕↖↽↸∖↑∐↸∖↻↥⋅∪↴⋝⋜∏⋝∐↕↑⋅↖↽↑↕⋯↑↑↕∐∖↥⋅↸∖⋜⋯∖⊀∖⊽↕⊔⋜↧∶↴∙⊾∐↸∖↑⋜∐⋅↴∖↴∶↴∙⊾↕↖⇁↸∖∐ that i) are fouud with outbursts at arate rin a fraction f of the Galaxy. probCNimi.rT.f). as in Equation 9..," Then, the probability that $m$ magnetars will be found is the joint probability that $n$ out of $N$ magnetars will lie in the survey region, and that $m$ out of $n$ will produce outbursts: This can be inverted using Bayes' theorem, in order to derive the probability that there are $N$ magnetars given that $m$ are found with outbursts at a rate $r$ in a fraction $f$ of the Galaxy, ${\rm prob}(N|m,rT,f)$, as in Equation \ref{eq:bayes}. ."
+ Tere. because we are dealing with transicut sources. f represents the fraction of the voung Calactic population searched per vear. which is the duration of a typical outburst.," Here, because we are dealing with transient sources, $f$ represents the fraction of the young Galactic population searched per year, which is the duration of a typical outburst."
+ We can estimate the rate of outbursts from trausicut ANPs from the four known examples (acre we assume, We can estimate the rate of outbursts from transient AXPs from the four known examples (here we assume
+adopt the popular values that Hocdriguezοἱal(1995);RodriguezandMirabel(1999) found io be consistent with the observations of the flare from Mareh 1994. Combining the results of equations (1) through (7) with the fits in Table 2. we have two equations (constraints) and four unknowns (o is known through Table 2).,"adopt the popular values that \citet{rod95,rod99} found to be consistent with the observations of the flare from March 1994, Combining the results of equations (1) through (7) with the fits in Table 2, we have two equations (constraints) and four unknowns $\alpha$ is known through Table 2)."
+ The constraints are eiven by frequency of the spectral peak and peak fIux density (equivalently. (he normalization of the background svncehrotvon power law and the Irequency of unit optical depth defined by pio)= 1).," The constraints are given by frequency of the spectral peak and peak flux density (equivalently, the normalization of the background synchrotron power law and the frequency of unit optical depth defined by $\mu({\nu})R=1$ )."
+" The unknowns can be described by 5. the magnetic field. Np. the normalization of the energy distribution of leptons. £,,;,. the minimum energy of the lepton energy. power law and A. (he radius of the spherical plasmoid."," The unknowns can be described by $B$ , the magnetic field, $N_{\Gamma}$, the normalization of the energy distribution of leptons, $E_{min}$, the minimum energy of the lepton energy power law and $R$, the radius of the spherical plasmoid."
+ The solution space is illustrated eraphically., The solution space is illustrated graphically.
+" Figures 5 and 6 are plots of B and ΑΝ. the total number of particles in the plasmoil. as a [function of Ro for each value of £5, lor the three components on December 6. 1993. where Five representative values of £2,,5, ave considered in (he family of models presented in the following that are equal to 1. 5. 10. 20 and 30."," Figures 5 and 6 are plots of $B$ and $N$, the total number of particles in the plasmoid, as a function of R for each value of $E_{min}$ for the three components on December 6, 1993, where Five representative values of $E_{min}$ are considered in the family of models presented in the following that are equal to 1, 5, 10, 20 and 30."
+" The integral in equation (8) is insensitive to the value of E,,,, [or large n as long as E,2»E.", The integral in equation (8) is insensitive to the value of $E_{max}$ for large $n$ as long as $E_{max}\gg E_{min}$.
+" A liberal value of E,,,,=LO? was chosen.", A liberal value of $E_{max} = 10^{6}$ was chosen.
+ Reasonable variations in this choice will make modest changes to the energy in C3. since η=2 is not large.," Reasonable variations in this choice will make modest changes to the energy in C3, since $n=2$ is not large."
+ However. (he energy in (his component appears to be very reduced compared to the other components. regardless of this choice.," However, the energy in this component appears to be very reduced compared to the other components, regardless of this choice."
+ The models of C3 are the most poorly constrained because of the sparse high frequency coverage., The models of C3 are the most poorly constrained because of the sparse high frequency coverage.
+ The (rends (o note are that D always increases dramatically as the radius of the model of the plasmoid increases for all three components., The trends to note are that $B$ always increases dramatically as the radius of the model of the plasmoid increases for all three components.
+" Furthermore. 2. is not a Iunction of £,,;,."," Furthermore, $B$, is not a function of $E_{min}$."
+ By contrast. ;/N. decreases dramatically as (he radius of the plasmoid increases.," By contrast, $N$ decreases dramatically as the radius of the plasmoid increases."
+ The total number of particles is very dependent on [τμ especially for CI and C2 that have the steep high frequency power laws (ecuivalentlv. a steep lepton energv distribution).," The total number of particles is very dependent on $E_{min}$ especially for C1 and C2 that have the steep high frequency power laws (equivalently, a steep lepton energy distribution)."
+ Ht is interesting to note the size of the largest component. Cl.," It is interesting to note the size of the largest component, C1."
+ The highest resolution ever used to view an ejected plasmoid in a major flare are the 2.0 em VLBA observation of Dhawanetal(2000)., The highest resolution ever used to view an ejected plasmoid in a major flare are the 2.0 cm VLBA observation of \citet{dha00}.
+. It is not clear from Dhawanetal(2000) if the 2 em observations reveal the emitted component il the data is nol tapered., It is not clear from \citet{dha00} if the 2 cm observations reveal the emitted component if the data is not tapered.
+ The VLBA observations at 2.0 em are capable of resolving components 2>x10H em., The VLBA observations at 2.0 cm are capable of resolving components $R > 2 \times 10^{14}$ cm.
+ Thus. only the largest models of CL could be marginally resolved by the VLBA and C2 and C3 could not be resolved on December 6.The Dhawanetal(2000) observations of two different flares were unable to resolve (he ejected region of hieh emissivity in the image plane," Thus, only the largest models of C1 could be marginally resolved by the VLBA and C2 and C3 could not be resolved on December 6.The \citet{dha00} observations of two different flares were unable to resolve the ejected region of high emissivity in the image plane"
+the metastable state. then thought to be the dominaut method of depopulation.,"the metastable state, then thought to be the dominant method of depopulation."
+ This 2-photon decay proceeds on a timescale of At2116days (Drake&Dal-earno LOG68)., This 2-photon decay proceeds on a timescale of $A^{-1}\gtrsim 116$days \citep{dra68}.
+. Iu the following vear. the radiative lifetime associated with I-photon decay was computed to be AtomY hows (πω1969).. nearly LOO times faster than the 2-photon decay.," In the following year, the radiative lifetime associated with 1-photon decay was computed to be $A^{-1}\approx 7$ hours \citep{gri69}, nearly 400 times faster than the 2-photon decay."
+ This value was later refined to A+=2.5 hours (Woodworth&Moos1975:Hata&Carant 1981).. still much faster. aud as a result the analyses performed by Scherb(1968) aud Reesetal.(1968) had ereatly overestimated the population in the metastable state (thispossibilitywasnotedbyReesetal. 1968).," This value was later refined to $A^{-1}=2.5$ hours \citep{woodworth75,hata81}, still much faster, and as a result the analyses performed by \citet{sch68} and \citet{ree68} had greatly overestimated the population in the metastable state \citep[this possibility was noted by][]{ree68}."
+. Consequeuth. the thought of observiug interstellar metastable helium was abaudoned.," Consequently, the thought of observing interstellar metastable helium was abandoned."
+ Because of the high ionization rate inferred. from.. we decided to revisit these calculations considering up-to-date rate coeffiicieuts aud improved telescope/detector capabilities.," Because of the high ionization rate inferred from, we decided to revisit these calculations considering up-to-date rate coefficients and improved telescope/detector capabilities."
+" Asstuine the same chemical scheme as in the past. we can derive the steady state equations for the eround. ionized. aud metastable states: IHere. ng. n;. aud a, denote the populations of the eround. ionized. aud metastable levels. respectively: Gn... the total ionization rate of He"" atoms due to cosmic raves. including the effects of secondary electrons: 04 and 3o3. the total. direct recombination rates to all singlet levels aud to all triplet levels. respectively: ος. the electron deusitv in the eas: and <A. the Einstein cocfiicicut for spontancous emission frou the iietastable level."," Assuming the same chemical scheme as in the past, we can derive the steady state equations for the ground, ionized, and metastable states: Here, $n_{g}$, $n_{i}$, and $n_m$ denote the populations of the ground, ionized, and metastable levels, respectively; $\zeta_{\rm He}$, the total ionization rate of $^0$ atoms due to cosmic rays, including the effects of secondary electrons; $\alpha_1$ and $\alpha_3$, the total, direct recombination rates to all singlet levels and to all triplet levels, respectively; $n_e$, the electron density in the gas; and $A$, the Einstein coefficient for spontaneous emission from the metastable level."
+" The units of cach term in equations (1)) through 3) are P 1,", The units of each term in equations \ref{eqssg}) ) through \ref{eqssm}) ) are $^{-3}$ $^{-1}$.
+ All reconibinations. to triplet levels above the inetastable level are assumed το produce subsequent radiative cascades to the metastable level that are effectively instantancous. owing to the long lifetime of the latter level.," All recombinations to triplet levels above the metastable level are assumed to produce subsequent radiative cascades to the metastable level that are effectively instantaneous, owing to the long lifetime of the latter level."
+ Similarly. all recombinations to sinelet levels are assumed to cascade promptly to the eround level," Similarly, all recombinations to singlet levels are assumed to cascade promptly to the ground level."
+" Equatious (2)) aud (3)) cau be solved for the ratios n; aud 7n; aud thus for i,Πο. well."," Equations \ref{eqssi}) ) and \ref{eqssm}) ) can be solved for the ratios $n_i/n_g$ and $n_m/n_i$, and thus for $n_m/n_g$ as well."
+" These ratios cau be substituted iuto the definition of the fractional population of the metastable level. in order to obtain the desired relation between the fractional metastable population fj, aud the ionization rate Gn. The triplet brauchiug fraction for recombinatious at TOW is b=aslfay|03)0.62. since ay=LO«LO ome bandas =6.6.10 ems 1 (BR. Porter 2009. private conuuunication)."," These ratios can be substituted into the definition of the fractional population of the metastable level, in order to obtain the desired relation between the fractional metastable population $f_m$ and the ionization rate $\zeta_{\rm He}$, The triplet branching fraction for recombinations at 70 K is $b=\alpha_3/(\alpha_1+\alpha_3)=0.62$, since $\alpha_1=4.0\times10^{-12}$ $^3$ $^{-1}$ and $\alpha_3=6.6\times10^{-12}$ $^3$ $^{-1}$ (R. Porter 2009, private communication)."
+ Radiative decay. to the eround level is by far the fastest of the three processes mentioned above. with 42911«10's +.," Radiative decay to the ground level is by far the fastest of the three processes mentioned above, with $A=1.1\times10^{-4}$ $^{-1}$."
+" In coutrast. yn=0.02 απὃν ogni,=1.3.10D sto and cy. bare representative values in diffuse clouds."," In contrast, $n_e=0.02$ $^{-3}$, $\alpha_3n_e=1.3\times10^{-13}$ $^{-1}$, and $\zeta_{\rm He}=3\times10^{-16}$ $^{-1}$ are representative values in diffuse clouds."
+ Caiven these values.1Atos.)ASC). aud equation (5)) cau be approximated by Owing to the verv larec differences amoung the rates. this approximation to equation (5)) is nearly exact.," Given these values,$1\ll A/(\alpha_3n_e)\ll A/(b\zeta_{\rm He})$, and equation \ref{eqinvmetafrac}) ) can be approximated by Owing to the very large differences among the rates, this approximation to equation \ref{eqinvmetafrac}) ) is nearly exact."
+ This holds tue as long as cy. υρίαι|03)10 Ps1 such that ionization of helium by cosuic-ravs is the rate-limiting step in the path to the metastable state.," This holds true as long as $\zeta_{\rm He}\ll n_e(\alpha_1+\alpha_3) \sim 10^{-13}$ $^{-1}$, such that ionization of helium by cosmic-rays is the rate-limiting step in the path to the metastable state."
+" Iu this Dit f, effectively depends on cy. alone = apart frou the well-determined atomic constants b and A tls sugecsting mctastable helium as a fairly robust indicator of the cosmic-ray ionization rate.", In this limit $f_m$ effectively depends on $\zeta_{\rm He}$ alone – apart from the well-determined atomic constants $b$ and $A$ – thus suggesting metastable helium as a fairly robust indicator of the cosmic-ray ionization rate.
+ The fundamental question remaining then is whether the interstellar lines of Πο arising from a suitable diffuse cloud are likely to be detectable., The fundamental question remaining then is whether the interstellar lines of He* arising from a suitable diffuse cloud are likely to be detectable.
+ The streugtlis of these lines are fixed by the cloud’s column deusity of nmuetastable atoms. ων which can be calculated frou where AN(Ile) is the total columm deusitv of hel atoms dn all states. AQ; is the total column deusitv of hydrogen unclei [Ny=ΑΠ and AGI)=Αποδι=0.097 is the relative ΔΟΥ).abuudauce of helimm with respect to hydrogen (Anders&Crevesse 1989)..," The strengths of these lines are fixed by the cloud's column density of metastable atoms, $N_m$, which can be calculated from where $N({\rm He})$ is the total column density of helium atoms in all states, $N_{\rm H}$ is the total column density of hydrogen nuclei $N_{\rm H} = N({\rm H}) + 2N({\rm H}_2)$ ], and $A({\rm He})=N({\rm He})/N_{\rm H}=0.097$ is the relative abundance of helium with respect to hydrogen \citep{and89}. ."
+ The fraction of interstellar helium sequestered in the graius has also been assumed negligible., The fraction of interstellar helium sequestered in the grains has also been assumed negligible.
+ Ifa direct 11easuremoeut of Nyy is not available. an alternative is to use Ay= V). where E(BV) is the observed color excess.," If a direct measurement of $N_{\rm H}$ is not available, an alternative is to use $N_{\rm H}=\beta E(B-V)$ , where $E(B-V)$ is the observed color excess,"
+of the V-band PL relation mapped to W would be reduced by a factor of —(—ο).,of the V-band $PL$ relation mapped to $W$ would be reduced by a factor of $-(R - \beta)/\beta$.
+ This is (he same factor that results in Che decreased scatter observed in WW as compared to the scalter at [ixed period in the V-band PL relation., This is the same factor that results in the decreased scatter observed in $W$ as compared to the scatter at fixed period in the V-band $PL$ relation.
+ In the extreme case where ;? equals R then W would not change its slope at all for anv change in tilt of the instabilitw strip., In the extreme case where $\beta$ equals $R$ then $W$ would not change its slope at all for any change in tilt of the instability strip.
+ The central portion of Figure 1 shows a V-(V-1) color-magnitude diagram with two possible realizations of the Cepheid instability strip cutting across lines of constant. period as laid down by the underlving period-huminositv-color relation., The central portion of Figure 1 shows a V-(V-I) color-magnitude diagram with two possible realizations of the Cepheid instability strip cutting across lines of constant period as laid down by the underlying period-luminosity-color relation.
+ The strip to the left ls significantly more vertical that the strip to the right., The strip to the left is significantly more vertical that the strip to the right.
+ Both instability strips share (he sime zero point in this realization. but thev could differ without changing the conclusions of this paper.," Both instability strips share the same zero point in this realization, but they could differ without changing the conclusions of this paper."
+ Fiducal points are marked in both strips bx open and filled circles., Fiducial points are marked in both strips by open and filled circles.
+ The filled circes in (he more highly inclined strip mark (he blue ancl red borders of the instability strip : two disünct periods., The filled circles in the more highly inclined strip mark the blue and red borders of the instability strip at two distinct periods.
+ The open circles in the left-hand strip mark the same color boundari al precisely the same (wo periods., The open circles in the left-hand strip mark the same color boundaries at precisely the same two periods.
+ The clownward-sloping dashed lines crossing the enti figure are lines of constant period., The downward-sloping dashed lines crossing the entire figure are lines of constant period.
+ The dashed horizontal lines carry the individual data points im the tilted instabililv strip over to the right and into its corresponding PL relation., The dashed horizontal lines carry the individual data points in the tilted instability strip over to the right and into its corresponding $PL$ relation.
+ Data from (he steep instabililv strip map to the left L relation., Data from the steep instability strip map to the left $PL$ relation.
+ Clearly by Uilling the instability strip one generates PL relations with different slopes., Clearly by tilting the instability strip one generates $PL$ relations with different slopes.
+ However. by projecüng the instability. strips along lines ol constant period. as is done bv following the broken lines to the upper left panel. one finds a dispersionless PL relation as defined by 1H; which. in addition to having no scatter. is incapable of distinguishing between stars originating either in (he vertical or the tilted instability strip.," However, by projecting the instability strips along lines of constant period, as is done by following the broken lines to the upper left panel, one finds a dispersionless $PL$ relation as defined by $W_{\beta}$ which, in addition to having no scatter, is incapable of distinguishing between stars originating either in the vertical or the tilted instability strip."
+ The solid lines emanating ουν the filled data points in the tilled instability strip show the trajectories of reddening lines., The solid lines emanating from the filled data points in the tilted instability strip show the trajectories of reddening lines.
+ Their sieht divergence [rom lines of constant period results in a coniposite W-logP relation that has some residual scatter aud a slight dependence of its slope on the originating slant of the instability strip., Their slight divergence from lines of constant period results in a composite W-logP relation that has some residual scatter and a slight dependence of its slope on the originating slant of the instability strip.
+ Unlike Che gross differences seen in the (left ancl right) monochromatic PL relations the divergence in the Wy relations can hardly be resolved and require an exploded view as given in Figure 2 (o see the impact., Unlike the gross differences seen in the (left and right) monochromatic PL relations the divergence in the $W_{\beta}$ relations can hardly be resolved and require an exploded view as given in Figure 2 to see the impact.
+ lt has been claimed (Sandage Tamimann 2008. their Table 1) that variations in the," It has been claimed (Sandage Tammann 2008, their Table 1) that variations in the"
+equation of state can be computed. with ellipticities as hich as 10 predicted: for some configurations (Llaskelletal. 2008).,"equation of state can be computed, with ellipticities as high as $10^{-5}$ predicted for some configurations \citep{haskellb08}."
+.. Ackelitionally. some exotic neutron star models (e.g. solid. strange quark stars) allow for cllipticities as large as 101 (Owen|2005).," Additionally, some exotic neutron star models (e.g. solid strange quark stars) allow for ellipticities as large as $10^{-4}$ \citep{owen05}."
+ Accreting neutron stars in binary svstenis can Dorm magnetic mountains. with c<10 (Melatos&Payne2005:Vigelius|Melatos20092) ).," Accreting neutron stars in binary systems can form magnetic mountains, with $\epsilon \leq 10^{-5}$ \citep{melpa05, vigmel09}."
+ Deformations of all kinds relax viseoclastically and resistively over time. so that voung neutron stars are expected to be generally stronger gravitational wave emitters.," Deformations of all kinds relax viscoelastically and resistively over time, so that young neutron stars are expected to be generally stronger gravitational wave emitters."
+ For example. thermoclastic deformations relax on the thermal conduction time-scale (~—103 vr). after the temperature gradient in the crust has switched olf (Brown&Bildsten1998:VigeliusAlelatos 2010).," For example, thermoelastic deformations relax on the thermal conduction time-scale $\sim 10^4$ yr), after the temperature gradient in the crust has switched off \citep{brown98, vigmel10}."
+". Magnetic mountains relax as the acereted matter dilfuses through the magnetic field on the Ohmic time-scale 10"". 107 vvr (Vigelius&Molatos 2009b).", Magnetic mountains relax as the accreted matter diffuses through the magnetic field on the Ohmic time-scale $10^5$ $10^8$ yr \citep{vigmel09b}.
+. A coherent search. for 7S known radio pulsars was »erformed. on 83 and 84 LIGO and GEO 600 data., A coherent search for 78 known radio pulsars was performed on S3 and S4 LIGO and GEO 600 data.
+" Upper imits on the ellipticities of these pulsars were obtained. the smallest beingex7.110"" for PSR J2124.3358 (Abbottetal. 2007¢)."," Upper limits on the ellipticities of these pulsars were obtained, the smallest being $\epsilon \leq 7.1 \times 10^{-7}$ for PSR J2124–3358 \citep{lscpulsars07}."
+. More recently. a coherent search for 116 known »ulsars was carried out using data from both the LIGO and Virgo detectors. placing an upper limit of«<7.0.10.7 for SR J2124. 3358.," More recently, a coherent search for 116 known pulsars was carried out using data from both the LIGO and Virgo detectors, placing an upper limit of $\epsilon < 7.0 \times 10^{-8}$ for PSR $-$ 3358."
+ The voungest isolated neutron star accessible to LIGO »obably resides in the supernova remnant SNR 1987. “Phe coincident. detection of neutrino bursts from the supernova w detectors all over the world. confirmed the core-collapse event. stronely indicating the formation of a neutron star (Agliettaetal.1987:LliratactBiontaLOST:Bahealletal.," The youngest isolated neutron star accessible to LIGO probably resides in the supernova remnant SNR 1987A. The coincident detection of neutrino bursts from the supernova by detectors all over the world confirmed the core-collapse event, strongly indicating the formation of a neutron star \citep{aglietta87, hirata87, bionta87, bahcall87}."
+LOST)..4 Constraints have been placed. on the magnetic field strength. spin period. and other birth properties of the »utative neutron star (Michel1994:OÓgelman&Alpar 2004): see Section 2. for details.," Constraints have been placed on the magnetic field strength, spin period, and other birth properties of the putative neutron star \citep{michel94, ogelman04}; see Section \ref{sn1987Adetails} for details."
+ However. searches for a »ulsar in SNR. LSTA have vielded no confirmed sightings: upper limits on its luminosity have been placed in the racio. optical and. N-ray bands. (Percivalctal.1995:Burrowsetal.2000:Manchester 2PO07).," However, searches for a pulsar in SNR 1987A have yielded no confirmed sightings; upper limits on its luminosity have been placed in the radio, optical and X-ray bands \citep{percival95, burrows00, manchester07}."
+. An unconfirmed. detection ofa transitory 467.5 Hz optical/infra-red. pulsation in SNR LOSTA was reported by Micddlecitehetal.(2000)., An unconfirmed detection of a transitory 467.5 Hz optical/infra-red pulsation in SNR 1987A was reported by \citet{middleditch00}.
+. The likely existence of a voung neutron star in SNR LOSTA makes it à good target for gravitational wave searches (Piran&Nakamura1988:1989).," The likely existence of a young neutron star in SNR 1987A makes it a good target for gravitational wave searches \citep{piran88, nakamura89}."
+. A coherent matched. filtering search was carried out in 2003 with the TAALA 300 detector. searching 1.2«105 hours of data from its lirst science run over a 1-Hz band centered on LILIz. assuming. a spin-down. range of 10n itz s+.," A coherent matched filtering search was carried out in 2003 with the TAMA 300 detector, searching $1.2 \times 10^3$ hours of data from its first science run over a 1-Hz band centered on Hz, assuming a spin-down range of $\times 10^{-10}$ Hz $^{-1}$."
+ JEPhe search vielded an upper limit on the wave strain of ο10E25 (Soidaetal.2003)., The search yielded an upper limit on the wave strain of $5 \times 10^{-23}$ \citep{soida03}.
+ An earlier matched filtering search was conducted: using 1t hours of data taken in 1989 by the Garching prototvpe laser interferometer., An earlier matched filtering search was conducted using $10^2$ hours of data taken in 1989 by the Garching prototype laser interferometer.
+ The latter search was carried. out over d-Hz bands near kklIIz ancl -ΕΚΚΙΙΣ. did not include any spin-down parameters. ancl vielded an upper limit of 91074 on the wave strain (Niebaueretal.1993).," The latter search was carried out over 4-Hz bands near kHz and kHz, did not include any spin-down parameters, and yielded an upper limit of $9 \times 10^{-21}$ on the wave strain \citep{niebauer93}."
+ There are two main tvpes of continuous-wave LIGO searches: coherent and. semi-coherent., There are two main types of continuous-wave LIGO searches: coherent and semi-coherent.
+ Ehe former demand phase coherence between the signal and search template over the entire time series., The former demand phase coherence between the signal and search template over the entire time series.
+ Although sensitive. they are restricted to small observation times ancl parameter ranges as they are computationally intensive.," Although sensitive, they are restricted to small observation times and parameter ranges as they are computationally intensive."
+ Semi-coherent searches break the full time series into many small chunks. analyse each chunk coherently. then sum the results incoherently. trading olf sensitivity for computational load.," Semi-coherent searches break the full time series into many small chunks, analyse each chunk coherently, then sum the results incoherently, trading off sensitivity for computational load."
+ Santostasictal.(2003) discussed the cletectability of gravitational waves from SNI OSTA. estimating that a coherent search. based. on. the Alicelleclitehetal.(2000) spin parameters requires 30 days of integration time and at least 107if search templates covering just the frequency ancl its first. derivative.," \citet{santostasi03} discussed the detectability of gravitational waves from SNR 1987A, estimating that a coherent search based on the \citet{middleditch00} spin parameters requires 30 days of integration time and at least $10^{19}$ search templates covering just the frequency and its first derivative."
+ Lo reality. the ask is even more daunting. because such a voung object spins clown so rapidly. that five or six higher-order frequency derivatives must be searched. in order to accurately track he gravitational wave phase.," In reality, the task is even more daunting, because such a young object spins down so rapidly, that five or six higher-order frequency derivatives must be searched in order to accurately track the gravitational wave phase."
+ A Bayesian Markov. Chain Alonte Carlo method was proposed as an alternative to Cover he parameter space efficientIv (Umstatteretal.2004:Um-stüetterctal. 2008).," A Bayesian Markov Chain Monte Carlo method was proposed as an alternative to cover the parameter space efficiently \citep{umstaetter04, umstaetter08}."
+.. As vet. though. SNR. 1987 has not »en considered a feasible search target. because even Monte Carlo methods are too arduous.," As yet, though, SNR 1987A has not been considered a feasible search target, because even Monte Carlo methods are too arduous."
+ In this paper. we show how o reduce the search space dramatically by assuming an astrophysically motivated phase model.," In this paper, we show how to reduce the search space dramatically by assuming an astrophysically motivated phase model."
+ In this paper. we discuss how to use a cross-correlation algorithm to search for periodic gravitational waves [rom a neutron star in SNR LOSTA. The search is semi-coherent (Dhurandharetal.2008).," In this paper, we discuss how to use a cross-correlation algorithm to search for periodic gravitational waves from a neutron star in SNR 1987A. The search is semi-coherent \citep{dhurandhar08}."
+. The signal-to-noise ratio is enhanced. by cross-correlating two cata sets. separated: by an adjustable time lag. or two simultaneous data sets from dillerent interferometers. thereby nullifving short-term timing noise (e.g. [rom rotational elitches).," The signal-to-noise ratio is enhanced by cross-correlating two data sets separated by an adjustable time lag, or two simultaneous data sets from different interferometers, thereby nullifying short-term timing noise (e.g. from rotational glitches)."
+ This is a modification of the method. used. in searches for a cosmological stochastic background. (Abbottctal.2009a) and for the low-mass X-ray binary Sco X-1 (Abbottetal. 2007b).," This is a modification of the method used in searches for a cosmological stochastic background \citep{lscstochastic07, lscstochastic09} and for the low-mass X-ray binary Sco X-1 \citep{lscscox107}."
+. In. Section 2.. we review the properties of SNR LSTA and its putative neutron star.," In Section \ref{sn1987Adetails}, we review the properties of SNR 1987A and its putative neutron star."
+ Section 32 briellv describes the cross-correlation algorithm and the data format., Section \ref{sec:crosscorr} briefly describes the cross-correlation algorithm and the data format.
+ We estimate the theoretical sensitivity of the search in Section 4.., We estimate the theoretical sensitivity of the search in Section \ref{sec:ccsensitivity}.
+ Section 5. describes an astrophysical model. which expresses the gravitational wave phase in terms of the initial spin. ellipticitv. magnetic field. ancl electromagnetic braking index of the neutron star.," Section \ref{searchparams} describes an astrophysical model, which expresses the gravitational wave phase in terms of the initial spin, ellipticity, magnetic field, and electromagnetic braking index of the neutron star."
+ We calculate the semi-coherent phase metric and the number of templates required [or the search. in the context of the astrophysical phase model., We calculate the semi-coherent phase metric and the number of templates required for the search in the context of the astrophysical phase model.
+ Given the computational resources available to us. we derive upper limits on the gravitational wave strain. cllipticity and maenctic field which can be placed. on a neutron star in SNR. 10στ with a cross-correlation search.," Given the computational resources available to us, we derive upper limits on the gravitational wave strain, ellipticity and magnetic field which can be placed on a neutron star in SNR 1987A with a cross-correlation search."
+ Finally. Section 7 summarises the results.," Finally, Section \ref{sec:ccconclusion} summarises the results."
+" SNA 12987, is the remnant of a Pype Lh core-collapse supernova which occurred in February LOST. kKkpe away"," SNR 1987A is the remnant of a Type II core-collapse supernova which occurred in February 1987, kpc away"
+As it is difficult to obtain a large sample of high quality. high resolution multi-band (e.g. both V. and H) data. we used the “structure mapy technique developed by Pogge to enhance the observable dust structure in each galaxy.,"As it is difficult to obtain a large sample of high quality, high resolution multi-band (e.g., both $V$ and $H$ ) data, we used the “structure map” technique developed by \citet{pogge02} to enhance the observable dust structure in each galaxy."
+" Structure maps can be thought of as ""single-band color maps: while requiring data in only one baud. they. provide a sharp contrast between dust aud star-forming regions."," Structure maps can be thought of as “single-band color maps;” while requiring data in only one band, they provide a sharp contrast between dust and star-forming regions."
+ Martini&Pogge(1999) note that (he regions of highest contrast in V.—1 color maps are primarily due to absorption by dust siructures in the V. image alone. while the smoother //-band images trace the underlving stellar population.," \citet{martini99b} note that the regions of highest contrast in $V-H$ color maps are primarily due to absorption by dust structures in the $V$ image alone, while the smoother $H$ -band images trace the underlying stellar population."
+ Therefore. the //-bancl images serve mostly to subtract larger-scale surface brightness variations Irom the V-band images in the V—// color maps.," Therefore, the $H$ -band images serve mostly to subtract larger-scale surface brightness variations from the $V$ -band images in the $V-H$ color maps."
+ Dy making structure maps from images taken al visible wavelengths. it is possible to enhance the structure due to dust and emission regions.," By making structure maps from images taken at visible wavelengths, it is possible to enhance the structure due to dust and emission regions."
+ Furthermore. as visible-wavelength images [rom HST have higher resolution (han those in the near infrared (both being diffraction limited). a structure map can make (he most of high quality data.," Furthermore, as visible-wavelength images from HST have higher resolution than those in the near infrared (both being diffraction limited), a structure map can make the most of high quality data."
+ The structure map is mathematical defined as where 5 is the structure map. J is the image. {2 is the point spread [nnction (PSF). 27! is (he transform of the PSF. and <2 is the convolution operator.," The structure map is mathematical defined as where $S$ is the structure map, $I$ is the image, $P$ is the point spread function (PSF), $P^t$ is the transform of the PSF, and $\otimes$ is the convolution operator."
+ Structure maps bring oul variations on the smallest resolvable scale of the image. i.e.. that of the PSF.," Structure maps bring out variations on the smallest resolvable scale of the image, i.e., that of the PSF."
+ For our objects. (his corresponds (to a projected scale of 115pe.," For our objects, this corresponds to a projected scale of 1–15pc."
+ The PSFs are modelled using the TinvTun software (Ixrist&Hook1999)., The PSFs are modelled using the TinyTim software \citep{krist99}.
+. Formally. structure maps are similar {ο the second-order iteration of the Richardson- (R-L) image correction (Richardson1972:Lucy1974).," Formally, structure maps are similar to the second-order iteration of the Richardson-Lucy (R-L) image correction \citep{richardson72, lucy74}."
+.. At this step in the image restoration. the first-order. smooth structure of the image has been removed. but the structures on the scale of the size of the PSF remain.," At this step in the image restoration, the first-order, smooth structure of the image has been removed, but the structures on the scale of the size of the PSF remain."
+ Structure maps for the 75 galaxies in our sample are given in Figure 1: in all of (he structure maps presented in (his paper. the dusty regions appear dark while enhanced stellar light and emission-line regions ave bright (e.g.. the F606W filter achnits several bright emission lines.)," Structure maps for the 75 galaxies in our sample are given in Figure \ref{fig:stmaps}; in all of the structure maps presented in this paper, the dusty regions appear dark while enhanced stellar light and emission-line regions are bright (e.g., the F606W filter admits several bright emission lines.)"
+Cosmic reionization represents one of the most. dramatic feedback processes occurring between Luminous sources and the intergalactic medium. (IGM).,Cosmic reionization represents one of the most dramatic feedback processes occurring between luminous sources and the intergalactic medium (IGM).
+ At some point before redshift six. nearly every hydrogen atom in the Universe was lonized. as evidenced by the transmission of pphotons in QSO spectra up to z~6 (??7)..," At some point before redshift six, nearly every hydrogen atom in the Universe was ionized, as evidenced by the transmission of photons in QSO spectra up to $z \sim 6$ \citep{becker2001, fan2002, fan2006b}."
+ Similarly. helium in the IG. is known to be fully ionized by 2~2.7 based. on observed transmission in the [forest at 2<3 (forasummaryofobservations.see.2)..," Similarly, helium in the IGM is known to be fully ionized by $z \sim 2.7$ based on observed transmission in the forest at $z < 3$ \citep[for a summary of observations, see][]{dixon2009}."
+ Understanding how and when these reionization events occurred is fundamental to understanding the nature of the earliest. galaxies and AGN. as well in assessing the impact of reionization on the evolution of luminous sources up to the present clay.," Understanding how and when these reionization events occurred is fundamental to understanding the nature of the earliest galaxies and AGN, as well in assessing the impact of reionization on the evolution of luminous sources up to the present day."
+ One of the main impacts of byelrogen and helium, One of the main impacts of hydrogen and helium
+hekl view is tlal gravitational interactions with an unseen embedded: planet are responsible for most Cases of cisk structure (Ozernoyetal.2000:Quillen&Thorudike2002:Wilner2002).,"held view is that gravitational interactions with an unseen embedded planet are responsible for most cases of disk structure \citep{oz00,qt02,whk02}."
+". Alternaive possible explauatious have been cousidered [or specili€ systems, such as clistwbance by a passlng star (e.g.HD111569:Clampinetal.2003).. a recewo large cometary/planetesimal collision (e.g.Fomalhaut:Wyatt&Dent2002) and coutaiminatioi by a background feature such as a distant gaoyaxy (e.g.Vega:Wilueretal.2002)."," Alternative possible explanations have been considered for specific systems, such as disturbance by a passing star \citep[e.g. HD141569:][]{clampin03}, a recent large cometary/planetesimal collision \citep[e.g. Fomalhaut:][]{wyatt02}
+ and contamination by a background feature such as a distant galaxy \citep[e.g. Vega:][]{whk02}."
+. A planetary companion has several efects ou an initially uuiorm debris disk., A planetary companion has several effects on an initially uniform debris disk.
+ Particles near the planet are often ejected due to a close encounter — an effect wich is enlaucecl as the jXanet's nass lk‘reases., Particles near the planet are often ejected due to a close encounter – an effect which is enhanced as the planet's mass increases.
+ Since particles interior to tle planet spiral increasinely rapidly iuto the cent‘al star. au iuterior cleared zoie is usually. present (Liou&Zook1999).," Since particles interior to the planet spiral increasingly rapidly into the central star, an interior cleared zone is usually present \citep{lz99}."
+. However. some particles spiralling in towards the star 1iiv become trapped i a series of Mean Motion Resonances (ÀINBsS) with tte platet.," However, some particles spiralling in towards the star may become trapped in a series of Mean Motion Resonances (MMRs) with the planet."
+ An MAIR exis«s when the ratio of orbital periods for two ojects orbiting a central mass S equal to nm:n. where m and i are integers.," An MMR exists when the ratio of orbital periods for two objects orbiting a central mass is equal to $m:n$, where $m$ and $n$ are integers."
+ The order of an MMR is given by the quantity [ii—n]., The order of an MMR is given by the quantity $\left| m-n \right|$.
+ Earh bas a dust rii& formed by cust particles in NMBs (Deriuotetal.1991).. aud Neptune is suspected to do tle sue to Ixuiper Belt dust (Liou&Zook1999).," Earth has a dust ring formed by dust particles in MMRs \citep{dermott94}, and Neptune is suspected to do the same to Kuiper Belt dust \citep{lz99}."
+. Whilst lockec ---wo au MMR with a planet. the augular monentunm of the particle lost to aaud solar wind «rag is balanced by the resonant forcing of the planets gravitation.," Whilst locked into an MMR with a planet, the angular momentum of the particle lost to and solar wind drag is balanced by the resonant forcing of the planet's gravitation."
+ A particle trappec in an MIR can exteud the particles lifetine ο many times the value predicted by aalkl solar wiud drag aloue., A particle trapped in an MMR can extend the particle's lifetime to many times the value predicted by and solar wind drag alone.
+ The resulaut ‘pile-up’ of »articles at cliscrete (tends O Cause p‘omninent radial anc angular structre 1 the debris disk ).., The resultant `pile-up' of particles at discrete tends to cause prominent radial and angular structure in the debris disk \citep{kh03}.
+ lu uultiple planet systems. it is tle ouermost planet that generally dominates the structure of the resttine dlebris disk.," In multiple planet systems, it is the outermost planet that generally dominates the structure of the resulting debris disk."
+ Neptune is suspececd to play tiis role iu interactions with Ixuiper Belt dust (Liou&Zooς1999)., Neptune is suspected to play this role in interactions with Kuiper Belt dust \citep{lz99}.
+. The interior plalets. however. can have significant effects ou j»articles which ‘leak past” tLe outermost planet.," The interior planets, however, can have significant effects on particles which `leak past' the outermost planet."
+ Iu tle nOar system. [or example. Jupiter aud Satrh eject many dust. particles which drift interior to Nepme (Liou&Zook1999).," In the solar system, for example, Jupiter and Saturn eject many dust particles which drift interior to Neptune \citep{lz99}."
+. For simplicity. we lnvestigate ouly single panet systems and assume the the effects of auy interior planets which jay be present is negligible outside the orbit of the outer:lost planet.," For simplicity, we investigate only single planet systems and assume that the effects of any interior planets which may be present is negligible outside the orbit of the outermost planet."
+ In this paper. we uumerically situlate a large uuuber of debrisdisk systel5. creating a syntetic debris disk catalogue.," In this paper, we numerically simulate a large number of debrisdisk systems, creating a synthetic debris disk catalogue."
+ Our utumerical methoc is clescribed in Section 2.. axl the testing of e code is detailed in Section 3).," Our numerical method is described in Section \ref{sec:numericmeth}, and the testing of the code is detailed in Section \ref{sec:verifycode}."
+ We present a sina| sample of the results [rom the svithetic Calaogue in Section L. aud use the catalogue to selec planetary configurations which could be res}iisible for the observed systems of Vega. aal Fomalhaut.," We present a small sample of the results from the synthetic catalogue in Section \ref{sec:synthcat}, and use the catalogue to select planetary configurations which could be responsible for the observed systems of Vega, and Fomalhaut."
+ The selected systems are then simulaed iu higher detail in Sectjon 5.. and the Jes=ts compared to observational data.," The selected systems are then simulated in higher detail in Section \ref{sec:compareprev}, , and the results compared to observational data."
+ Our couclusious are presented in Section 6.., Our conclusions are presented in Section \ref{sec:conclusions}. .
+Most of the large-scale. undirected pulsar survevs conducted since (he time of discovery of the millisecond pulsars (Backer et al.,"Most of the large-scale, undirected pulsar surveys conducted since the time of discovery of the millisecond pulsars (Backer et al."
+ 1982) have been driven by an astrophvsically well molivated desire to detect more of these rapidly rotating neutron stus (e.g. Phinney Kulkarni 1994: Lorimer 2001)., 1982) have been driven by an astrophysically well motivated desire to detect more of these rapidly rotating neutron stars (e.g. Phinney Kulkarni 1994; Lorimer 2001).
+ This is certainly true for the spectacularly successful southern hemisphere survevs with the Parkes Gti (telescope (Manchester et al., This is certainly true for the spectacularly successful southern hemisphere surveys with the Parkes 64-m telescope (Manchester et al.
+ 2001: Edwards Bailes 2001). as well as for a number of less extensive. but more sensitive Arecibo surveys that have contributed several milestone pulsar discoveries (Camilo 1995).," 2001; Edwards Bailes 2001), as well as for a number of less extensive, but more sensitive Arecibo surveys that have contributed several milestone pulsar discoveries (Camilo 1995)."
+ In fact. the progress in this field has become particularly dramatic after the first discoveries of millisecond pulsars awav from the ealactic plane (Wolszezan 1991). which indicated an approximately isotropic distribution of these objects in the solar neighborhood.," In fact, the progress in this field has become particularly dramatic after the first discoveries of millisecond pulsars away from the galactic plane (Wolszczan 1991), which indicated an approximately isotropic distribution of these objects in the solar neighborhood."
+ Along with the millisecond pulsars. the surveys ustally discover the objects of the slow pulsar variety in much greater numbers.," Along with the millisecond pulsars, the surveys usually discover the objects of the so-called slow pulsar variety in much greater numbers."
+ For example. the Arecibo drift-scan searches carried out since 1991 have found 9 millisecond pulsars and at least TO slow pulsars (Foster οἱ al.," For example, the Arecibo drift-scan searches carried out since 1991 have found 9 millisecond pulsars and at least 70 slow pulsars (Foster et al."
+ 1995: Camilo et al., 1995; Camilo et al.
+ 1996: Rav et al., 1996; Ray et al.
+ 1996: Nilouris et al., 1996; Xilouris et al.
+ 2000: MeLaughlin. Cordes Arzoumanian 2000: Lommen οἱ al.," 2000; McLaughlin, Cordes Arzoumanian 2000; Lommen et al."
+ 2000: Lorimer et al., 2000; Lorimer et al.
+ 2002: Chandler 2003)., 2002; Chandler 2003).
+ Newly discovered slow pulsars certainly deserve to be subjected to close scrutiny., Newly discovered slow pulsars certainly deserve to be subjected to close scrutiny.
+ For example. there exists a possibility that some slow pulsars have black hole or planetary companions (e.g. Thorsett Dewev 1993: Lipunov et al.," For example, there exists a possibility that some slow pulsars have black hole or planetary companions (e.g. Thorsett Dewey 1993; Lipunov et al."
+ 1994: Ixalogera et al., 1994; Kalogera et al.
+ 2001)., 2001).
+ Furthermore. timing nmeasurenients of pulsars. especially (he vounger ones. remain the only source of experimental evidence concerning the neutron star seismology (seeChengetal.1983).," Furthermore, timing measurements of pulsars, especially the younger ones, remain the only source of experimental evidence concerning the neutron star seismology \citep[see][]{cheng88}."
+. Finally. there is also an open issue of the pulsar emission mechanism (Melrose2000).. whose final clarification nmav require new. possibly unanticipated kinds of evidence.," Finally, there is also an open issue of the pulsar emission mechanism \citep{mel2000}, whose final clarification may require new, possibly unanticipated kinds of evidence."
+ The Penn State/NRL survey (Foster et al., The Penn State/NRL survey (Foster et al.
+ 1995) and the subsequent untargeted observations bv the Penn State group throughout the Arecibo upgrade period have effectively covered about 850 square degrees of the skv in a manner dictated bv (he logistics of the telescope, 1995) and the subsequent untargeted observations by the Penn State group throughout the Arecibo upgrade period have effectively covered about 850 square degrees of the sky in a manner dictated by the logistics of the telescope
+"Zhekov&Palla(2007) showed that the high-resolutionChandra spectrum can be fitted by a shock model where the dominant temperature is kkeV wwas actually the star with the coolest dominant plasma), though some contribution from plasma with kkeV was needed to achieve a good fit.","\citet{zhe07} showed that the high-resolution spectrum can be fitted by a shock model where the dominant temperature is keV was actually the star with the coolest dominant plasma), though some contribution from plasma with keV was needed to achieve a good fit."
+ This modelling clearly shows that the plasma in lis rather cool., This modelling clearly shows that the plasma in is rather cool.
+" We decided to use a simpler formalism, which does not try to reproduce shocks, but simply considers the addition of optically-thin thermal plasma (without any assumption on their origins)."," We decided to use a simpler formalism, which does not try to reproduce shocks, but simply considers the addition of optically-thin thermal plasma (without any assumption on their origins)."
+" In Xspec, the models using a distribution of thermal plasma (e.g. cópvmkI) fail to provide a fit close to the data: we therefore had to fit the data using a sum of individual optically-thin plasmas."," In Xspec, the models using a distribution of thermal plasma (e.g. $c6pvmkl$ ) fail to provide a fit close to the data: we therefore had to fit the data using a sum of individual optically-thin plasmas."
+" As could be expected, one, two or three temperature fits do not provide fits close to the data, and similar conclusions are reached for four-temperature fits with a single absorption (which is unsurprising in view of the c6pvmkl result)."," As could be expected, one, two or three temperature fits do not provide fits close to the data, and similar conclusions are reached for four-temperature fits with a single absorption (which is unsurprising in view of the $c6pvmkl$ result)."
+" Our model of choice is thus a four temperature model with individual absorptions: tbabsx3vphabsx vapec, where the first component represents the interstellar absorption, fixed to 8.9x10? ccm? (Diplas&Sav-age, 1994),, and the abundances of the absorption and emission components are assumed to be equal."," Our model of choice is thus a four temperature model with individual absorptions: $tbabs\times \sum vphabs\times vapec$ , where the first component represents the interstellar absorption, fixed to $\times 10^{19}$ $^{-2}$ \citep{dip94}, and the abundances of the absorption and emission components are assumed to be equal."
+" Adding a fifth thermal components (e.g. at kkeV) does not significantly improve the fit, and we thus sticked to the decision of using 4 components."," Adding a fifth thermal components (e.g. at keV) does not significantly improve the fit, and we thus sticked to the decision of using 4 components."
+" After the best-fit temperatures and absorptions were found, we investigated the impact of using non-solar abundances."," After the best-fit temperatures and absorptions were found, we investigated the impact of using non-solar abundances."
+" We began by keeping the helium abundance to 2 times solar (Repolustetal,2004), and let the nitrogen and oxygen abundances vary."," We began by keeping the helium abundance to 2 times solar \citep{rep04}, and let the nitrogen and oxygen abundances vary."
+ We perfom such a simultaneous fit to all available EPIC spectra of each revolution., We perfom such a simultaneous fit to all available EPIC spectra of each revolution.
+ A few conclusions could be drawn from this trial., A few conclusions could be drawn from this trial.
+" First, the pn spectra obtained in Large Window mode with the Medium filter (Revs 0156 and 0552, and second pn observation of Revs 535, 538 and 542) are clearly deviant from other data, showing the impact of pile-up."," First, the pn spectra obtained in Large Window mode with the Medium filter (Revs 0156 and 0552, and second pn observation of Revs 535, 538 and 542) are clearly deviant from other data, showing the impact of pile-up."
+" Since all other data (pn or MOS, LW or SW modes with Thick filter) overall agree with one another, we only discarded the pn data taken with Medium filter from further analysis."," Since all other data (pn or MOS, LW or SW modes with Thick filter) overall agree with one another, we only discarded the pn data taken with Medium filter from further analysis."
+" Second, the absorptions and temperatures of the fits do not vary much: we therefore fixed them to KT; of 0.09, 0.27, 0.56 and kkeV and Ny; of 0.1, 0.1, 0.71, 0."," Second, the absorptions and temperatures of the fits do not vary much: we therefore fixed them to $kT_i$ of 0.09, 0.27, 0.56 and keV and $N_{{\rm H},i}$ of 0.1, 0.1, 0.71, 0."
+" x10? ccm"".", $\times 10^{22}$ $^{-2}$.
+" Finally, the origins of the high x? can be better pinpointed."," Finally, the origins of the high $\chi^2$ can be better pinpointed."
+" On the one hand, MOS and pn data do not always agree, especially at 0.4kkeV (see Fig. 1))."," On the one hand, MOS and pn data do not always agree, especially at keV (see Fig. \ref{epicspec}) )."
+" This explains why some fits deviate from the mean behaviour: Revs 0731 and 1071 only provided pn data, while Revs 0156 and 1164 consist only of MOS data."," This explains why some fits deviate from the mean behaviour: Revs 0731 and 1071 only provided pn data, while Revs 0156 and 1164 consist only of MOS data."
+" This difference between MOS and pn is probably due to remaining calibration problems, not pile-up since iis far from the pile-up limit for data taken in SW mode with the Thick filter."," This difference between MOS and pn is probably due to remaining cross-calibration problems, not pile-up since is far from the pile-up limit for data taken in SW mode with the Thick filter."
+" On the other hand, all three instruments (MOS1, MOS2, and pn) sometimes display óy? of the same sign."," On the other hand, all three instruments (MOS1, MOS2, and pn) sometimes display $\delta \chi^2$ of the same sign."
+" This is often the case near strong lines, which points toward two explanations: (1) atomic parameters are imperfect (even for APEC - Fig."," This is often the case near strong lines, which points toward two explanations: (1) atomic parameters are imperfect (even for APEC - Fig."
+" 1 shows two examples: at kkeV, there is flux in the data but not the model - a line is probably missing; at 1.5keV, there is a line in the model but not the data) and (2) the asymetric line shape of iinfluence even EPIC data."," \ref{epicspec} shows two examples: at keV, there is flux in the data but not the model - a line is probably missing; at 1.5keV, there is a line in the model but not the data) and (2) the asymetric line shape of influence even EPIC data."
+" We cannot do much for the former, and the latter will actually be studied in detail in the third paper of this series - it is thus beyond the scope of this contribution."," We cannot do much for the former, and the latter will actually be studied in detail in the third paper of this series - it is thus beyond the scope of this contribution."
+" The next step is to free more abundances, but this could yield erratic and/or problematic and/or not better results."," The next step is to free more abundances, but this could yield erratic and/or problematic and/or not better results."
+" For example, the carbon abundance clearly gets unrealistic (whether freed last or first)."," For example, the carbon abundance clearly gets unrealistic (whether freed last or first)."
+" Indeed, low-resolution, broad-band"," Indeed, low-resolution, broad-band"
+In this subsection we show a basic logarithmic Sobolev inequality.,In this subsection we show a basic logarithmic Sobolev inequality.
+ la particular. we show the following lemma.," In particular, we show the following lemma."
+ First. let us mention that the ideas of the proof of this lemma are inspired from the proof of Ogawa [?.Corollary2.{]..," First, let us mention that the ideas of the proof of this lemma are inspired from the proof of Ogawa \cite[Corollary
+2.4]{Ogawa}."
+ The proof is divided into three steps. and the constants in the proof may vary [roin line to (Estimate of IT. Let 5>0. aud NEN be two arbitraryvariables.," The proof is divided into three steps, and the constants in the proof may vary from line to (Estimate of $\|u\|_{\wt{F}^{0}_{\infty,1}}$ Let $\g>0$, and $N\in \N$ be two arbitraryvariables."
+" We compute: where C>Q is a positive (Optimization in ;V). We optimize the previous inequality in No by setting: In this case we cau easily check that: In the case where ||js>οp, . we choose 153<2° such that"," We compute: where $C_{\g}>0$ is a positive (Optimization in $N$ We optimize the previous inequality in $N$ by setting: In this case we can easily check that: In the case where $\|u\|_{F^{\g}_{\infty,2}} > 2^{\g}
+\|u\|_{\wt{F}^{0}_{\infty,2}}$ , we choose $1\leq \b < 2^{\g}$ such that"
+spherical curvature in the cap while retaining the strong convergence to the pole.,spherical curvature in the cap while retaining the strong convergence to the pole.
+ Then the outer boundary of the cylinder is identified with low latitude boundary of the spherical polar cap., Then the outer boundary of the cylinder is identified with low latitude boundary of the spherical polar cap.
+" We recognize that imposing such a boundary is somewhat artificial, since no such physical boundary exists in the Sun."," We recognize that imposing such a boundary is somewhat artificial, since no such physical boundary exists in the Sun."
+" But we can think of it as the location where the physics that determines the meridional flow changes from the forces that are dominant in low and mid-latitudes, to those that are most important at high latitudes."," But we can think of it as the location where the physics that determines the meridional flow changes from the forces that are dominant in low and mid-latitudes, to those that are most important at high latitudes."
+" We do not know exactly the latitude where the dominant forces change, so we treat the location of the boundary as a parameter of the problem."," We do not know exactly the latitude where the dominant forces change, so we treat the location of the boundary as a parameter of the problem."
+ One guide to its placement comes from numerical simulations of global convection in deep rotating spherical shells (Gilman1979;Mieschetal2008) ," One guide to its placement comes from numerical simulations of global convection in deep rotating spherical shells \citep{g79,mbdt08}
+ "
+on sinulatio1 resolution. wor as accurate al uuderstandiug of the gas state such as metalicity. lemperatu'e. velocity auc iolization equilibrunm. as X-rays and Lymau-o properties do.,"on simulation resolution, nor as accurate an understanding of the gas state such as metalicity, temperature, velocity and ionization equilibrium, as X-rays and $\alpha$ properties do."
+ Pressure is the total hermal energy coleut of the gas. and finite volume flux conservative simulations are partictla ratmenable to 1 uodelliug.," Pressure is the total thermal energy content of the gas, and finite volume flux conservative simulations are particular amenable to its modelling."
+ It s also easier to moclel analytically. (, It is also easier to model analytically. (
+3) It las strong observation: potential.,3) It has strong observational potential.
+ In t yaper we will primarily cousider the thermal SZ ellect. which is easier {ο 0erve.," In this paper we will primarily consider the thermal SZ effect, which is easier to observe."
+ [ts uniqi »eudeuce ou requency allows us in princije to disentangle the SZ effect frou ile contamiuiati he primary CAB and other noise sources 2000).., Its unique dependence on frequency allows us in principle to disentangle the SZ effect from the contamination of the primary CMB and other noise sources \citep{Cooray00b}.
+ He'eafter. if not olerwise specified. ie. SZ ellect always meas tje thermal SZ effect.," Hereafter, if not otherwise specified, the SZ effect always means the thermal SZ effect."
+ Iu SZ observations. al redshifts are entaneled aud sinear some key info‘nation of the intervening ICM: dts spatia distribution aud time evotion.," In SZ observations, all redshifts are entangled and smear some key information of the intervening IGM: its spatial distribution and time evolution."
+ These properties around redshift z1 are sensitive to maly cosmological parameters :ul tlus potentially πεig to break degeneracies iu cosmologies f.‘om the primary CAB experilelts., These properties around redshift $z\sim 1$ are sensitive to many cosmological parameters and thus potentially promising to break degeneracies in cosmologies from the primary CMB experiments.
+ Oue cau Lope to resolve the eqration of the state of the dark energy., One can hope to resolve the equation of the state of the dark energy.
+ We are currently peToriing simulations for dillerent cosmologies. which are degeuerate 11i the primary CMB. to test le pxtential of such a procedure (Zhaugetal.2002j.," We are currently performing simulations for different cosmologies, which are degenerate in the primary CMB, to test the potential of such a procedure \citep{Zhang02}."
+ If comπο wit1 other observations. the SZ ellec will become even more powerlul.," If combined with other observations, the SZ effect will become even more powerful."
+ For example. with tje aid of a galaxy photometric redshif survey. the evolution of the 3D gas pressure power speciruu aud its cross correlation with the galaxy. distribution can be extracted. )..," For example, with the aid of a galaxy photometric redshift survey, the evolution of the 3D gas pressure power spectrum and its cross correlation with the galaxy distribution can be extracted \citep{Zhang01a}."
+ Tie. above analysis depends ou a detailed quantitative theoretical understancdii& of tle SZ statistics., The above analysis depends on a detailed quantitative theoretical understanding of the SZ statistics.
+ Analytical approaches considered in the past stronely depend on various assuiplous., Analytical approaches considered in the past strongly depend on various assumptions.
+ The P'ess-Schechter approac1 as adopted by Cole&Ixaiser(2000) requires ad hoc mocles fo “the eas profile iu halos. whose shape aud evolution a‘e urcertalu.," The Press-Schechter approach as adopted by \citet{Cole88,Makino93,Atrio99,Komatsu99,Cooray00,Molnar00} requires ad hoc models for the gas profile in halos, whose shape and evolution are uncertain."
+ The hierarchical method as oxoposed by Zhaug&Pen(2001) strougly depends ou he gas-dark matter correlation. which is :ilso »oorly understood.," The hierarchical method as proposed by \citet{Zhang01a} strongly depends on the gas-dark matter correlation, which is also poorly understood."
+ These estimates can be significanly inuproved by high-resolution simulatior s., These estimates can be significantly improved by high-resolution simulations.
+But past simulation results (Scaramellaetal.1993:daSilva9000:Relregieretal.2000:SeljasBurwell&Pen2001:Springel2001) disagreed. on bo aiplitude and shape of the SZ power spectrum (see Springeletal.(2001) for a recent. cliscussior," But past simulation results \citep{Scaramella93,daSilva00,Refregier00, Seljak01,Springel01}
+ disagreed on both amplitude and shape of the SZ power spectrum (see \citet{Springel01} for a recent discussion)."
+ Sunulatiou resolution may play a key ‘ole in this discrepancy since the SZ effect is sensitive to sinal structures. as discussed by Seljak.Burwell&Peu(2001).," Simulation resolution may play a key role in this discrepancy since the SZ effect is sensitive to small structures, as discussed by \citet{Seljak01}."
+.. Dillerences between code algoritluls nay also cause part of the ciscrepatrcy., Differences between code algorithms may also cause part of the discrepancy.
+ To address these problems. we rau the largest. adaptive SZ simulation to date. a direct 512° 1loving 1jesli hydro (MMH) code (Pen1998a) simulation. reinvestigate the SZ statistics.," To address these problems, we ran the largest adaptive SZ simulation to date, a direct $512^3$ moving mesh hydro (MMH) code \citep{Pen98a} simulation, to reinvestigate the SZ statistics."
+ For a yetter iuxcerstaudiug of numerical and code issues in the SZ effect. our group Is carrying out a series of siljulations with identical cosiiological parameters aia initial coudition but using different codes wih different resolutions 2002).," For a better understanding of numerical and code issues in the SZ effect, our group is carrying out a series of simulations with identical cosmological parameters and initial condition but using different codes with different resolutions \citep{Codes}."
+ Apart from simulation issues. a problem hat has received little attenion is analysis strategies for SZ data.," Apart from simulation issues, a problem that has received little attention is analysis strategies for SZ data."
+ Several interferometric arrays stch as AMIBA and SZA are under construction. but detailed data analysis models are still iu their infaucy.," Several interferometric arrays such as AMIBA and SZA are under construction, but detailed data analysis models are still in their infancy."
+ With our sinulated SZ taps. we cau," With our simulated SZ maps, we can"
+the mean-field theory. predicts the magnetic fiekl generation effects due to the interaction ol the Coriolis force (QxJ-effect) and the differential rotation (WVxJ-effect. where WV. is the large-scale velocity shear) with large-scale electric current (see. Radler.1969:Krause&Raicller.1980:Rogachevskii&IxXleeorin. 2003)).,"the mean-field theory predicts the magnetic field generation effects due to the interaction of the Coriolis force $\Omega\times J$ -effect) and the differential rotation $W\times J$ -effect, where $W$ is the large-scale velocity shear) with large-scale electric current (see, \citealp{rad69,krarad80,kle-rog:04a}) )."
+ The role of these mechanisms for the solar dvnamo is not well understood though it was shown (Pipin&Seehlafer.2008.2009)) that il is possible to construct the solar-tvpe dvnamo models including these effects.," The role of these mechanisms for the solar dynamo is not well understood though it was shown \citealp{2008arXiv0812.1466P,2009A&A...493..819P}) ) that it is possible to construct the solar-type dynamo models including these effects."
+ Our model includes the a- amd OxJ effects., Our model includes the $\alpha$ - and $\Omega\times J$ effects.
+" We introduce free parameters C, and CU io control their strength.", We introduce free parameters $C_{\alpha}$ and $C_{\delta}^{(\Omega)}$ to control their strength.
+" We use the solarconvection zone model computed by Stix(2002).. in which mixing length 6=oapA1. where A""=Vlogp _.is the inverse⋅ pressure scale height.⋅ and Aye=2."," We use the solarconvection zone model computed by \citet{stix:02}, in which mixing length $\ell=\alpha_{MLT}\left|\Lambda^{(p)}\right|^{-1}$, where $\mathbf{\boldsymbol{\Lambda}}^{(p)}=\boldsymbol{\nabla}\log\overline{p}$ is the inverse pressure scale height, and $\alpha_{MLT}=2$."
+ We confine the integration domain between 0.712. and 0.97222. in radius., We confine the integration domain between $R_{\odot}$ and $R_{\odot}$ in radius.
+ It extends [rom the pole to pole in latitude., It extends from the pole to pole in latitude.
+ The differential rotation profile. O=OyfoCr.pr). pcΠιν = cos@is a modified version of the analviical approximation to helioseismology data proposed by Antiaetal.(1998).. see Figure La.," The differential rotation profile, $\Omega=\Omega_{0}f_{\Omega}\left(x,\mu\right)$, $x=r/R_{\odot}$, $\mu=\cos\theta$ is a modified version of the analytical approximation to helioseismology data proposed by \citet{1998MNRAS.298..543A}, see Figure 1a."
+ The meridional flow is modeled in the form of two stationary circulation cells. one in the northern ancl one in the southern hemisphere. with a poleward motion in the upper and equator-ward motion in the lower part of the convection zone.," The meridional flow is modeled in the form of two stationary circulation cells, one in the northern and one in the southern hemisphere, with a poleward motion in the upper and equator-ward motion in the lower part of the convection zone."
+ Following Ixitchatinov (2011).. the meridional circulation velocity components were approximated via the orthogonal Chebyshev polvnomial decompositions: where," Following \citet{2011MNRAS.411.1059K}, , the meridional circulation velocity components were approximated via the orthogonal Chebyshev polynomial decompositions: where"
+eextends [rom the horigor Rey to the light cylinder located at Hye=¢/Qy.,extends from the horizon $R_{\rm BH}$ to the light cylinder located at $R_{LC}= c/\Omega_H$.
+ General relativistic effects of time clilation would effectively freeze out the reconnection process near the horizou., General relativistic effects of time dilation would effectively freeze out the reconnection process near the horizon.
+ Let us next estimate the recomection rate near the light cylinder «X the resulting hole., Let us next estimate the reconnection rate near the light cylinder of the resulting .
+. Taking ὃ as a thickuess of the resistive laver of total leueth £L~Rye. we can estimate the Luudquist (nagnetic Reynolds) nuujer Within. the Sweet-Parker (Ixusrud2005) model of reconnection. S1=/ο...(Ic).," Taking $\delta$ as a thickness of the resistive layer of total length $L \sim R _{LC}$, we can estimate the Lundquist (magnetic Reynolds) number within, the Sweet-Parker \citep{Kulsrud} model of reconnection, $S=(R_{LC}/\delta)^2$."
+ In a higt.lyv inaguetized.magnetized. planaplasma the Sweet-Parser recounectiou tuiflow velocity is tin~ofVS., In a highly magnetized plasma the Sweet-Parker reconnection inflow velocity is $v_{\rm in} \sim c / \sqrt{S}$.
+ The expected deusity cali be scaled to the local Colcdreich-.Jian density 1960).. = Anc;j.," The expected density can be scaled to the local Goldreich-Julian density \citep{GoldreichJulian}, $n = \lambda n_{GJ}$ ."
+ Taking the angular velocity of field lines equal to Qy. then ney=BOY/(2nec).," Taking the angular velocity of field lines equal to $\Omega_H$, then $n_{GJ} = B
+\Omega_H/(2 \pi e c)$."
+ The multiplicity A is expected tobe high. A<>L (RuclermanaidSutherland1975)..," The multiplicity $\lambda$ is expected tobe high, $\lambda \gg 1$ \citep{1975ApJ...196...51R}."
+" We then find ⋅ ,aay? ↕∐↩∟⋯∐⇂≺↽↓⋯⊳∖↕⋯∐∐∣⋈↵↓⋅⊳↔∶∐⋗⋋−⋜↕∐≺⇂↕⋅≺↵∢∙∩∐∐≺↵∢∙⋃∩∐↕∐∐≺↵⊳∖∢∙⋜↕↥≺↵⊤↓−≺⋅⋅⊲∶↓⋋∫∎≱⊓≬∖∖↽⇂∐∢∙∐∢∙⋜↕∐∣⋈↵∖∖↽⊔⊓≺↵∐⋅ ⋅ / ⋅ ⋅"," We then findthe Lundquist number $S= 16 \lambda^2 $ andreconnection time scale $\tau_{\rm rec} = 4 \lambda R_{LC}/c$ , which can be written"
+ollowiug recombination is typically supersouic.,following recombination is typically supersonic.
+ This relative velocity arises because dark matter is accelerated by eravitational potential eradieuts. while barvons are Jeans stabilized against gravitational collapse duc to their tight coupling with the photon radiation field. until recombination.," This relative velocity arises because dark matter is accelerated by gravitational potential gradients, while baryons are Jeans stabilized against gravitational collapse due to their tight coupling with the photon radiation field, until recombination."
+ Because DM ancl barvous suffer different acceleratious. hey acquire significant relative velocities that are predominantly sourced by potential fluctuatious on scales of order he sound horizon at recombination. ~150 AIpe.," Because DM and baryons suffer different accelerations, they acquire significant relative velocities that are predominantly sourced by potential fluctuations on scales of order the sound horizon at recombination, $\sim 150$ Mpc."
+ At recombination. the barvon sound speed and Jeans leugth fall xecipitouxlv. allowing barvous to respond to the same gravitational potential wells that accelerate dark matter.," At recombination, the baryon sound speed and Jeans length fall precipitously, allowing baryons to respond to the same gravitational potential wells that accelerate dark matter."
+" Subsequently, the lavge-scale relative velocity is unsomceed aud decays as «+ following recombination."," Subsequently, the large-scale relative velocity is unsourced and decays as $a^{-1}$ following recombination."
+ At redshifts 21000. the relative DM-baryon motion is highly supersonic. with Mach umubers Mo5.," At redshifts $z\approx
+1000$, the relative DM-baryon motion is highly supersonic, with Mach numbers ${\cal M} \approx 5$."
+ The gas sound speed does rot initially decay as quickly asa 1 because of residual thermal coupling to the CMB. so the Mach wmuber diminishes over time to ντο2 at 2~100. and remains nearly constant thereafter.," The gas sound speed does not initially decay as quickly as $a^{-1}$ because of residual thermal coupling to the CMB, so the Mach number diminishes over time to ${\cal M}\sim 2$ at $z\sim 100$, and remains nearly constant thereafter."
+ Because the relative motions between barvous and DAL are supersonic. they. have significaut effects on the erowth of structure.," Because the relative motions between baryons and DM are supersonic, they have significant effects on the growth of structure."
+ TII computed the effects of relative DALbarvon velocities at Ligh redshift. iu the perturbative regime of structure formation.," TH computed the effects of relative DM-baryon velocities at high redshift, in the perturbative regime of structure formation."
+ They showed that these motious cause a ~1056 suppression of the matter power spectruii at kom200Mpe+ compared to standard linear perturbation theory calculations., They showed that these motions cause a $\sim 10\%$ suppression of the matter power spectrum at $k\approx 200 {\rm Mpc}^{-1}$ compared to standard linear perturbation theory calculations.
+ Using a Press-Schechter approach. TII also sugeested that the abuudauce of dark matter lialos of mass AL~LOCAL. could be suppressed by a factor of ~2 at redshift :=10.," Using a Press-Schechter approach, TH also suggested that the abundance of dark matter halos of mass $M\sim 10^6 M_\odot$ could be suppressed by a factor of $\sim 2$ at redshift $z=40$."
+ At later redshifts closer to reionization. :z1020. the halo abundance would be ach closer to standard predictions.," At later redshifts closer to reionization, $z\approx 10-20$, the halo abundance would be much closer to standard predictions."
+ Iu this paper. we consider a similar effect of the supersouic relative velocities that TII discussed.," In this paper, we consider a similar effect of the supersonic relative velocities that TH discussed."
+ Instead of cousideriug the effect ou dark matter halos. we examine the imipact on barvonic objects; which Likely determine the properties of observable quantities like 21 cm absorption. ciuuission. etc.," Instead of considering the effect on dark matter halos, we examine the impact on baryonic objects, which likely determine the properties of observable quantities like 21 cm absorption, emission, etc."
+ The supersonic flow changes the mass threshold of biurvouicallv populated halos. aud this effect cau be exponentially huge for rare objects.," The supersonic flow changes the mass threshold of baryonically populated halos, and this effect can be exponentially large for rare objects."
+ The relative bulk flows between DAL and baryous are modulated ou large scales. of order LOO Προς meaning that the 1inibialo abundance is stnuilarly modulated.," The relative bulk flows between DM and baryons are modulated on large scales, of order $\sim 100$ Mpc, meaning that the minihalo abundance is similarly modulated."
+ This provides a large-scale signature of the effect of mimihalos., This provides a large-scale signature of the effect of minihalos.
+" All calculations preseuted lere asiunue WALAP (Larsonetal.2010) cosinological piuuneters: O,,=0.27. O4=0.73. QO,=0.015. Fh= 0.2. n4=0.96. σς=0.5."," All calculations presented here assume WMAP \citep{WMAP7} cosmological parameters: $\Omega_m=0.27$, $\Omega_\Lambda=0.73$, $\Omega_b=0.045$, $h=0.7$ , $n_s=0.96$, $\sigma_8=0.8$."
+" Iu this section. we compute the statistics of the collapsed barvouic density. which we parametz0üze using f,.. the collapsed fraction of barvous."," In this section, we compute the statistics of the collapsed baryonic density, which we parametrize using $f_c$, the collapsed fraction of baryons."
+ The barvonic collapsed fraction is different than the CDAL collapsed fraction because barvous. uulike CDM. have a nonzero temperature. and therefore cannot fall iuto shallow poteutial wells.," The baryonic collapsed fraction is different than the CDM collapsed fraction because baryons, unlike CDM, have a nonzero temperature, and therefore cannot fall into shallow potential wells."
+" Naively. we mieht expect that the characteristic mass AZ, unable to attract barvous would scale like the Jeans mass Mjx3GUL. however cosmological bydrodvuamic simulations (Caedin2000:Naozetal.2010) have found instead that AZ. scales like the filter scale (Cnedin&Thi1998).. For a standard WALAP cosmology. AL.~9ὃν104AZ, at 2~20 (Naozetal.20L0}."," Naively, we might expect that the characteristic mass $M_c$ unable to attract baryons would scale like the Jeans mass $M_J \propto c_s^3/GH$, however cosmological hydrodynamic simulations \citep{Gnedin00,Naoz10} have found instead that $M_c$ scales like the filter scale \citep{Gnedin98}, For a standard WMAP cosmology, $M_c\sim 2-3\times 10^4 M_\odot$ at $z\sim 20$ \citep{Naoz10}."
+. Just as gas pressure impedes the collapse of barvous. we expect that bulk relative velocity between gas and dark matter halos will also suppress the accretion of barvous.," Just as gas pressure impedes the collapse of baryons, we expect that bulk relative velocity between gas and dark matter halos will also suppress the accretion of baryons."
+ We can make a simple estimate of the minimum halo mass that can acerete eas moving at some bulk velocity ο relative to the halo through analogy to the above argument., We can make a simple estimate of the minimum halo mass that can accrete gas moving at some bulk velocity $v$ relative to the halo through analogy to the above argument.
+" Assunmiug that bulk kinetic energy is converted iuto thermal euerev when eas falls into the halo. we expect that ((1)) will still hold if we replace the sound speed e; with eig=(ος|02382,"," Assuming that bulk kinetic energy is converted into thermal energy when gas falls into the halo, we expect that \ref{Mc}) ) will still hold if we replace the sound speed $c_s$ with $c_{s,{\rm eff}}=(c_s^2+v^2)^{1/2}$."
+" Since e and e both scale as a+ (see below). this effectively umltiplies the eritieal mass scale by the factor (L|07ο3,"," Since $c_s$ and $v$ both scale as $a^{-1}$ (see below), this effectively multiplies the critical mass scale by the factor $(1+v^2/c_s^2)^{3/2}$."
+ Given AL... we can compute the barvon collapsed fraction f. bv πορταιο the halo mass function. where we use the fitting function of Shethetal.(2001). to describe the halo mass function d»/dAM.," Given $M_c$, we can compute the baryon collapsed fraction $f_c$ by integrating the halo mass function, where we use the fitting function of \citet{ShethTormen} to describe the halo mass function $dn/dM$."
+" Since AL. is a function of the local relative velocity between CDM and baryons. e,5—v.-—vy. we see that the collapse fraction f, is also modulated by this velocity. fia)=οία)]."," Since $M_c$ is a function of the local relative velocity between CDM and baryons, $\v_{cb}\equiv\v_c-\v_b$, we see that the collapse fraction $f_c$ is also modulated by this velocity, $f_c(\bm{x})=f_c(v_{cb}(\bm{x}))$."
+ We have assumed that the velocity dependence of the collapse fraction is a sharp cut-off in the mass function. whereas bydrodvuamic simulations fud a smooth transition in barvou content of halos below and above AL.," We have assumed that the velocity dependence of the collapse fraction is a sharp cut-off in the mass function, whereas hydrodynamic simulations \citep[e.g.][]{Gnedin00,Naoz10} find a smooth transition in baryon content of halos below and above $M_c$."
+ This distinction will not be siguificaut for our results: the important poiut is that the collapse fraction ix now a function of the local relative bulk velocity between dark matter aud barvous., This distinction will not be significant for our results: the important point is that the collapse fraction is now a function of the local relative bulk velocity between dark matter and baryons.
+ The relative velocity between CDM andbarvous can be computed using the linearized continuity equation using comoving coordinates rather than proper coordinates., The relative velocity between CDM andbaryons can be computed using the linearized continuity equation using comoving coordinates rather than proper coordinates.
+ In the lincar regiae. where we have potential flow. the velocity is," In the linear regime, where we have potential flow, the velocity is"
+or an expression of couservation of moientuum of the magnetized medium.,or an expression of conservation of momentum of the magnetized medium.
+" Ou the right-hiaud side olf Equation (11)). we have used 7,=1rg upstream."," On the right-hand side of Equation \ref{up_down}) ), we have used $r_n=1=r_B$ upstream."
+" Equatious (12)) aud (13)) represent the ""slrong coupling"" approximation. in which the full maguetic force on the ious is conveyed to the neutrals. Le. (Shu1992.Equation (27.3)).. "," Equations \ref{drB2dx}) ) and \ref{drn+dB2}) ) represent the “strong coupling"" approximation, in which the full magnetic force on the ions is conveyed to the neutrals, i.e., \citep[Equation~(27.8)]{1992phas.book.....S}. ."
+We cau solve Equation (1 1)) to obtain once rjGr) is known. this gives rgGe).," We can solve Equation \ref{up_down}) ) to obtain once $r_n(x)$ is known, this gives $r_B(x)$."
+" The compression ratio ry for both neutrals and magnetic field lines in the post-shock region is obtained by setting rj—ry= in Equation (16]). yielding Note that if 3,M?5»1. for a strong shock. In dimensional form. this is Note that for oblique C shocks. ry4rp dn the post-shock regiou. aud they both dependon the angle 0 between B aud v."," The compression ratio $r_f$ for both neutrals and magnetic field lines in the post-shock region is obtained by setting $r_B = r_f = r_n$ in Equation \ref{rB}) ), yielding Note that if $\beta{\cal M}^2 \gg 1$, for a strong shock, In dimensional form, this is Note that for oblique C shocks, $r_B \neq r_f$ in the post-shock region, and they both dependon the angle $\theta$ between $\mathbf{B}$ and $\mathbf{v}$."
+ Appendix A provides expressions for generalized rj(0) with Equation (A16)) and ryj(0) with Equation CA15)).," Appendix \ref{sec:appendix} provides expressions for generalized $r_f(\theta)$ with Equation \ref{rfQuadra}) ) and $r_{B,f}(\theta)$ with Equation \ref{rBfApprox}) )."
+" Combining Equatious (16)) and (12)). we obtain au ODEfor r,."," Combining Equations \ref{rB}) ) and \ref{drB2dx}) ), we obtain an ODEfor $r_n$ ."
+ The governing equation is, The governing equation is
+scale but. individual superclusters must be located. Hexibly within this network.,scale but individual superclusters must be located flexibly within this network.
+ The scale of the model is determined by the period. of oscillations., The scale of the model is determined by the period of oscillations.
+ To check the correspondence of the model to reality one can use small deviations of the correlation function from strict regularity expressed through parameters As; and Ass., To check the correspondence of the model to reality one can use small deviations of the correlation function from strict regularity expressed through parameters $\Delta_{21}$ and $\Delta_{32}$.
+ As we see (Paper LL). deviations of the observed. correlation function. from strict regularity are what is expected. for. models with constant scale and a certain scatter in individual superclusters positions.," As we see (Paper II), deviations of the observed correlation function from strict regularity are what is expected for models with constant scale and a certain scatter in individual superclusters positions."
+ This coincidence cannot be accidental. since a visual inspection of the distribution of real superclusters also shows the presence of such a network of superclusters and. voids (Paper L).," This coincidence cannot be accidental, since a visual inspection of the distribution of real superclusters also shows the presence of such a network of superclusters and voids (Paper I)."
+ The basic dillerence between the regular and. Voronoi models lies in the correlation of positions of high-density regions on large scales: in the regular model positions are correlated and in the Voronoi niocel not., The basic difference between the regular and Voronoi models lies in the correlation of positions of high-density regions on large scales: in the regular model positions are correlated and in the Voronoi model not.
+ lteal cluster samples are limited. to a certain. maximal distance. and clusters in the galactic equatorial zone are not seen due to obscuration.," Real cluster samples are limited to a certain maximal distance, and clusters in the galactic equatorial zone are not seen due to obscuration."
+ Thus real observational samples generally have ai double. conical shape., Thus real observational samples generally have a double conical shape.
+ To investigate the influence of such a sample shape we selected: from cubical cluster samples double conical subsaniples., To investigate the influence of such a sample shape we selected from cubical cluster samples double conical subsamples.
+ These subsamples have a conical height equal to half the size of the total sample. and the equatorial belt. (galactic latitudes less than 307) is exeluded.," These subsamples have a conical height equal to half the size of the total sample, and the equatorial belt (galactic latitudes less than $\pm
+30^{\circ}$ ) is excluded."
+ The volume of these subsamples is about one-quarter the volume of the whole cube. and they contain about one-quarter the number of clusters than is in the original sample.," The volume of these subsamples is about one-quarter the volume of the whole cube, and they contain about one-quarter the number of clusters than is in the original sample."
+ In Figure 5 we compare correlation functions of. two models — the regular net model and the random supereluster model calculated. for the whole cubical volume. and or à double conical volume.," In Figure 5 we compare correlation functions of two models – the regular net model and the random supercluster model – calculated for the whole cubical volume, and for a double conical volume."
+ The first. model is regular and the corresponding correlation function is oscillating: he second model is not., The first model is regular and the corresponding correlation function is oscillating; the second model is not.
+ This Figure shows that the oscillating properties of the sample are. preserved in. the smaller double conical sample., This Figure shows that the oscillating properties of the sample are preserved in the smaller double conical sample.
+ Only the error. corridor is arecr in the smaller samples. as expected.," Only the error corridor is larger in the smaller samples, as expected."
+ This similarity of the oscillating properties is due to the fact that the small samples are large enough to contain in sullicient. quantity he structural elements responsible for oscillations., This similarity of the oscillating properties is due to the fact that the small samples are large enough to contain in sufficient quantity the structural elements responsible for oscillations.
+ Spectra for models are. plotted in. Figures 2 4., Spectra for models are plotted in Figures 2 – 4.
+ Power spectra have three regions. corresponding (ο inhomogeneities with high. medium. anc low spatial frequeney (or small. medium. and lone wavelengths).," Power spectra have three regions, corresponding to inhomogeneities with high, medium, and low spatial frequency (or small, medium, and long wavelengths)."
+ This, This
+similar to Run 1.,similar to Run 1.
+ However. in Run 7 the temperatures are ower.," However, in Run 7 the temperatures are lower."
+ There are two reasons for this., There are two reasons for this.
+ First. because the disc is more extended. the material hitting the accretion shock at the boundary. of the disc has fallen less far into he gravitational potential well of the core and is therefore moving more slowly: it is therefore not heated so much in he accretion shock.," First, because the disc is more extended, the material hitting the accretion shock at the boundary of the disc has fallen less far into the gravitational potential well of the core and is therefore moving more slowly; it is therefore not heated so much in the accretion shock."
+ Second. again because the disc is more extended and therefore has lower surface density. but also »ecause it is cooler. its cooling radiation can escape more readily.," Second, again because the disc is more extended and therefore has lower surface density, but also because it is cooler, its cooling radiation can escape more readily."
+ This in turn is the main reason why the disc in ltun 7 ds more prone to fragmentation., This in turn is the main reason why the disc in Run 7 is more prone to fragmentation.
+ Since the rate of infall onto the dise greatly exceeds the rate of accretion onto the primary protostar (see below). the heating at the accretion shock again dominates the energy budget of the disc.," Since the rate of infall onto the disc greatly exceeds the rate of accretion onto the primary protostar (see below), the heating at the accretion shock again dominates the energy budget of the disc."
+ Further heating by compression and viscous dissipation ensures that the vertical temperature gradient is negative. dT/d|z|«0.," Further heating by compression and viscous dissipation ensures that the vertical temperature gradient is negative, $dT/d|z|<0$."
+ Llowever. because the outer disc is very cool. we only expect to see the 10jun feature in absorption from the inner disc.," However, because the outer disc is very cool, we only expect to see the $10\,\mu{\rm m}$ feature in absorption from the inner disc."
+ ‘The spikes superimposed on the temperature profiles are due, The spikes superimposed on the temperature profiles are due
+independent coustraius.,independent constrains.
+ The main results are suuimiaized below., The main results are summarized below.
+ 1) Our profile suuple is consistent with a beam shape that is a function of a. circular at à=0 and increasingly ME iu the latitudinal direction as à increases. as sugeested by Biges (1990).," 1) Our profile sample is consistent with a beam shape that is a function of $\alpha$, circular at $\alpha =0$ and increasingly compressed in the latitudinal direction as ${\alpha}$ increases, as suggested by Biggs (1990)."
+ ILowever. the data is equally consistent with the possibility that the beam is circular for all values of à.," However, the data is equally consistent with the possibility that the beam is circular for all values of ${\alpha}$."
+ 2) We identity three nested cones of eniüssion based on a normalized distribution of outer compoucuts., 2) We identify three nested cones of emission based on a normalized distribution of outer components.
+ The observed eap (05 607) in the distribution independently suggests three cones in the form of annular rings whose widths are typically about of the cone radii., The observed gap $\theta\gsim 60^{\circ}$ ) in the distribution independently suggests three cones in the form of annular rings whose widths are typically about of the cone radii.
+ We consider this circumstance as au important evidence for the nested-coue structure., We consider this circumstance as an important evidence for the nested-cone structure.
+ Auv further significant progress in such imnodelliug would necessarily need refined estimates of the observables. particularly the corce-widths.," Any further significant progress in such modelling would necessarily need refined estimates of the observables, particularly the core-widths."
+ We thank V. Racthakrishuan. Raja Nitvananda aud Joanna Rakin for fruitful discussions and for several sugecstious that have helped in improving the mamuscript.," We thank V. Radhakrishnan, Rajaram Nityananda and Joanna Rankin for fruitful discussions and for several suggestions that have helped in improving the manuscript."
+ We acknowledge Ashish Aseckar. D. Bhattacharva aud BR. Ramachandran for useful discussions and thank our referee. J. A. Cal. for critical comunents aud sugecstious.," We acknowledge Ashish Asgekar, D. Bhattacharya and R. Ramachandran for useful discussions and thank our referee, J. A. Gil, for critical comments and suggestions."
+ We acknowledge Ashish Aseckar. D. Bhattacharva aud BR. Ramachandran for useful discussions and thank our referee. J. A. Cal. for critical comunents aud sugecstious.O," We acknowledge Ashish Asgekar, D. Bhattacharya and R. Ramachandran for useful discussions and thank our referee, J. A. Gil, for critical comments and suggestions."
+ We acknowledge Ashish Aseckar. D. Bhattacharva aud BR. Ramachandran for useful discussions and thank our referee. J. A. Cal. for critical comunents aud sugecstious.OO," We acknowledge Ashish Asgekar, D. Bhattacharya and R. Ramachandran for useful discussions and thank our referee, J. A. Gil, for critical comments and suggestions."
+G8003) slit spectra with a dispersion of per pixel.,) slit spectra with a dispersion of per pixel.
+ Thev used Principal Compoucut Analvsis (PCA) to pre-oxocess their spectra. and reduce the uummber of input rodes.," They used Principal Component Analysis (PCA) to pre-process their spectra, and reduce the number of input nodes."
+ They obtained an accuracy of 2.31 spectral types over the full MIN range (OAD)., They obtained an accuracy of 2.34 spectral types over the full MK range (O–M).
+ ? used again IHouk plate naterial. digitized with the APM plate scanner. vielding a wavelength range of and a dispersion of per pixel.," \cite{Bailer-Jonesetal:1998a} used again Houk's plate material, digitized with the APM plate scanner, yielding a wavelength range of and a dispersion of per pixel."
+ Their best ANN configuration classified hese spectra with au eror distribution having a vpes. and the Iuminositv classification was correct for salnple spectra.," Their best ANN configuration classified these spectra with an error distribution having a types, and the luminosity classification was correct for sample spectra."
+ Recently. ? used a MBPN for derivation of the stellar yaralcters Tig. logg aud |Fe/TII] from moderate resolution per pixel) spectra.," Recently, \cite{Snideretal:2001} used a MBPN for derivation of the stellar parameters $T_{\mbox{\scriptsize eff}}$, $\log g$ and [Fe/H] from moderate resolution per pixel) spectra."
+ Althoueh their aim is to assign paramucter values to cach spectrum. while we as well as the above mentioned authors carried out classifications. we mclude their work in our review because Suider et al.," Although their aim is to assign parameter values to each spectrum, while we as well as the above mentioned authors carried out classifications, we include their work in our review because Snider et al."
+ applied. their technique to iietal-poor stars. which is also the object type we are mainly concerned with iu this paper.," applied their technique to metal-poor stars, which is also the object type we are mainly concerned with in this paper."
+ Snider ct al., Snider et al.
+" report classification accuracies of στ=135 190 ΤΙΝ. σι=0.25 0.30 ος aud o4,4,=0.15 0.20 ddex."," report classification accuracies of $\sigma_{T_{\mbox{\tiny eff}}}=135$ $150$ K, $\sigma_{\log
+ g}=0.25$ $0.30$ dex and $\sigma_{\mbox{\tiny [Fe/H]}}=0.15$ $0.20$ dex."
+ However. i appears from the upper paucl of heir Fie.," However, it appears from the upper panel of their Fig."
+ I tha subeiauts aud horizoutal xonch stars have been excluded from the sample of stars hey studied., 4 that subgiants and horizontal branch stars have been excluded from the sample of stars they studied.
+ A rough graphical analysis of their Fiso., A rough graphical analysis of their Fig.
+ 1 reveals that unlike iu real samples of stars eimiereius οιB.ο roni wide-angle spectroscopic surveys.which coutaiu subeiants and horizoutabbranch stars. their sample cau © Classified in logg with a similar precision bv dividing it iuto two classes “by laud”. that is. assigning logy=2.5 ο all stars with Tig<DODO IN. and lozg=L5 to all stars with Digg>SOOO TTS. Furthermore. it is questionable that there is any feature present in their set of spectra which does allow for a eravityv classification. since they used coutinuun divided spectra.," 4 reveals that unlike in real samples of stars emerging e.g. from wide-angle spectroscopic surveys,which contain subgiants and horizontal-branch stars, their sample can be classified in $\log g$ with a similar precision by dividing it into two classes “by hand”, that is, assigning $\log g=2.5$ to all stars with $T_{\mbox{\scriptsize eff}}<5000$ K, and $\log g=4.5$ to all stars with $T_{\mbox{\scriptsize eff}}>5000$ K. Furthermore, it is questionable that there is any feature present in their set of spectra which does allow for a gravity classification, since they used continuum divided spectra."
+ The Balmer jump. which Is a eravity indicator in cool stars. is therefore removed.," The Balmer jump, which is a gravity indicator in cool stars, is therefore removed."
+ Iu conclusion. while Snider ct al.," In conclusion, while Snider et al."
+ succeeded i sine ANNs for automated classification in τω aud [Fe/TI]. it roinaius to be demonstrated with a realistic sample that rectified moderate-resolutiou spectra inecd contain the information needed for à useful eravity classification.," succeeded in using ANNs for automated classification in $T_{\mbox{\scriptsize eff}}$ and [Fe/H], it remains to be demonstrated with a realistic sample that rectified moderate-resolution spectra indeed contain the information needed for a useful gravity classification."
+ ANN techniques iux “classical” statistical methods such as Bayes aux uaa cost rule classifications ofteu perform equally well. iu terms of e.g. minimising the total uuniber of misclassificatious.," ANN techniques and “classical” statistical methods such as Bayes and minimum cost rule classifications often perform equally well, in terms of e.g. minimising the total number of misclassifications."
+ In the present work. we cluplov statistical methods. because their mathematical properties are well-studied. and the formulation of classification rules im the framework of mathematical statistics makes them very transparent.," In the present work, we employ statistical methods, because their mathematical properties are well-studied, and the formulation of classification rules in the framework of mathematical statistics makes them very transparent."
+ Before we go into details of the methods we developed (Sect. 27?)).," Before we go into details of the methods we developed (Sect. \ref{Sect:Autoclass}) ),"
+ we give a brief overview of the Ibuuburg/ESO Survey CIES) in Sect. ??..," we give a brief overview of the Hamburg/ESO Survey (HES) in Sect. \ref{Sect:HES},"
+ for better readibilitv., for better readibility.
+ Tn Sect., In Sect.
+ TT we investigate the classification performance for stars in the effective temperature range 5200].cTigGSOOTS achievable iu the TES. by a simulation study.," \ref{Sect:Performance} we investigate the classification performance for stars in the effective temperature range $5200\,\mbox{K}<T_{\mbox{\scriptsize
+ eff}}<6800\,\mbox{K}$ achievable in the HES, by a simulation study."
+ We sunnunarnze our conclusions im Sect. 6.., We summarize our conclusions in Sect. \ref{Sect:Conclusions}. .
+ The TES (27) is an objective-prisin survey designed to select bright (12.592By 17.5) quasars in the southern," The HES \citep{hespaperI,hespaperIII} is an objective-prism survey designed to select bright $12.5 \gtrsim B_J \gtrsim 17.5$ ) quasars in the southern"
+When analyzing velocity dispersion and surface density data of star clusters. confusion over. nomenclature often arises.,"When analyzing velocity dispersion and surface density data of star clusters, confusion over nomenclature often arises."
+ Therefore we strictly stick to the following three terms:profiles.. andfemplates.," Therefore we strictly stick to the following three terms:, and."
+. Phe first. will be used. for the data. which in our case comes from N-bocly computations but might also originate [rom observations.," The first will be used for the data, which in our case comes from $N$ -body computations but might also originate from observations."
+ The term. will be used. for physically motivated distribution. functions such as those proposed. bv Wine(1966) or Wilson(1975)., The term will be used for physically motivated distribution functions such as those proposed by \citet{King66} or \citet{Wilson75}.
+. The word covers the analvtic expressions which have been found empirically and which do not originate [rom a physical derivation by being solutions of the Collisionless Boltzmann Equation. (e.g. nine1962:Elson.Fall&FreemanLOST:Laueretal. 1995)).," The word covers the analytic expressions which have been found empirically and which do not originate from a physical derivation by being solutions of the Collisionless Boltzmann Equation (e.g. \citealt{King62, Elson87,
+Lauer95}) )."
+ The word template hereby emphasises their artificial Further confusion arises about the term., The word template hereby emphasises their artificial Further confusion arises about the term.
+radius.. Again there is a physically motivated. ancl an empirical version., Again there is a physically motivated and an empirical version.
+ The former is the radius at which the internal eravitational acceleration equals the tidal acceleration fron the host galaxy (see e.g. Binney&“Tremaine 2008))., The former is the radius at which the internal gravitational acceleration equals the tidal acceleration from the host galaxy (see e.g. \citealt{Binney87}) ).
+ This will be named in our text., This will be named in our text.
+ Ehe empirical tidal radius is defined through the eut-olf radius of the mocels and templates which are fitted to the profiles., The empirical tidal radius is defined through the cut-off radius of the models and templates which are fitted to the profiles.
+ We will refer to this cut-olf radius as here and would like to emphasise that Jacobi radius and edge radius are two cdilferent. concepts and therefore are not necessarily equal., We will refer to this cut-off radius as here and would like to emphasise that Jacobi radius and edge radius are two different concepts and therefore are not necessarily equal.
+ In fact. we will show that the two are often significantly dillerent and that conclusions from a itted edge radius should be drawn with Note also that the results of this investigation cannot oe easily scaled. to more massive e&lobular. clusters. as the relative importance of tidal features depends on the mass of he cluster. since the mass in the tidal debris is proportional o the mass-loss rate of the cluster and the mass-loss rate does not scale linearly with cluster mass. M. but rather with A47 (Baumegaredt&Makino2003).," In fact, we will show that the two are often significantly different and that conclusions from a fitted edge radius should be drawn with Note also that the results of this investigation cannot be easily scaled to more massive globular clusters, as the relative importance of tidal features depends on the mass of the cluster, since the mass in the tidal debris is proportional to the mass-loss rate of the cluster and the mass-loss rate does not scale linearly with cluster mass, $M$, but rather with $M^{1/4}$ \citep{Baumgardt03}."
+". Aloreover. even though we have some concentrated clusters in our sample this investigation focuses on less concentrated clusters. with a ratio of projected. half mass radius to Jacobi radius between 0.1κἐς03 because Baumegardtetal.(2010) showed that clusters with ratios of feu,/I;«0.05 cannot be properly fitted by either Wing(1962) templates or Wine(1966) models as their ratio of eiufI; is too high."," Moreover, even though we have some concentrated clusters in our sample this investigation focuses on less concentrated clusters, with a ratio of projected half mass radius to Jacobi radius between $0.1 <
+R_{hp}/R_J < 0.3$ because \citet{Baumgardt10} showed that clusters with ratios of $R_{hp}/R_J < 0.05$ cannot be properly fitted by either \citet{King62} templates or \citet{King66} models as their ratio of $R_{hp}/R_J$ is too high."
+ Therefore. our. conclusions mainly hold for less concentrated clusters which Baunmgarelt named thepopulalion.. and which are likely to be on the main sequence of cluster evolution (Ixüpper.INroupa&Baumegarct2008).. the final and. universal stage of cluster evolution after core collapse. where the ratio Rip{τμ only depends on the mass of the In the following sections we will first give some details about our data set. before. we explain how we extract the velocity clispersion and surface density. profiles.," Therefore, our conclusions mainly hold for less concentrated clusters which \citet{Baumgardt10} named the, and which are likely to be on the main sequence of cluster evolution \citep{Kuepper08b}, the final and universal stage of cluster evolution after core collapse, where the ratio $R_{hp}/R_J$ only depends on the mass of the In the following sections we will first give some details about our data set before we explain how we extract the velocity dispersion and surface density profiles."
+ Thereafter we brielly describe the analytical templates which are used in this investigation and specify our fitting moethocd., Thereafter we briefly describe the analytical templates which are used in this investigation and specify our fitting method.
+ The velocity dispersion and surface density profiles are mainty taken from our recent investigation of star clusters ancl their tidal tails (Ixüpperetal.2010)., The velocity dispersion and surface density profiles are mainly taken from our recent investigation of star clusters and their tidal tails \citep{Kuepper09}.
+.. For this analysis we computed a comprehensive set of star clusters on various orbits about a Milky-Way In addition to the 16 clusters taken from Kipperοἱ (2010).. we computed 4 concentrated clusters which are of specific interest. for this investigation.," For this analysis we computed a comprehensive set of star clusters on various orbits about a Milky-Way In addition to the 16 clusters taken from \citet{Kuepper09}, we computed 4 concentrated clusters which are of specific interest for this investigation."
+ ALL clusters were set-up using the publicly available code (Ixüàpppoer et al..," All clusters were set-up using the publicly available code (Küppper et al.,"
+ in prep.).," in prep.),"
+ and integrated over time with the W-bocly code CXarseth2003). on the computers at ALLA Bonn (Fukushige.Makino&Ixawai 2005)., and integrated over time with the $N$ -body code \citep{Aarseth03} on the computers at AIfA Bonn \citep{Fukushige05}.
+. The code was modified so that the Milkv-M'av potential suggested:by Allen&Santillan(1991)— could be used., The code was modified so that the Milky-Way potential suggestedby \citet{Allen91} could be used.
+ Due to its analytic form this potential is useful for calculating the Jacobi radius. £7. of the star clusters.," Due to its analytic form this potential is useful for calculating the Jacobi radius, $R_J$, of the star clusters."
+" ""Throughout the computations 27 was evaluated. using equation 12 of Ixüpperetal.(2010)... ie. where Gis the gravitational constant. AZ is the cluster mass. © is the angular velocity of the cluster on its orbit about the galaxy. and O7/OL, is the second. derivative of the gravitational potential of the galaxy with respect to the ealactocentric radius. fe."," Throughout the computations $R_J$ was evaluated using equation 12 of \citet{Kuepper09}, i.e. where $G$ is the gravitational constant, $M$ is the cluster mass, $\Omega$ is the angular velocity of the cluster on its orbit about the galaxy, and $\partial ^2\Phi/\partial R_G^2$ is the second derivative of the gravitational potential of the galaxy with respect to the galactocentric radius, $R_G$."
+ We always determine the Jacobi radius iteratively. assuming first a mass of the cluster. then applving eq. 10...," We always determine the Jacobi radius iteratively, assuming first a mass of the cluster, then applying eq. \ref{eq:R_J}. ."
+ With this estimate we determine the mass of the cluster again. counting all masses around the cluster centre within 27. and compute the Jacobi radius once," With this estimate we determine the mass of the cluster again, counting all masses around the cluster centre within $R_J$ , and compute the Jacobi radius once"
+the decreasing criieal svnchirotron frequency. which both increase the radio emission: low effective UV dust opacity leads to the decreasing FIR emission. which balances the decrease in radio emission caused by CW escape for lower surface densitv galaxies.,"the decreasing critical synchrotron frequency, which both increase the radio emission; low effective UV dust opacity leads to the decreasing FIR emission, which balances the decrease in radio emission caused by CR escape for lower surface density galaxies."
+ In fact. this kind of calorimetry theory (hat galaxies act in a reasonable approximation as calorimeters for the stellar UV radiation aud for the energy [lux of the CHR. electrons was proposed 20 vears ago (Voelk 1989).," In fact, this kind of calorimetry theory that galaxies act in a reasonable approximation as calorimeters for the stellar UV radiation and for the energy flux of the CR electrons was proposed 20 years ago (Voelk 1989)."
+ Although5 the molecular 5gas connection to the FIR-RC correlation in 5galaxies is proposed by Muregia5 (2005). we demonstrated here that the addition of the total dense molecular eas tracecl by ICN appears to be globally insignificant in providing the dense star-lorminge gas connection between the FIR. (SFR) and RC (the(μον.Lin. )) model. Equations (2). (5). and D(3)).," Although the molecular gas connection to the FIR-RC correlation in galaxies is proposed by Murgia (2005), we demonstrated here that the addition of the total dense molecular gas traced by HCN appears to be globally insignificant in providing the dense star-forming gas connection between the FIR (SFR) and RC (the, ) model, Equations (2), (5), and B(3))."
+" Yet. the contribution of CO is non-neegligible and appears to be more important than ICN,"," Yet, the contribution of CO is non-negligible and appears to be more important than HCN."
+ Nevertheless. the star formation materials. particularly the dense gas (UCN). definitely. play an important role in the obvious connection of SFR to both the FIR and RC emission.," Nevertheless, the star formation materials, particularly the dense gas (HCN), definitely play an important role in the obvious connection of SFR to both the FIR and RC emission."
+ Both ICN and RC contribute nearly equally to FUR ancl there is no additional contribution to FUR by adding CO (theIR(Lyex..μοι. Leo)) model. eqs. (," Both HCN and RC contribute nearly equally to FIR and there is no additional contribution to FIR by adding CO (the, ) model, eqs. ("
+4). (7) and D(3)).,"4), (7) and B(3))."
+ Thus. CO seems to be the least important one among all the four parameters RC. FIR. CO. and HCN in relating to SFR even though they all show quite good correlations with each other.," Thus, CO seems to be the least important one among all the four parameters RC, FIR, CO, and HCN in relating to SFR even though they all show quite good correlations with each other."
+ In conventional models. (he RC emission produced fom synchrotron radiation is related wilh massive star formation via several steps: star formation —supernovae and SNRs —relativistic electrons accelerated by SNRs.," In conventional models, the RC emission produced from synchrotron radiation is related with massive star formation via several steps: star formation $\rightarrow$ supernovae and SNRs $\rightarrow$ relativistic electrons accelerated by SNRs."
+ So. too. is the FIR. emission related: star [ormation—- UV raciation—absorpüon and re-emission bv dust enshrouding new-born stars.," So, too, is the FIR emission related: star $\rightarrow$ UV $\rightarrow$ absorption and re-emission by dust enshrouding new-born stars."
+ These should all be expected to be the dominant. FUR and. RC emission in active stur-Iorming galaxies., These should all be expected to be the dominant FIR and RC emission in active star-forming galaxies.
+ Nevertheless. CR. electrons involve with other radiation processes. generating mostly the background RC emission that is not related to massive star formation (the general interstellar radiation field).," Nevertheless, CR electrons involve with other radiation processes, generating mostly the background RC emission that is not related to massive star formation (the general interstellar radiation field)."
+ Additionally. the old stars also heat up the dust and additionally contribute to the general interstellar radiation field.," Additionally, the old stars also heat up the dust and additionally contribute to the general interstellar radiation field."
+ The contribution of the general interstellar racliation field to the total FIR. emission in galaxies might be significant at the eend. where the general infrared interstellar raciation fiekl is comparable to or even more dominant than the FIR radiation from the active star formation.," The contribution of the general interstellar radiation field to the total FIR emission in galaxies might be significant at the end, where the general infrared interstellar radiation field is comparable to or even more dominant than the FIR radiation from the active star formation."
+ The temporal and spatial inconsistencies between these processes could be the obstacles to understand fly the physical origin of the FIR and RC correlation., The temporal and spatial inconsistencies between these processes could be the obstacles to understand fully the physical origin of the FIR and RC correlation.
+ Given the obvious star lormation connection of both FIR and RC ancl the important role of the star formation, Given the obvious star formation connection of both FIR and RC and the important role of the star formation
+Rotation curves plaved an esseutial role iu establishing the need for dark matter halos (Rubin. Thounarcd. Ford 1978: Dosma 1978).,"Rotation curves played an essential role in establishing the need for dark matter halos (Rubin, Thonnard, Ford 1978; Bosma 1978)."
+ Απο Newtonian eravity holds. dark matter is required in all spirals to explain the asviiptotically flat rotation curves observed at laree radi.," Assuming Newtonian gravity holds, dark matter is required in all spirals to explain the asymptotically flat rotation curves observed at large radii."
+ The need for dark matter is less obvious at sinall raclii., The need for dark matter is less obvious at small radii.
+ Rotation curves provide stroug coustraints on the radial shape of the eravitational potential of ealaxies., Rotation curves provide strong constraints on the radial shape of the gravitational potential of galaxies.
+ This translates to a coustraiut on the distribution the various mass components. both Iuninous aud dark.," This translates to a constraint on the distribution the various mass components, both luminous and dark."
+ Though the coustrait on the potential is strong. there can be degeneracies between the contributions of the various niass components.," Though the constraint on the potential is strong, there can be degeneracies between the contributions of the various mass components."
+ Iu particular. the mass of the stellar disk can often be traded off against tle mass in dark matter. leaving cousiderable room or differences of opinion about disk masses aud halo distributions.," In particular, the mass of the stellar disk can often be traded off against the mass in dark matter, leaving considerable room for differences of opinion about disk masses and halo distributions."
+ This disk-halo degeueracyv has plagued the field for a good while now., This disk-halo degeneracy has plagued the field for a good while now.
+ Consequently. there is at present a wide diversity of opinion as to how important he huninous and dark mass are at stuall radii," Consequently, there is at present a wide diversity of opinion as to how important the luminous and dark mass are at small radii."
+ These range between two casily ideuti&able extremes: maximal disks aud. cuspy halos., These range between two easily identifiable extremes: maximal disks and cuspy halos.
+ Iu between these is a ranee occupied by models with stellar miass-to-helt ratios which are alausible for the composite stellar populations of spiral galaxies., In between these is a range occupied by models with stellar mass-to-light ratios which are plausible for the composite stellar populations of spiral galaxies.
+" Which of hese various optious one prefers is larecly a matter of how one weighs the evidence,", Which of these various options one prefers is largely a matter of how one weighs the evidence.
+ Dynamicists find much to recommend imiaxinal disks., Dynamicists find much to recommend maximal disks.
+ Cosimologists expect cuspy dark matter halos. for which the disk mass must be minimal.," Cosmologists expect cuspy dark matter halos, for which the disk mass must be minimal."
+" Those who study stellaz populations prefer intermediate disk masses,", Those who study stellar populations prefer intermediate disk masses.
+model of solar mass and CZ.Y) values of (0.0001.0.24). (0.001.0.24). (0.015.020). (0.05.0.30) and (0.1.0.30). other physical and numerical parameters remaining fixed.,"model of solar mass and $(Z,Y)$ values of (0.0001,0.24), (0.001,0.24), (0.018,0.29), (0.05,0.30) and (0.1,0.30), other physical and numerical parameters remaining fixed."
+ The results are presented in Fig., The results are presented in Fig.
+ 4 bv complete. continuous tracks in the ΠΠ. diagram.," \ref{fig:metal} by complete, continuous tracks in the H-R diagram."
+ We note (hat an increase in metallicity has a similar ellect to a decrease in the initial stellar mass: luminosities are lower and (he durations of evolutionary. phases are longer., We note that an increase in metallicity has a similar effect to a decrease in the initial stellar mass: luminosities are lower and the durations of evolutionary phases are longer.
+ For example. the AIS phase lasts wp to over 3 limes longer. when Z increases [rom 10.| to 0.1.," For example, the MS phase lasts up to over 3 times longer, when $Z$ increases from $10^{-4}$ to $0.1$."
+ This result is mostly the consequence of the dependence of opacity on composition; at a lower metallicity. the opacily decreases. the star is able (ο radiate away ils energv wilh greater elliciency. the stellar luminosity is therefore higher ancl timescales are correspondingly shorter.," This result is mostly the consequence of the dependence of opacity on composition; at a lower metallicity, the opacity decreases, the star is able to radiate away its energy with greater efficiency, the stellar luminosity is therefore higher and timescales are correspondingly shorter."
+ Apart [rom the apparent shift of the evolutionary tracks in the H-R diagram. and the different timescales. metallicity also affects the final masses.," Apart from the apparent shift of the evolutionary tracks in the H-R diagram, and the different timescales, metallicity also affects the final masses."
+ For A;=1AL.. a final mass of 0.57AL. was obtained for the lowest metallicity (7=0.0001). and 0.52A. for the highest one (Z= 0.1). as compared with 0.55AL... obtained for solar metallicityan overall spread of almost1054.," For $M_i=1\ \Msun$, a final mass of $0.57\ \Msun$ was obtained for the lowest metallicity $Z=0.0001$ ), and $0.52\ \Msun$ for the highest one $Z=0.1$ ), as compared with $0.55\ \Msun$, obtained for solar metallicity—an overall spread of almost."
+.. We consider Population I (Pop., We consider Population I (Pop.
+I) and Population II (Pop.,I) and Population II (Pop.
+II) stars. adopting metallicities ol Z=0.01 and Z=0.001. respectively. and initial masses in the range 0.25—9M... leading to cooling WDs.,"II) stars, adopting metallicities of $Z=0.01$ and $Z=0.001$, respectively, and initial masses in the range $0.25-9\,\Msun$, leading to cooling WDs."
+ The complete evolutionary (racks are shown in the (wo panels οἱ Fig. 5.., The complete evolutionary tracks are shown in the two panels of Fig. \ref{fig:hrds}.
+ Timescales and the final WD masses and composition are given in (he accompanying Table 1., Timescales and the final WD masses and composition are given in the accompanying Table 1.
+ It should be noted that the AIS and RGB cdurations as shown in the table depend on the definition of the MS-turnoff point and beginning of the RGB. which involves some arbitrariness.," It should be noted that the MS and RGB durations as shown in the table depend on the definition of the MS-turnoff point and beginning of the RGB, which involves some arbitrariness."
+ The criterion we use for the MS turnoll is as follows: let /4 be the time when X. has decreased below 10. *: let ry=logT;yy(ty) and yy=logL(/4)., The criterion we use for the MS turnoff is as follows: let $t_1$ be the time when $X_c$ has decreased below $10^{-6}$ ; let $x_1=\log T_{eff}(t_1)$ and $y_1=\log L(t_1)$.
+ The turnolf time fs is the earliest time for which the distance between the points Lr»=logTipr(/o). yo=log L(fa)] and tra. yf in the logΤε. log L] plane exceeds 0.1.," The turnoff time $t_2$ is the earliest time for which the distance between the points $x_2=\log T_{eff}(t_2)$, $y_2=\log L(t_2)$ ] and $x_1$, $y_1$ ] in the $\log T_{eff}$, $\log L$ ] plane exceeds 0.1."
+ Similar criteria are used for other transitions between evolutionary stages., Similar criteria are used for other transitions between evolutionary stages.
+ Time scales depend strongly ou composition. especially on Z. decreasing wilh decreasing Z.," Time scales depend strongly on composition, especially on $Z$, decreasing with decreasing $Z$."
+ Given dillerences in composition adopted in different studies. as well as dillerences in criteria delining evolutionary stages. a precise comparison between models is difficult to achieve.," Given differences in composition adopted in different studies, as well as differences in criteria defining evolutionary stages, a precise comparison between models is difficult to achieve."
+ Nevertheless. we lind excellent agreement. for example. between our low-mass Pop.," Nevertheless, we find excellent agreement, for example, between our low-mass Pop."
+"lI models and corresponding ones caleulated by others: for the 0.8 M. and LAL. models. we find 74,5=1.44LOM vr and 6.02x10? vr. respectively (see Table 1).","II models and corresponding ones calculated by others: for the 0.8 $\Msun$ and 1$\Msun$ models, we find $\tau_{MS}=1.44\times10^{10}$ yr and $6.02\times10^9$ yr, respectively (see Table 1),"
+(2 <1 Reuzini(2006))). Sssojn>Dandy) (Simailetal.2002:BlainTacconi2008).. (22:Williamsetal.2009:I&otibüne,"$z > 1$ $< 1$ \citet{Renz06}) $S_{850\mu m} > 3$ \citep{Smai02, Blai02, Tacc08}. \citep[$z > 2$;][]{Dadd05, Krie08, Toft09, vDok08, Kurk09, Coll09, Will09, Koti09},"
+n2009).. (2>5:Riech- IRAM 3011 telescope with the bolometer array NAMDBO (Lestradeetal.2009).," \citep[$z > 5$;][]{Riec10, Capa11}
+ $S_{1200 \mu m}=30 \pm 3$ IRAM 30m telescope with the bolometer array MAMBO \citep{Lest09}."
+.. Although no optical counterpart was known for MMISI23]5938. its redshift :.=3.92960+0.00013 was spectroscopically measured through the detection of several strong CO aud CT cluission lines with the wide spectrometer EMIB also at this telescope (Lestradeetal.2010).," Although no optical counterpart was known for MM18423+5938, its redshift $z=3.92960\pm 0.00013$ was spectroscopically measured through the detection of several strong CO and CI emission lines with the wide spectrometer EMIR also at this telescope \citep{Lest10}."
+. These line detections indicate a massive reservoir of molecular gas in a galaxv oulv 1.6 Cer after the big bang., These line detections indicate a massive reservoir of molecular gas in a galaxy only 1.6 Gyr after the big bang.
+ The observed CO ladder (CO(1-3). CO(G6-5) aud CO(t1-6). CO(9-8) upper lait) peaks around 7=5. indicative of a starburst-powered SAIG that is not douinanutlv excited by au active nucleus (AGN).," The observed CO ladder (CO(4-3), CO(6-5) and CO(7-6), CO(9-8) upper limit) peaks around $J=5$, indicative of a starburst-powered SMG that is not dominantly excited by an active nucleus (AGN)."
+ Recently. bright. strongly leused SAIGs have been discovered iun large area surveys (Vieiractal.2010:Neercloctal.2010:Fraver 2011).," Recently, bright, strongly lensed SMGs have been discovered in large area surveys \citep{Viei10, Negr10, Fray11}."
+" Even compared to these. NNIISI23|5938 is very bright. and its well established redshift makes if a prime tarect for high-resolution observations to study the cause of its ligh brightness. to understand its place in this enmereme population. to verity the hvpothesis that it is eravitationally lensed. aud to study further the similarities of NAMIS123|5938 with the other very bright SMC SMNLJ2135-0102. at 2=2.3 (58595,7106 niv) which has been shown to be stronely lensed (Swinhauk.etal. 2010).."," Even compared to these, MM18423+5938 is very bright, and its well established redshift makes it a prime target for high-resolution observations to study the cause of its high brightness, to understand its place in this emerging population, to verify the hypothesis that it is gravitationally lensed, and to study further the similarities of MM18423+5938 with the other very bright SMG SMMJ2135-0102 at $z=2.3$ $S_{850\mu m}$ =106 mJy) which has been shown to be strongly lensed \citep{Swin10}. ."
+ Iu this Letter we report observations of the CO(I1- and CO(2-1) line eimissious in MMISI23|5938 to 0.67. resolution using the Expanded Verv Large. Array (EVLA: Perleyetal. (2011))). as well as inulticolor optical and ucar-IR imagine using the Cauada-Frauce- Telescope (CFUT).," In this Letter we report observations of the CO(1-0) and CO(2-1) line emissions in MM18423+5938 to $0.6''$ resolution using the Expanded Very Large Array (EVLA; \citet{Perl11}) ), as well as multicolor optical and near-IR imaging using the Canada-France-Hawaii Telescope (CFHT)."
+" These observations coufirii that AIATLS123,5938 is leused.providing a uuique opportunity to study the eas clvnamics on small physical scale with further high sensitivity. aud lich spatial"," These observations confirm that MM18423+5938 is lensed,providing a unique opportunity to study the gas dynamics on small physical scale with further high sensitivity and high spatial"
+elliptical isophotes leading to large residuals.,elliptical isophotes leading to large residuals.
+ Also. the fitted surface brightness profiles. lead to unphysical parameters for the fit profiles (exponential. de Vaucouleurs. Ixing-like core mocel and. Hubble-Ievnolds). show that this component is irregular in its whole extent.," Also, the fitted surface brightness profiles, lead to unphysical parameters for the fit profiles (exponential, de Vaucouleurs, King-like core model and Hubble-Reynolds), show that this component is irregular in its whole extent."
+ The group's gravitational potential is dominated by the eroup's dark matter halo and we assume that the light traces the mass., The group's gravitational potential is dominated by the group's dark matter halo and we assume that the light traces the mass.
+ Hf the IGL is due to stripped stars that ave distributed following the group's dark matter dominad potential. the irregular shape of the IGL lets us conclude that the groups. are definitely not. virialized structures.," If the IGL is due to stripped stars that are distributed following the group's dark matter dominated potential, the irregular shape of the IGL lets us conclude that the groups are definitely not virialized structures."
+ Thev would be in a compact configurations for long enough to have part of the stars stripped. [rom the ealaxics aud dispersed in the group potential. but not enough to ect in ecquilibrium.," They would be in a compact configurations for long enough to have part of the stars stripped from the galaxies and dispersed in the group potential, but not enough to get in equilibrium."
+ The non-equilibrium situation. together with the small number of galaxies. makes the estimate of group cdvnamical parameters. such as mass and crossing time. very uncertain and sensitive to small variations of the input data. as mentioned in section 4.2..," The non-equilibrium situation, together with the small number of galaxies, makes the estimate of group dynamical parameters, such as mass and crossing time, very uncertain and sensitive to small variations of the input data, as mentioned in section \ref{grpparm}."
+ An important parameter is the total mass of the eroup. which can only be determined dynamically: using the member galaxies. due to the inexistence of other determinations. such as through X-ray observations.," An important parameter is the total mass of the group, which can only be determined dynamically, using the member galaxies, due to the inexistence of other determinations, such as through X-ray observations."
+ In the case of compact groups of galaxies. the dynamical mass estimators can only be used as an indicator. since the structure. as we showed before. is not. virialized.," In the case of compact groups of galaxies, the dynamical mass estimators can only be used as an indicator, since the structure, as we showed before, is not virialized."
+" Also. they are only sensitive to the mass inside the area occupied. by the galaxies (the ""galaxies circle”). which in the case of a compact group. can leave a considerable part of the mass unaccounted for."," Also, they are only sensitive to the mass inside the area occupied by the galaxies (the “galaxies circle”), which in the case of a compact group, can leave a considerable part of the mass unaccounted for."
+ 2? studied the properties of compact groups. of ealaxies identified in a mock galaxv catalogue based on the Millennium Run simulation (?).., \citet*{mcc08} studied the properties of compact groups of galaxies identified in a mock galaxy catalogue based on the Millennium Run simulation \citep{spr05}.
+ These authors found that the member galaxies are in the inner part of the common dark matter halo of the eroup and that there is no correlation between the halo virial radius and how the galaxies are Concentrated towards the center of the group., These authors found that the member galaxies are in the inner part of the common dark matter halo of the group and that there is no correlation between the halo virial radius and how the galaxies are concentrated towards the center of the group.
+ We could approach this question in two extreme wavs., We could approach this question in two extreme ways.
+ One is assuming that the dark matter halo clominates he mass in the whole extent of the group anc that the galaxies do not contribute significantly to it., One is assuming that the dark matter halo dominates the mass in the whole extent of the group and that the galaxies do not contribute significantly to it.
+" In this case he dark matter mass could be directly related to the IGL uminosityv. so that the fraction of mass unaccounted would »e proportional to the fraction of Lib outside the ""galaxies circle""."," In this case the dark matter mass could be directly related to the IGL luminosity, so that the fraction of mass unaccounted would be proportional to the fraction of IGL outside the “galaxies circle”."
+ Phe missing fraction. in this case. would be about 60 o of the total mass.," The missing fraction, in this case, would be about 60 to of the total mass."
+ The second approach is to assume hat the galaxies dominate the mass in the inner part. of he group and that we have a constant AZ/L in the whole structure., The second approach is to assume that the galaxies dominate the mass in the inner part of the group and that we have a constant $M/L$ in the whole structure.
+ In this case the missing fraction would not be arger than of the total mass., In this case the missing fraction would not be larger than of the total mass.
+ We are. not. able to pin. down a number for the unaccounted mass. as we are not able to pin down a number or the total mass itself using dynamical estimators.," We are not able to pin down a number for the unaccounted mass, as we are not able to pin down a number for the total mass itself using dynamical estimators."
+ The non-equilibrium state of the compact groups. together with he fraction of the total mass that can be measured. by dynamical mass indicators. prevent us from a reliable mass determination.," The non-equilibrium state of the compact groups, together with the fraction of the total mass that can be measured by dynamical mass indicators, prevent us from a reliable mass determination."
+ Other mass estimators. as X-rays. assunie hat the structure is in equilibrium. what can lead: to he same kind of uncertainties.," Other mass estimators, as X-rays, assume that the structure is in equilibrium, what can lead to the same kind of uncertainties."
+ For lensing estimates the small clistance of the groups and their low masses turn the detection of the shear signal. if possible. very. complicated.," For lensing estimates the small distance of the groups and their low masses turn the detection of the shear signal, if possible, very complicated."
+ A recent work on the Coma cluster (?).. which is about the same distance as most of the LCGs. determined. its mass using weak lensing detected: with deep. large field. of view. images from the SDSS.," A recent work on the Coma cluster \citep{kub07}, which is about the same distance as most of the HCGs, determined its mass using weak lensing detected with deep, large field of view, images from the SDSS."
+ Even though the distances are similar. the difference in total mass between Coma and the LCGs. would still be a problem for detecting the shear signal.," Even though the distances are similar, the difference in total mass between Coma and the HCGs, would still be a problem for detecting the shear signal."
+ Xecording to ?.. the Εςας would need a distance 10 times larger to produce a detectable lensing signal.," According to \citet{men94b}, the HCGs would need a distance 10 times larger to produce a detectable lensing signal."
+ In a qualitative wav (our sample is too small for a more quantitative result). we could show that the fraction of LGL agrees with other dynamical evolution indicators such as the fraction of early-type galaxies ancl the crossing time.," In a qualitative way (our sample is too small for a more quantitative result), we could show that the fraction of IGL agrees with other dynamical evolution indicators such as the fraction of early-type galaxies and the crossing time."
+ The crossing time. as the mass. is sensitive to small variations on the input parameters. as we showed for the case of the previously studied group LCC ss.," The crossing time, as the mass, is sensitive to small variations on the input parameters, as we showed for the case of the previously studied group HCG 88."
+ We can also see that the magnitude dillerence between the first ancl the second-ranked ealaxies does not clearly correlate with any of the other indicators., We can also see that the magnitude difference between the first and the second-ranked galaxies does not clearly correlate with any of the other indicators.
+ ? shows that the crossing time correlates with the fraction of spirals (complementary to the fraction of early-type galaxies) and Amy. (figures 5 and 6 in their paper). but for objects in bins of crossing time.," \citet{hic92} shows that the crossing time correlates with the fraction of spirals (complementary to the fraction of early-type galaxies) and $\Delta m_{12}$ (figures 5 and 6 in their paper), but for objects in bins of crossing time."
+ The correlation using individual objects was not shown., The correlation using individual objects was not shown.
+ So the correlation of Amys. on average. could eventually be seen when we have a considerably larger saniple.," So the correlation of $\Delta m_{12}$, on average, could eventually be seen when we have a considerably larger sample."
+ The new ciscordant velocity of LCG 15€. which has a much larger radial velocity. (by 2720kms+) than the median. velocity of the group. is a point to be considered with care.," The new discordant velocity of HCG 15C, which has a much larger radial velocity (by $2720~{\rm km~s^{-1}}$ ) than the median velocity of the group, is a point to be considered with care."
+ Galaxies in Εςος are bright enough for a good recession velocity determination and both sources. the UZC ane 2.. have reliable redshift’ determinations.," Galaxies in HCGs are bright enough for a good recession velocity determination and both sources, the UZC and \citet{hic92}, have reliable redshift determinations."
+ Altering the original group configuration. as excluding LCG. 15€ (due to its new published recession. velocity) or including LCG 51€ (extending the definition of “discordant velocity”) puts a challenge on defining which is the “correct” configuration for cach of the groups.," Altering the original group configuration, as excluding HCG 15C (due to its new published recession velocity) or including HCG 51C (extending the definition of “discordant velocity”) puts a challenge on defining which is the “correct” configuration for each of the groups."
+ In figure οr we can see the location of cach of the two configurations for both systems., In figure \ref{figcorr} we can see the location of each of the two configurations for both systems.
+ For HCG 15. the relation between FE and foross leaves all the points too close to the line for any consideration. as so does the relation between foc. and Pen. but in the relation between Lega: and Lye. the quintet configuration shows a better agreement with the main trend. of the dynamical evolution indicators.," For HCG 15, the relation between $F_{Egal}$ and $t_{cross}$ leaves all the points too close to the line for any consideration, as so does the relation between $t_{cross}$ and $F_{IGL}$, but in the relation between $F_{Egal}$ and $F_{IGL}$, the quintet configuration shows a better agreement with the main trend of the dynamical evolution indicators."
+ La the case of HCG 51. the relation between Lever and Loss. ad. between A404; and Lye would favor the sextet configuration. while the relation between feyar and Lye. would favor the quintet one.," In the case of HCG 51, the relation between $F_{Egal}$ and $t_{cross}$ and between $t_{cross}$ and $F_{IGL}$ would favor the sextet configuration, while the relation between $F_{Egal}$ and $F_{IGL}$, would favor the quintet one."
+" Of course this kind. of ""selection? is very speculative. since both configurations mostly agree between the errors."," Of course this kind of “selection” is very speculative, since both configurations mostly agree between the errors."
+ In a possible more quantitative analysis. we can be able to extrapolated this kind. of relation as a test to verify if a ealaxy that has an accordant redshift. but do not seem to be involved in the interactions occurring in the eroup. is à member of the group in a dynamical sense.," In a possible more quantitative analysis, we can be able to extrapolated this kind of relation as a test to verify if a galaxy that has an accordant redshift, but do not seem to be involved in the interactions occurring in the group, is a member of the group in a dynamical sense."
+ We can update our evolutionary sequence envisioned in 7.. where we placed LCC: 79 in the advanced stage of dynamical evolution. WCC 95 in the intermediary stage and LCC. SS in the initial stage.," We can update our evolutionary sequence envisioned in \citet{dar05}, where we placed HCG 79 in the advanced stage of dynamical evolution, HCG 95 in the intermediary stage and HCG 88 in the initial stage."
+ Phe three groups presented here would be placed in the intermediary stage., The three groups presented here would be placed in the intermediary stage.
+" ALL the objects have fraction of IGL between 10 and (with the exception. of LCC 51. when the quintet configuration is considered it has of IGL fraction). the crossing times are in the “niidelle region"" (not as long as"," All the objects have fraction of IGL between 10 and (with the exception of HCG 51, when the quintet configuration is considered it has of IGL fraction), the crossing times are in the “middle region” (not as long as"
+with the optical spectrum showing the presence of a nuclear starburst.,with the optical spectrum showing the presence of a nuclear starburst.
+" Also, such a signature could be produced by a very strong molecular outflow."," Also, such a signature could be produced by a very strong molecular outflow."
+" Finally another, although unlikely, possibility is that in the most central regions, CO is dense enough to be seen in absorption."," Finally another, although unlikely, possibility is that in the most central regions, CO is dense enough to be seen in absorption."
+" Regions located within the galaxy have a large CO width, which is likely due to the rotation of the galaxy as for NC the width of the CO(2-1) line, observed with a smaller beam, is narrower than that of the CO(1-0) line, observed with a beam twice as large."," Regions located within the galaxy have a large CO line--width, which is likely due to the rotation of the galaxy as for NC the width of the CO(2–1) line, observed with a smaller beam, is narrower than that of the CO(1–0) line, observed with a beam twice as large."
+" The line width reduction is minimal in the case of NE, showing the the bulk of the emission must come from the nucleus rather than from the disk of the galaxy."," The line width reduction is minimal in the case of NE, showing the the bulk of the emission must come from the nucleus rather than from the disk of the galaxy."
+ HI emission displays a wider profile due to blending by nearby regions and by the fact that HI is more loosely bound compared to molecular gas., HI emission displays a wider profile due to blending by nearby regions and by the fact that HI is more loosely bound compared to molecular gas.
+ This is particularly visible in C5., This is particularly visible in C5.
+ The NE region shows two features., The NE region shows two features.
+ One corresponds to the double-peaked component mentioned earlier whereas the other one is at a recession velocity around 4800 km s!., One corresponds to the double–peaked component mentioned earlier whereas the other one is at a recession velocity around 4800 km $^{-1}$.
+" This probably corresponds to gas funneled to the nucleus, feeding the current starburst."," This probably corresponds to gas funneled to the nucleus, feeding the current starburst."
+ TDGs are by definition bound objects., TDGs are by definition bound objects.
+ Rotation curves have been clearly detected in a number of objects (??)..," Rotation curves have been clearly detected in a number of objects \citep{bournaud2004a,bournaud2007a}."
+" The observations of Arp 158, unfortunately, do not have enough resolution to measure the rotation curve of C5."," The observations of Arp 158, unfortunately, do not have enough resolution to measure the rotation curve of C5."
+" However, the CO line profile nevertheless provides us with an indication of the mass within the CO beam."," However, the CO line profile nevertheless provides us with an indication of the mass within the CO beam."
+" Indeed the CO traces the densest regions which are most likely gravitationally bound, as opposed to the HI which traces lower density gas."," Indeed the CO traces the densest regions which are most likely gravitationally bound, as opposed to the HI which traces lower density gas."
+ This allows us to compute the virial mass assuming that C5 is kinematically detached from the interacting galaxies., This allows us to compute the virial mass assuming that C5 is kinematically detached from the interacting galaxies.
+" We estimate the mass using the following relation: My;=RAV?/G, where My, is the virial mass, R is half of the CO beam size, AV is the line FWHM and G is the gravitational constant."," We estimate the mass using the following relation: $\mathrm{M_{vir}=R\Delta V^2/G}$, where $\mathrm{M_{vir}}$ is the virial mass, R is half of the CO beam size, $\mathrm{\Delta V}$ is the line FWHM and G is the gravitational constant."
+" For the C5 clump with AV=40 km s! we find My;,=1.3x10? Mo.", For the C5 clump with $\mathrm{\Delta V}=40$ km $^{-1}$ we find $\mathrm{M_{vir}}=1.3\times10^9$ $_\odot$.
+ Such a virial mass is close to the molecular mass of Arp 1055 for instance (?).., Such a virial mass is close to the molecular mass of Arp 105S for instance \citep{braine2001a}.
+ This mass is smaller than the total HI mass in C5 according to ?.., This mass is smaller than the total HI mass in C5 according to \cite{iyer2004a}.
+ This is not surprising as the HI is more extended., This is not surprising as the HI is more extended.
+" We have seen in Table 6 that for C5, 2N(H2)/N(HD=0.25."," We have seen in Table \ref{tab:masses} that for C5, $\mathrm{2N(H_2)/N(HI)=0.25}$."
+" This leads to a minimal mass M(HI+H5He)=4.7x105 Μο within the CO beam, which is about of the virial mass."," This leads to a minimal mass $\mathrm{M(HI+H_2+He)=4.7\times10^8}$ $_\odot$ within the CO beam, which is about of the virial mass."
+" As the HI beam covers nearly 4 times the CO beam area, the HI column density at higher resolution could easily be several times higher or indicative of the presence of dark baryons."," As the HI beam covers nearly 4 times the CO beam area, the HI column density at higher resolution could easily be several times higher or indicative of the presence of dark baryons."
+" Indeed, the stellar mass derived from the SED (Spectral Energy Distribution) fit in Sec."," Indeed, the stellar mass derived from the SED (Spectral Energy Distribution) fit in Sec."
+" 4 is in the range of 0.7x108Μο to 1.3x105Mo, which is too small to explain the discrepancy."," \ref{sec:discussion} is in the range of $\mathrm{0.7\times10^8~M_\odot}$ to $\mathrm{1.3\times10^8~M_\odot}$, which is too small to explain the discrepancy."
+ New high resolution HI observations are required to answer this question., New high resolution HI observations are required to answer this question.
+ What is the exact nature of the clumps in the system?, What is the exact nature of the clumps in the system?
+" We have seen that we are in the presence of a 2-body merger at an intermediate stage in the ? sequence, with the disks already merging but the 2 nuclei still clearly separated."," We have seen that we are in the presence of a 2–body merger at an intermediate stage in the \cite{toomre1977a} sequence, with the disks already merging but the 2 nuclei still clearly separated."
+ Spectroscopic observations have shown that NE is a nuclear starburst in which the efficiency of star formation should be higher than in the rest of the system., Spectroscopic observations have shown that NE is a nuclear starburst in which the efficiency of star formation should be higher than in the rest of the system.
+" These same spectroscopic observations have also shown that NC is an AGN, and as such this region cannot be used to discuss star formation."," These same spectroscopic observations have also shown that NC is an AGN, and as such this region cannot be used to discuss star formation."
+ We examine here the SED of these clumps which provide us with information regarding their stellar content as well as their SFH (Star Formation History)., We examine here the SED of these clumps which provide us with information regarding their stellar content as well as their SFH (Star Formation History).
+observing point (Fieure 10).,observing point (Figure 10).
+ Figure 11 shows the expected correlation between 4 and ty., Figure 11 shows the expected correlation between $q$ and $t_0$.
+ If enough observations are found. it should be possible to determine the 4 value with the accuracy as high as ld0.16 units.," If enough observations are found, it should be possible to determine the $q$ value with the accuracy as high as $ 1 \pm 0.16$ units."
+ The 1995 outburst was already used in our solution., The 1995 outburst was already used in our solution.
+ There was au intensive monitorius canrpaigu of OJ287 (called OJ91) during this outburst season. but uufortunatelv there exists a gap in these observations just at the crucial time (Figure 12).," There was an intensive monitoring campaign of OJ287 (called OJ94) during this outburst season, but unfortunately there exists a gap in these observations just at the crucial time (Figure 12)."
+ It may still be possible that there are measurements somewhere which are not recorded in the OJOL campaign light curve. aud which would be valuable in narrowing down q even from these data.," It may still be possible that there are measurements somewhere which are not recorded in the OJ94 campaign light curve, and which would be valuable in narrowing down $q$ even from these data."
+ The line in Figure 12 is drawn using the standard light curve of Figure 3 ax a template to compare with the 1995 observations., The line in Figure 12 is drawn using the standard light curve of Figure 3 as a template to compare with the 1995 observations.
+ Tn set 11 the value of fj has been shifted down by 3.5 hours. aud in set 12 it is shifted up by the sale amount.," In set 11 the value of $t_0$ has been shifted down by 3.5 hours, and in set 12 it is shifted up by the same amount."
+" The shifts lead to shifting q,4,,4. up or down by ~15%.", The shifts lead to shifting $q_{center}$ up or down by $\sim 15\%$.
+ It should be noted that even a few more measurements of 1995 could narrow down the range of q., It should be noted that even a few more measurements of 1995 could narrow down the range of $q$.
+ Let us now turn our attention to the expected future outbursts ii our binary black hole τους]., Let us now turn our attention to the expected future outbursts in our binary black hole model.
+ We expect three more outbursts during the next two decades occuring in 2015. 2019 and 2022.," We expect three more outbursts during the next two decades occuring in 2015, 2019 and 2022."
+ As we inentioned above. the 2015 outburst should be an easv one to detect. as it is expected iu December of that vear.," As we mentioned above, the 2015 outburst should be an easy one to detect, as it is expected in December of that year."
+ The exact date will iu fact eive us a good spin value., The exact date will in fact give us a good spin value.
+ The dependence on 4 is secondary. and thus it is of no use by itself for the testing of the no-hair ούτος.," The dependence on $q$ is secondary, and thus it is of no use by itself for the testing of the no-hair theorems."
+ The 2019 outburst is sensitive tothe q and with yo=0.275. it should begin at 2019.53175 ifq=1 (Figue 13).," The 2019 outburst is sensitive tothe $q$ and with $\chi = 0.275$, it should begin at 2019.53175 if $q=1$ (Figure 13)."
+ With good timing the 4 value is determined with the accuracy of 1430.16 (Figure 11)., With good timing the $q$ value is determined with the accuracy of $1 \pm 0.16$ (Figure 14).
+ Tf by good luck we will fine the necessary historical data to time both the 1959 and 1971 outbursts in addition to observing the 2019 outburst. we get close to the LO% accuracy in 4 (see the combined sets 1)215-9 in Table 3).," If by good luck we will find the necessary historical data to time both the 1959 and 1971 outbursts in addition to observing the 2019 outburst, we get close to the $10\%$ accuracy in $q$ (see the combined sets 1+2+5-9 in Table 3)."
+ huprovenieuts in uuderstanding the astrophysical processes in OJ287 may then bring the accuracy even below 10:54., Improvements in understanding the astrophysical processes in OJ287 may then bring the accuracy even below $10\%$ .
+ Without the accurate timing iu, Without the accurate timing in
+"T=10(7/0.02pe) MIs. In the radiative transfer program we used a to I»abundance of 2x10."" and CS to II» abundance of 1x10.7.",$T=10(r/0.02\rm{~pc})^{-1.0}$ K. In the radiative transfer program we used a to $_2$abundance of $2\times 10^{-9}$ and CS to $_2$ abundance of $1\times 10^{-9}$.
+ For isotopic ratios. we used [PC/MC] =45 and [S/S] =15.," For isotopic ratios, we used $^{12}$ $^{13}$ C] $= 45$ and $^{32}$ $^{34}$ S] $=15$."
+ mince the fitting is done “by eve’. i€ is difficult to claim that the best fit model is a unique fit to the observed data.," Since the fitting is done “by eye”, it is difficult to claim that the best fit model is a unique fit to the observed data."
+ While some model parameters are only loosely constrained. ihe modeling exercise does set strong constraints on other parameters.," While some model parameters are only loosely constrained, the modeling exercise does set strong constraints on other parameters."
+ For example. the direction and size of the gradient seen in the centroid velocity maps shown in Figure 6 can only be reproduced in the models. by assuming a rotational axis al an ~45? position angle.," For example, the direction and size of the gradient seen in the centroid velocity maps shown in Figure \ref{centall} can only be reproduced in the models, by assuming a rotational axis at an $\sim 45$ position angle."
+ Assuming a north-south or east-west orientation of the rotational axis results in model centroid velocity maps that have gradients west to east and north (o south. not the (o south-east sense seen in Figure 6..," Assuming a north-south or east-west orientation of the rotational axis results in model centroid velocity maps that have gradients west to east and north to south, not the north-west to south-east sense seen in Figure \ref{centall}."
+ The angle that the rotational axis makes to the plane of the sky is also well constrained., The angle that the rotational axis makes to the plane of the sky is also well constrained.
+ Significant departures from (he assumed orientation. viz..," Significant departures from the assumed orientation, viz.,"
+ (he rotational axis is in the plane of the sky. results in bad fits to the ratio of blue-to-red peaks of the lines.," the rotational axis is in the plane of the sky, results in bad fits to the ratio of blue-to-red peaks of the lines."
+" The best fit infall size of 0.015 pe corresponds to an angular diameter of ~20"" at the distance of the Serpens cloud.", The best fit infall size of $0.015$ pc corresponds to an angular diameter of $\sim 20$ at the distance of the Serpens cloud.
+ The estimated core size (as discussed above) isLO”. so (his indicates that the expansion wave might have already escaped the densest regions of the core.," The estimated core size (as discussed above) is, so this indicates that the expansion wave might have already escaped the densest regions of the core."
+ As seen in Figure 3.. other protostellar sources are located quite close to SMMA.," As seen in Figure \ref{intco}, other protostellar sources are located quite close to SMM4."
+ The TSC model is applicable only (o isolated collapse resulting in a single protostar., The TSC model is applicable only to isolated collapse resulting in a single protostar.
+ The values clerived above for the infall parameters of SMMA using the TSC model should (hus be interpreted. wilh some caution., The values derived above for the infall parameters of SMM4 using the TSC model should thus be interpreted with some caution.
+ If we assume that (he constrained infall parameters are not jeavilv affected bv the cluster environment. we can (quantify some other plivsical parameters in (he SMMMA system.," If we assume that the constrained infall parameters are not heavily affected by the cluster environment, we can quantify some other physical parameters in the SMM4 system."
+ From Table 3.. the infall timescale is given by |=μερα—3x10! ves.," From Table \ref{modelfits}, the infall timescale is given by $t =
+r_{inf}/a = 3\times 10^4$ yrs."
+ This number is in good agreement with the lower limit to the age of the outflow svstem derived as LO! vrs in this work (see above). as well as the estimate of 3x104 vears [rom the CO observations of Davisetal.(1999).," This number is in good agreement with the lower limit to the age of the outflow system derived as $\sim
+10^4$ yrs in this work (see above), as well as the estimate of $3\times 10^4$ years from the CO observations of \citet{davis99}."
+. This would indicate that outflow was triggered almost simultaneously with the onset of collapse in SMMA., This would indicate that outflow was triggered almost simultaneously with the onset of collapse in SMM4.
+" The rotational velocitv. ο=7xLO! ! implies that the turnover radius r,=a/Q0.23 pe =150""."," The rotational velocity, $\Omega =
+7\times 10^{-14}$ $^{-1}$ implies that the turnover radius $r_c =
+a/\Omega = 0.23$ pc $= 150\arcsec$."
+ From the sell-similar solution of Shu (1977). a mass accretion rate of 0.975a°/G=2.9x10? + is derived.," From the self-similar solution of Shu (1977), a mass accretion rate of $0.975a^3/G = 2.9\times 10^{-5}$ $^{-1}$ is derived."
+ The present mass of the star-disk system can then be ealeulated by MI giving a mass of ~0.9 M... which falls well within the dust mass limits derived by IHogerheijdeetal(1999) for SMMA.," The present mass of the star-disk system can then be calculated by $\dot{M}t$ giving a mass of $\sim 0.9$ , which falls well within the dust mass limits derived by \citet{hog99} for SMM4."
+ The centrifugal radius (or the accretion disk radius) of the SMALL system at the present epoch is, The centrifugal radius (or the accretion disk radius) of the SMM4 system at the present epoch is
+lieht-curve derivatives.,light-curve derivatives.
+ The cllective transverse velocity obtained. is 20.in units. of⋅ lmsceοἐςn2 where (nm) is. the mean microlens mass.," The effective transverse velocity obtained is in units of $km\,sec^{-1}\,\langle m\rangle^{-\frac{1}{2}}$ where $\langle m\rangle$ is the mean microlens mass."
+" When multiplied by the length. of the monitoring period. this vields the length. of caustic structure sampled 62,5,5:) in units of the ER of the average microlens mass (g,)."," When multiplied by the length of the monitoring period, this yields the length of caustic structure sampled $\eta_{period}$ ) in units of the ER of the average microlens mass $\eta_{o}$ )."
+ This is determined. as a function of source size S., This is determined as a function of source size $S$.
+ Note that. Ενω(9). is free of assumptions about time-scale and microlens mass., Note that $\eta_{period}(S)$ is free of assumptions about time-scale and microlens mass.
+ We have compute the probability pyCppericd|5) for source sizes 5— 0.0015625 0.003125. 0.00625 0.0125. 0.025. 0.05. 0.1. 0.2. 0.4. OLS anc l.65g..," We have computed the probability $p_{\eta}(\eta_{period}|S)$ for source sizes $S=$ 0.0015625 0.003125, 0.00625 0.0125, 0.025, 0.05, 0.1, 0.2, 0.4, 0.8 and $\eta_o$."
+ For S«θ and SzOye. Ριςμον). was computed from 5000 ancl 1000 mock sets of monitoring data respectively at cach of 50 assumed sampling lengths η.," For $S<0.1\eta_o$ and $S\ge0.1\eta_o$, $p_{\eta}(\eta_{period}|S)$ was computed from 5000 and 1000 mock sets of monitoring data respectively at each of 50 assumed sampling lengths $\eta$."
+ The cumulative distributions are shown in the upper plot of figure 2— (thicker lines indicate larger sources)., The cumulative distributions are shown in the upper plot of figure \ref{sample_length} (thicker lines indicate larger sources).
+" For Sx0.0554, (solid lines). the measurement of sample length is approximately independent of source size (this independence is discussed in WWT90)."," For $S\leq0.05\eta_{o}$ (solid lines), the measurement of sample length is approximately independent of source size (this independence is discussed in WWT99)."
+ LIowever. the measured. sample length increases," However, the measured sample length increases"
+sinultaueouslv dn a spectral line wmode with polarization products VN and VY in 1021 channels across a 1. MITZ total bandwidth.,simultaneously in a spectral line mode with polarization products $XX$ and $YY$ in 1024 channels across a 4 MHz total bandwidth.
+ The conutimuun data are centered on v=13al MIIz. whereas the spectral line data are centered on ;=1120.0 MIIz and have channel separation of Av=O082kins+ (3.9 ΚΠΣ].," The continuum data are centered on $\nu = 1384$ MHz, whereas the spectral line data are centered on $\nu= 1420.0$ MHz and have channel separation of $\Delta v = 0.82~{\rm km~s^{-1}}$ $3.9$ kHz)."
+ Data editing. calibration. aud imaging of the ATCA data are carried out in the data reduction package (Sault&τοσα2000).," Data editing, calibration, and imaging of the ATCA data are carried out in the data reduction package \citep{sault00}."
+. The source PINS D19031-638 is used for flux deusitv aud baudpass calibration and was observed. once per observing session., The source PKS B1934-638 is used for flux density and bandpass calibration and was observed once per observing session.
+ À flux density of LL9OL Jv at 1120 ΑΠΣ is aà9uined (Reynolds1991)., A flux density of 14.94 Jy at 1420 MHz is assumed \citep{reynolds94}.
+. PINS D1931-6358 has no detectable liuear polarization and can therefore be used to solve for the polarization leakages., PKS B1934-638 has no detectable linear polarization and can therefore be used to solve for the polarization leakages.
+ Polarization leakages are also caleulated for the sources MIRC B1613-586 and AIRC Bl131-18 which were observed approximately ouce every hour for calibration of the tine-variation im complex auteuna gaius., Polarization leakages are also calculated for the sources MRC B1613-586 and MRC B1431-48 which were observed approximately once every hour for calibration of the time-variation in complex antenna gains.
+ The individual poiutiues were luearly combined aud nuaged using a standard erid-aud-FFT scheme with super-uniforii weighting., The individual pointings were linearly combined and imaged using a standard grid-and-FFT scheme with super-uniform weighting.
+ Like uniforiii weighting. supeor-uufor weighting minimizes sidelobe levels to improve the dynamic range aud seusitivitv to extended structure.," Like uniform weighting, super-uniform weighting minimizes sidelobe levels to improve the dynamic range and sensitivity to extended structure."
+ Uniform weighting reduces to natural weielting. however. if the field of view is πιο larger than the primary beam. as is the case for large mosaics.," Uniform weighting reduces to natural weighting, however, if the field of view is much larger than the primary beam, as is the case for large mosaics."
+ Super-muiform weighting overcomes this limitation by decoupling the weighting ron the field size (Sault&Isillecu2000)., Super-uniform weighting overcomes this limitation by decoupling the weighting from the field size \citep{sault00}.
+. Th this mamuer super-unitorm weighting attempts to minimize sidelobe contributions from strong sources over a region smaller hau the full field being imaged aud therefore is typically nore successful than uniformi weighting on large mosaics (Saultetal.1996)., In this manner super-uniform weighting attempts to minimize sidelobe contributions from strong sources over a region smaller than the full field being imaged and therefore is typically more successful than uniform weighting on large mosaics \citep{sault96}.
+. The specifies of the calibration aud imagine of ATCA xolarization data are described in (Gacusleretal.2000a).. rere we discuss ouly the Stokes I coutimmun aud ddata.," The specifics of the calibration and imaging of ATCA polarization data are described in \citep{gaensler00a}, here we discuss only the Stokes I continuum and data."
+ Two ddata cubes were produced., Two data cubes were produced.
+ Iu one cube the continu enission was subtracted from the οσα using the task uvlin., In one cube the continuum emission was subtracted from the data using the task uvlin.
+ The second cube. for use iu absorption studies. contains contiuuun cussion.," The second cube, for use in absorption studies, contains continuum emission."
+ The resultaut svuthesized bean for both cubes aud the continuuni image is 12179ς10775 (a« 8).," The resultant synthesized beam for both cubes and the continuum image is $124\farcs9 \times 107\farcs5$ $\alpha \times
+\delta$ )."
+ For the data presented here the 6 km baselines of the ATCA were not used., For the data presented here the 6 km baselines of the ATCA were not used.
+ However. the long baseliues are retained for some absorption studies. such as those presented iu Dickeyctal. (2000).," However, the long baselines are retained for some absorption studies, such as those presented in \citet{dickey00}."
+. Joiut decouvolution was performed on the full linear mosaics using a niaxinmnà entropy aleorithiu inplemenuted m task mosimem (Saultetal.1996)., Joint deconvolution was performed on the full linear mosaics using a maximum entropy algorithm implemented in task mosmem \citep{sault96}.
+. The imethod of joint decouvolution is very effective for maximizing thewe coverage attainable iu mosaiced observations (Saultctal.1996)., The method of joint deconvolution is very effective for maximizing the coverage attainable in mosaiced observations \citep{sault96}.
+". Despite deconvolution. sole sidelobes are visible around strong poiut sources m the continmun and Mages,"," Despite deconvolution, some sidelobes are visible around strong point sources in the continuum and images."
+ Though 1uosaicius allows us to recover angular scales larger than nonual interferometric observations bv reducing the effective shortest projected. baseline. we are nouctheless limited im practice to angular scales smaller than 0=Af(dD2). where d=30.6 1 is the shortest baseline of the ATCA. and D=22 ii is the diameter of a single auteuna.," Though mosaicing allows us to recover angular scales larger than normal interferometric observations by reducing the effective shortest projected baseline, we are nonetheless limited in practice to angular scales smaller than $\theta = \lambda / (d -D/2)$, where $d=30.6$ m is the shortest baseline of the ATCA, and $D= 22$ m is the diameter of a single antenna."
+" This limits the ATCA data to angular scales sunaller than 02236""."," This limits the ATCA data to angular scales smaller than $\theta
+\approx 36\arcmin$."
+ In order to recover information on leger size scales. the ATCA mosaic niust be combined with sinele dish data from the Parkes Radio Telescope.," In order to recover information on larger size scales, the ATCA mosaic must be combined with single dish data from the Parkes Radio Telescope."
+ The Parkes Radio Telescope is a Glin antenna situated near Parkes NSW. Australia.," The Parkes Radio Telescope is a 64 m antenna situated near Parkes NSW, Australia."
+ It has a thirteen lea. A2l-cin receiver package at prine focus (Stavelev-Siuithetal.1996).," It has a thirteen beam, $\lambda$ 21-cm receiver package at prime focus \citep{staveley-smith96}."
+. The SCPS Test Region was observed on 1998 December 15-16., The SGPS Test Region was observed on 1998 December 15-16.
+ The Parkes survey Was subsequentIv expanded to 5=410° for more complete coverage of large scale structures durine additional observiug sessjons: 1998 June 18-21. 1999 September 18-27. aud 2000 March 10-15.," The Parkes survey was subsequently expanded to $b=\pm
+10\arcdeg$ for more complete coverage of large scale structures during additional observing sessions: 1998 June 18-21, 1999 September 18-27, and 2000 March 10-15."
+ Observations were made by the process of mapping “ou-the-fw with the immer seven beams of the multibeam svstem., Observations were made by the process of mapping “on-the-fly” with the inner seven beams of the multibeam system.
+ Data were xccorded iu 5 s samples. while scanning through three deerees in Galactic latitude.," Data were recorded in 5 s samples, while scanning through three degrees in Galactic latitude."
+ The data were taken in frequeney switching mode using the narrowband back-end (Πάνοςctal.1998).. with a total bandwidth of δ MITz across 2018 chauncls.," The data were taken in frequency switching mode using the narrowband back-end \citep{haynes98}, with a total bandwidth of 8 MHz across 2048 channels."
+ Though the Parkes data are frequency switched. total power information for cach sample is retained.," Though the Parkes data are frequency switched, total power information for each sample is retained."
+ Each sample was divided bv the previous frecuency switched sample and the residual bandpass shape fitted with a series of Fourier compoucuts., Each sample was divided by the previous frequency switched sample and the residual bandpass shape fitted with a series of Fourier components.
+ The spectra were then multipled by the mean of the reference signa over the spectiun to reconstruct the continu einissiou with a flat bascline., The spectra were then multiplied by the mean of the reference signal over the spectrum to reconstruct the continuum emission with a flat baseline.
+ Absolute brightness temperature calibration of the line data was performed frou observations of the LAU standard reeious S6 and 89 (WiljamsL973)., Absolute brightness temperature calibration of the line data was performed from observations of the IAU standard regions S6 and S9 \citep{williams73}.
+. A detailed escription of the observing stratecv. calibration. and nuages procedures ds found 1ji MeChure-Cuffiths.," A detailed description of the observing strategy, calibration, and imaging procedures is found in \citet{mcgriff00b}."
+et Off-line chaunels were used for coutinuuin subtraction and to produce the coutiuuuni nuage use the AIPS task IMLIN., Off-line channels were used for continuum subtraction and to produce the continuum image using the AIPS task IMLIN.
+ The Parkes data ou the SGPS Test Reegiou have au angular resohitico of ~15’., The Parkes data on the SGPS Test Region have an angular resolution of $\sim 15\arcmin$.
+ The final. calibrated data have a bandwidth of 1 MITz with 1021 channels. such that the velocity resolution of 0.82 uunatches the ATCA data.," The final, calibrated data have a bandwidth of 4 MHz with 1024 channels, such that the velocity resolution of $0.82$ matches the ATCA data."
+ It should be noted that the diita presented here have not been corrected for the effects of stray radiation., It should be noted that the data presented here have not been corrected for the effects of stray radiation.
+ Strav radiation leakage from bright eenission through the back sidelolvos of a single dish beam introduces baseline errors which are typically between 0.5 I& and 2 K (Ixalberla.Mebold&Reich1980)., Stray radiation leakage from bright emission through the back sidelobes of a single dish beam introduces baseline errors which are typically between $0.5$ K and $2$ K \citep{kalberla80}.
+. When compared to low latitude Calactic spectra. this is a small percentage. but it does nonetheless limut the seusitivitv and accurate representation of exteuded features in the Parkes data.," When compared to low latitude Galactic spectra, this is a small percentage, but it does nonetheless limit the sensitivity and accurate representation of extended features in the Parkes data."
+ The data for the eutire survey will have a first order stray radiation correction applied., The data for the entire survey will have a first order stray radiation correction applied.
+ A complete correction requires a low resolution survey of the eutire sky with a known antenna pattern in order to reconstruct the strav radiation at every point on the sky. at any azimuth auc elevation. aud at any time of the vear.," A complete correction requires a low resolution survey of the entire sky with a known antenna pattern in order to reconstruct the stray radiation at every point on the sky, at any azimuth and elevation, and at any time of the year."
+ Such a task is bevond the scope of this project., Such a task is beyond the scope of this project.
+ A first order correction. however. can be done by re-obseving the sirvey region at different times of the vear iux comparing the spectra.," A first order correction, however, can be done by re-observing the survey region at different times of the year and comparing the spectra."
+ The velocity shifts caused by the Ταντς motion around the Sun result iu excess cussion at different velocities., The velocity shifts caused by the Earth's motion around the Sun result in excess emission at different velocities.
+" We have. therefore. re-observed the full survey region four times at three mouth iutervals around the vear aud we will compare cach προςπα,"," We have, therefore, re-observed the full survey region four times at three month intervals around the year and we will compare each spectrum."
+ The niununni value at cach velocity will be a reasonable τιo»per Tait to the stray radiation corrected profile. eood ο within 0.5 I< of Tp.," The minimum value at each velocity will be a reasonable upper limit to the stray radiation corrected profile, good to within $\sim 0.5$ K of $T_B$ ."
+"An initial fit was performed in which the eecentricity was allowed to float, yielding e = 0.036 + 0.021.","An initial fit was performed in which the eccentricity was allowed to float, yielding $e$ = 0.036 $\pm$ 0.021."
+" From the Lucy-Sweeney test (2),, we find it is only significant at the 1o level, hence a circular orbit was adopted Que - 18.3, "" = 14.0 with 13 RVs and 5 [ree parameters: A, sCORALLE. sopluu- €cosc and € sinc)."," From the Lucy-Sweeney test \citep{Lucy71}, we find it is only significant at the $\sigma$ level, hence a circular orbit was adopted $\chi^{2}_{\rm circ}$ = 18.3, $\chi^{2}_{\rm ecc}$ = 14.0 with 13 RVs and 5 free parameters: $K$, $\gamma_{\rm CORALIE}$, $\gamma_{\rm SOPHIE}$, $e\cos \omega$ and $e \sin \omega$ )."
+ We determined an upper limit for the eccentricily as eupper=€+300.075., We determined an upper limit for the eccentricity as $e_{\rm upper} = e + 3\sigma = 0.078$.
+" We also fitted a linear trend in the RVs to search for a third body in the system, but found that the decrease in 47 did not warrant the extra [ree parameter."," We also fitted a linear trend in the RVs to search for a third body in the system, but found that the decrease in $\chi^{2}$ did not warrant the extra free parameter."
+" From the F-test, we found the significance of the trend was less than Io (42218. V"" = 15.4 with 13 RVs and 4 free parameters: A, copi. sop and linear trend 7)."," From the F-test, we found the significance of the trend was less than $\sigma$ $\chi^{2}$ = 18.3, $\chi^{2}_{\rm trend}$ = 15.4 with 13 RVs and 4 free parameters: $K$, $\gamma_{\rm CORALIE}$, $\gamma_{\rm SOPHIE}$ and linear trend $\dot{\gamma}$ )."
+ We performed the final MCMC fit with a chain length of 20.000 points and the resulung best-fit parameters and uncertainties are shown in Table 4..," We performed the final MCMC fit with a chain length of 20,000 points and the resulting best-fit parameters and uncertainties are shown in Table \ref{pparams}."
+" WASP-37b has P? = 3.577469 + 0.000011 d, == 1.804 0.17Mj. -- 0.070.06 aand EN-- 0.12quzpj«"," WASP-37b has $P$ = 3.577469 $\pm$ 0.000011 d, = 1.80 $\pm$ 0.17, = $^{+ 0.07}_{- 0.06}$ and = $^{+ 0.12}_{- 0.15}$."
+ We report the discovery of a new transiting hot Jupiter. WASP-37b.," We report the discovery of a new transiting hot Jupiter, WASP-37b."
+ The planet has a radius comparable to Jupiter == n Rj)) but is more massive == 1.80 + 0.17 M))) and hence is very similar to WASP-5 (22).," The planet has a radius comparable to Jupiter = $^{+ 0.07}_{- 0.06}$ ) but is more massive = 1.80 $\pm$ 0.17 ) and hence is very similar to WASP-5 \citep{Anderson08, Southworth09}."
+" The surlace gravity of WASP-37b is high compared to other planets with similar orbital periods gj, = A gp = n 7, == 0.12I p4).NN which⋅ hes⋅ significantly⋅UN above the period-gravily⋅ correlation⋅ proposed by ‘?.."," The surface gravity of WASP-37b is high compared to other planets with similar orbital periods $_{\rm p}$ = $^{+ 0.03}_{- 0.04}$, $g_{\rm p}$ = $^{+2.0}_{-2.8}$ $^{-2}$, = $^{+ 0.12}_{- 0.15}$ ), which lies significantly above the period-gravity correlation proposed by \citet{Southworth10}."
+" By comparison with the theoretical models of ?. and 2,, we find that WASP-37b has an inflated radius."," By comparison with the theoretical models of \citet*{Fortney07} and \citet*{Baraffe08}, we find that WASP-37b has an inflated radius."
+" The models do not cover the age range of WASP-37, however given that planetary radii are thought to decrease with age, the expected radius of WASP-37b will be smaller than those given for the oldest model."," The models do not cover the age range of WASP-37, however given that planetary radii are thought to decrease with age, the expected radius of WASP-37b will be smaller than those given for the oldest model."
+" For an equivalent semi-major axis with solar insolation of 0.445 AU, the measured mass and radius do not suggest the presence ol a core of heavy elements for models older than 1 Gyr, and this 15 consistent with the correlation between core mass and metallicity (2?) given the low stellar metallicity."," For an equivalent semi-major axis with solar insolation of 0.445 AU, the measured mass and radius do not suggest the presence of a core of heavy elements for models older than 1 Gyr, and this is consistent with the correlation between core mass and metallicity \citep{Guillot06, Burrows07} given the low stellar metallicity."
+ We detect no significant orbital eccentricity and find an upper limit of 0.098., We detect no significant orbital eccentricity and find an upper limit of 0.098.
+ Nor do we find evidence of any long-term trend caused by a third body., Nor do we find evidence of any long-term trend caused by a third body.
+ Our radial velocity observations span ~ 70 days and continued monitoring would reveal the presence of any longer-term trends and further constrain the eccentricily., Our radial velocity observations span $\sim$ 70 days and continued monitoring would reveal the presence of any longer-term trends and further constrain the eccentricity.
+" With a stellar effective temperature of 5800 + 150 K and orbital period of 3.577469 + 0.000011 d, the planet is⋅ predicted⋅ to have a equilibrium⋅⋅⋅ temperature of- OF,15 K (assuming⋅ zero albedo and [ull- redistribution of the incident radiation)."," With a stellar effective temperature of 5800 $\pm$ 150 K and orbital period of 3.577469 $\pm$ 0.000011 d, the planet is predicted to have a equilibrium temperature of $^{+25}_{-15}$ K (assuming zero albedo and full redistribution of the incident radiation)."
+ Secondary eclipse measurements of WASP-37 may be a way of constraining the influence of metallicity on the presence or not of atmospheric temperature inversions., Secondary eclipse measurements of WASP-37 may be a way of constraining the influence of metallicity on the presence or not of atmospheric temperature inversions.
+" The relaüvely large uncertainties in the stellar parameters are caused by the low S/N of the combined spectrum used for spectral analysis, due to the faintness of the target ην = 12.7)."," The relatively large uncertainties in the stellar parameters are caused by the low S/N of the combined spectrum used for spectral analysis, due to the faintness of the target $m_{\rm v}$ = 12.7)."
+ Further spectra would, Further spectra would
+The point of this section is that Eq.,The point of this section is that Eq.
+ 5 provides a simple and general explanation lor the observed variation of ρε) with P5»; (Eq., 5 provides a simple and general explanation for the observed variation of $\langle E_{peak} \rangle$ with $P_{256}$ (Eq.
+ 6)., 6).
+ In words. the brighter bursts will have some combination of higher luminosities (with a higher LP) and be closer (with a smaller 1+2) than fainter bursts. so that the jets blue shift and cosmological red shift will work togetherto give a higher ρω for the brighter bursts.," In words, the brighter bursts will have some combination of higher luminosities (with a higher $\Gamma$ ) and be closer (with a smaller $1+z$ ) than fainter bursts, so that the jet's blue shift and cosmological red shift will work togetherto give a higher $E_{peak}$ for the brighter bursts."
+ The values of M. and N can tell us about the physics of gammia rav bursts., The values of $M$ and $N$ can tell us about the physics of gamma ray bursts.
+ Observations can place reliable constraints on both AZ and AN., Observations can place reliable constraints on both $M$ and $N$.
+ Sections 2 and 3 have already provided limits. ancl (wo more observational constraints will now be presented.," Sections 2 and 3 have already provided limits, and two more observational constraints will now be presented."
+" For the 112 DATSE bursts with known L and = (based on the measured lag and variabilitv) reported in Schaefer.Deng.&Band(2001).. the E,,,5 value al the time of peak flux is known for 84 events. ("," For the 112 BATSE bursts with known $L$ and $z$ (based on the measured lag and variability) reported in \citet{sdb01}, the $E_{peak}$ value at the time of peak flux is known for 84 events. ("
+The reason lor why the 28 events do not have a measured Peak 15 nol known due to the untimely death of R. Mallozzi.),The reason for why the 28 events do not have a measured $E_{peak}$ is not known due to the untimely death of R. Mallozzi.)
+ Eq., Eq.
+ 5 shows that. Εμ.4-2) should varv with L as a power law with index (AJ+1)/N., 5 shows that $E_{peak}(1+z)$ should vary with $L$ as a power law with index $(M+1)/N$.
+ As shown in Figure 3. the logarithans of Γρτε2) and L are indeed accurately related linearly.," As shown in Figure 3, the logarithms of $E_{peak}(1+z)$ and $L$ are indeed accurately related linearly."
+ A chi-square fit gives CMVN=0.36£0.03., A chi-square fit gives $(M+1)/N=0.36 \pm 0.03$.
+ This same method can be applied to 20 bursts which have the Iuminositiesbased on optically measured. red shifts., This same method can be applied to 20 bursts which have the luminositiesbased on optically measured red shifts.
+ Figure 3 shows these bursts. with the data taken from the literature (primarily Schaefer(2002) anc Amatiοἱal. (2002))).," Figure 3 shows these bursts, with the data taken from the literature (primarily \citet{sch02} and \citet{ama02}) )."
+ This sample also shows a power law dependence. with a highly significant correlation (7=0.90. for a chance probability of 3x10 7).," This sample also shows a power law dependence, with a highly significant correlation $r=0.90$, for a chance probability of $3
+\times 10^{-8}$ )."
+ The slope is (AZ+D)/N. which is fitted to be 0.38--0.11.," The slope is $(M+1)/N$, which is fitted to be $0.38 \pm
+0.11$."
+ The variabilitv/Iuminositwv relation (Fenimore&Ramirez-Ruiz2000) can be used to constrain N., The variability/luminosity relation \citep{frr00} can be used to constrain $N$.
+ Variability (V) is a measure of how ‘spikew a burst lightcurve is. which empirically varies as (he inverse of (he rise limes in the light curves.," Variability $V$ ) is a measure of how 'spikey' a burst lightcurve is, which empirically varies as the inverse of the rise times in the light curves."
+" The rise (mes will be limited by the geometric delavs which scale as DP.7. κο V.xΓή, Scha"," The rise times will be limited by the geometric delays which scale as $\Gamma^{-2}$, so $V \propto \Gamma^2$."
+efer(2002) has shown from 9 bursts with measuredred shifts that Lxο , \citet{sch02} has shown from 9 bursts with measuredred shifts that $L \propto V^{1.57 \pm 0.17}$.
+Thus; we have LoxThs! which is to sav NV=3.1440.34.," Thus, we have $L \propto \Gamma^{3.14 \pm 0.34}$, which is to say $N=3.14
+\pm 0.34$."
+ ] now have four observational restrictions and two theoretical restrictions on NV and M. (, I now have four observational restrictions and two theoretical restrictions on $N$ and $M$. (
+1) The narrowness of the observed [ρω distribution requires CV+1)/N.<0.45. (,1) The narrowness of the observed $E_{peak}$ distribution requires $(M+1)/N<0.45$. (
+"2) The variation of (E,ok) Wilh P535 is easily explained if CM+1)/N=0404 0.14. (",2) The variation of $\langle E_{peak} \rangle$ with $P_{256}$ is easily explained if $(M+1)/N=0.40 \pm 0.14$ . (
+3) The variation ol ροκ2) with L for two independent set of bursts shows that CU--1)/.N equals 0.36dE0.03 and 0.38 £0.12. (,"3) The variation of $E_{peak}(1+z)$ with L for two independent set of bursts shows that $(M+1)/N$ equals $0.36 \pm
+0.03$ and $0.38 \pm 0.12$ . ("
+4) The variabilityluminosity relation gives NV=3.144 0.34. (,4) The variability/luminosity relation gives $N=3.14 \pm 0.34$ . (
+5) Ordinary relativistic effects for any internal shock model implies that V= 3. (,5) Ordinary relativistic effects for any internal shock model implies that $N \ga 3$ . (
+6) The,6) The
+acceleration of fast chareed particles at a shock which wopagating normal to the average upstream magnetic Seld.,acceleration of fast charged particles at a shock which propagating normal to the average upstream magnetic field.
+ We next consider the case where the shock xopagation direction isrot normal to the magnetic field., We next consider the case where the shock propagation direction is normal to the magnetic field.
+ Clearly. if the varviug direction of the upstream uaenetie field is such that at some places the local anele of the magnetic field relative to the average field direction exceeds the augle of the average maguctic field o the shock plane. we will have situations similar to that discussed in the previous sections.," Clearly, if the varying direction of the upstream magnetic field is such that at some places the local angle of the magnetic field relative to the average field direction exceeds the angle of the average magnetic field to the shock plane, we will have situations similar to that discussed in the previous sections."
+ There will be places where the connection poiuts of the maguetic field to the shock move further apart or closer together., There will be places where the connection points of the magnetic field to the shock move further apart or closer together.
+ Heuce we expect the same plivsies can be applicable., Hence we expect the same physics can be applicable.
+ An example is eiven in Figure 7., An example is given in Figure 7.
+" In this case the ratio of 9D/By is taken to be 0.5. the averaged shock normal augle Op,=TO’."," In this case the ratio of $\delta B/B_0$ is taken to be $0.5$, the averaged shock normal angle $\theta_{Bn} = 70^\circ$."
+ It cau be seen from this plot that the connection points can still move toward cach other in sole reeious., It can be seen from this plot that the connection points can still move toward each other in some regions.
+ Fieure 8 shows the density contours of accelerated particles the same as Figure 2. but for the case of the oblique shock.," Figure 8 shows the density contours of accelerated particles the same as Figure 2, but for the case of the oblique shock."
+ We fiud that in this case the process we discussed in the last section is still persistent. even for a oblique shock aud relative smaller 6BBy.," We find that in this case the process we discussed in the last section is still persistent, even for a oblique shock and relative smaller $\delta B/B_0$."
+" The ""hot spot” forms correspond to the convereiug /magnetic connection points aud particle acceleration is suppressed in the region where connection points separate from each other."," The “hot spot"" forms correspond to the converging magnetic connection points and particle acceleration is suppressed in the region where connection points separate from each other."
+" We may couchicde that. for a shock which is oblique. if some magnetic field Hues can intersect the shock 1nultiple times. we have ""hot spots” of accelerated particles forms where the connection poiuts convereiug together."," We may conclude that, for a shock which is oblique, if some magnetic field lines can intersect the shock multiple times, we have “hot spots"" of accelerated particles forms where the connection points converging together."
+ The acceleration of charged particles in shock waves is oue of the most miportaut unsolved problems iu space phivsies and astrophysics., The acceleration of charged particles in shock waves is one of the most important unsolved problems in space physics and astrophysics.
+ The charged particle transport in turbulent magnetic field and acceleration m shock region are two inseparable problems., The charged particle transport in turbulent magnetic field and acceleration in shock region are two inseparable problems.
+ Iu this paper we illustrate the effect of a large-scale sinusoidal magnetic field variation., In this paper we illustrate the effect of a large-scale sinusoidal magnetic field variation.
+ This siuple model allows a detailed exalination of the plivsical effects;, This simple model allows a detailed examination of the physical effects.
+" As the maguctic ficld lines pass through the shock. the counection poiuts vetween field lines on the shock surface will move accordingly,"," As the magnetic field lines pass through the shock, the connection points between field lines on the shock surface will move accordingly."
+" We fud that the region where counectiou »oiuts approaching each other will trap aud prefercutially accelerate particles to ligh cuereies aud form “hot spots” along the shock surface. somewhat in analogy to he ""hot spots” postulated by Jokipi&Iota(2008).."," We find that the region where connection points approaching each other will trap and preferentially accelerate particles to high energies and form “hot spots"" along the shock surface, somewhat in analogy to the “hot spots"" postulated by \citet{Jokipii2008AIP}."
+ The shock acceleration will be suppressed at places where the connection points move apart cach other., The shock acceleration will be suppressed at places where the connection points move apart each other.
+" Some of the particles injected im those regions will transport to the ""hot spots” aud ect further accelerated."," Some of the particles injected in those regions will transport to the “hot spots"" and get further accelerated."
+ The resulting. distribution function is lighly spatial dependent at the energies we studied. which could eive a possible explanation to the observation. of anomalous cosmüc rays.," The resulting distribution function is highly spatial dependent at the energies we studied, which could give a possible explanation to the observation of anomalous cosmic rays."
+ Although we have discussed a simplified. illustrative model. the resulting spectra aud radial distributions show qualitative similarity with the observations.," Although we have discussed a simplified, illustrative model, the resulting spectra and radial distributions show qualitative similarity with the observations."
+ Thus. the intensities do not in general. peak at the the shock aud the energy spectra are not power laws.," Thus, the intensities do not in general, peak at the the shock and the energy spectra are not power laws."
+ We show this process is robust even im the case of oblique shocks with relatively sinall maguetic field variations., We show this process is robust even in the case of oblique shocks with relatively small magnetic field variations.
+ Large scale magnetic feld variation. which could be due to magnetic structures like maguctic clouds. or the ubiquitous large scale field line random walk. will strongly modify the simple planar shock solution.," Large scale magnetic field variation, which could be due to magnetic structures like magnetic clouds, or the ubiquitous large scale field line random walk, will strongly modify the simple planar shock solution."
+ This effect could work in a nuuber of situations for laree scale shock acceleration imceludius maguetic variations. for example. the solar wind termination shock and supernova blast waves.," This effect could work in a number of situations for large scale shock acceleration including magnetic variations, for example, the solar wind termination shock and supernova blast waves."
+ F. C. would like to thauk Joe Cdacalone for the tutorial and discussion of stochastic iuteeration method., F. G. would like to thank Joe Giacalone for the tutorial and discussion of stochastic integration method.
+ F. €. also thank Clunsheng Pei aud Erica. MceEvox. for valuable discussion on stochastic differential equatious., F. G. also thank Chunsheng Pei and Erica McEvoy for valuable discussion on stochastic differential equations.
+ We acknowledge partial support fron NASA ποτ erants NNNOSATIS5C and NNNOSAQTLC. Iu a 2-D system we considered. the Parker transport equation (1) can be re-written as. The equivalent stochastic differcutial equatiou is. Where rp. ro. and rs are different sets ο παπάο uimnmnibers which⋅ satisfvur <r;>=O aud «or?»>=:L.," We acknowledge partial support from NASA under grants NNX08AH55G and NNX08AQ14G. In a 2-D system we considered, the Parker transport equation $(1)$ can be re-written as, The equivalent stochastic differential equation is, Where $r_1$, $r_2$, and $r_3$ are different sets of random numbers which satisfy $<r_i> = 0$ and $<r_i^2> = 1$."
+" In order to study. diffusive shock acceleration. we approximate the shock laver as a sharp variation (=(0C,|C5)/2ΜΗο witha thickness tf 1iucli smaller conrpared with the characteristic leneth of diffusion acceleration &,,4/04."," In order to study diffusive shock acceleration, we approximate the shock layer as a sharp variation $U_x =
+(U_1+U_2)/2 - (U_1-U_2)tanh(x/th)/2$ with a thickness $th$ much smaller compared with the characteristic length of diffusion acceleration $\kappa_{xx1}/U_1$."
+ At the same time. we have to male sure the time step Aft is small enough to resolve the motion in the shock laver.," At the same time, we have to make sure the time step $\Delta t$ is small enough to resolve the motion in the shock layer."
+9. show a flattening toward fainter magnitudes.,"9, show a flattening toward fainter magnitudes."
+ The local LF of Londish (2002) lor ivpe 1 AGN reveals a much flatter slope at the faint end as compared to the bright end and can be well lit with a two power-law function., The local LF of Londish (2002) for type 1 AGN reveals a much flatter slope at the faint end as compared to the bright end and can be well fit with a two power-law function.
+ Our increasing numbers at the faint end may be an evolutionary effect or might be attrvibutecl to the inclusion of a greater fraction ol obsetrecd) AGN., Our increasing numbers at the faint end may be an evolutionary effect or might be attributed to the inclusion of a greater fraction of obscured AGN.
+ More data are needed as well as better classification into AGN tvpes to improve (he statistics and verily this trend., More data are needed as well as better classification into AGN types to improve the statistics and verify this trend.
+ We have investigated. variability within the IIubble Deep Field North with the aim of identifving low-luminosity AGN in the nuclei of field galaxies., We have investigated variability within the Hubble Deep Field North with the aim of identifying low-luminosity AGN in the nuclei of field galaxies.
+ We compare photometry of the nuclei of galaxies in the initial V-band image obtained in 1995 and a second image obtained in 2000 and [ind evidence for significant variability in 16 galaxies., We compare photometry of the nuclei of galaxies in the initial V-band image obtained in 1995 and a second image obtained in 2000 and find evidence for significant variability in 16 galaxies.
+" C'orrecting for possible spurious detections and incompleteness results in a total of ~18 sources which make up ol galaxy nuclei (o. V,,,.221.5.", Correcting for possible spurious detections and incompleteness results in a total of $\sim$ 18 sources which make up of galaxy nuclei to $_{nuc}$ =27.5.
+ The galaxies hosting variable nuclei range in redshift [rom 0.00«z«1.8 with most ling between 0.4«z«1.1., The galaxies hosting variable nuclei range in redshift from $<$ $<$ 1.8 with most lying between $<$ $<$ 1.1.
+ Nuclear colors of the variable nuclei are consistent. with those of non-variable nuclei., Nuclear colors of the variable nuclei are consistent with those of non-variable nuclei.
+ Seven of our variable sources are coincident (within ~ 1.2”) with ταν sources detected in the 2Ms Chandra exposure (Alexander 2003)., Seven of our variable sources are coincident (within $\sim$ $\arcsec$ ) with X-ray sources detected in the 2Ms Chandra exposure (Alexander 2003).
+ We find that of the 2Ms Chandra detections in (he original IDE. are significant. variables., We find that of the 2Ms Chandra detections in the original HDF are significant variables.
+ The optical variability observed in these X-ray sources provides important evidence to confirm the presence of an AGN and could potentially help discriminate between low-Iuminosity X-rav. emitting AGN and other X-ray emitting phenomena such as supernova remnants or X-ray binaries., The optical variability observed in these X-ray sources provides important evidence to confirm the presence of an AGN and could potentially help discriminate between low-luminosity X-ray emitting AGN and other X-ray emitting phenomena such as supernova remnants or X-ray binaries.
+ In addition. of our sources are associated with mid-IR detections at 15jmi and are detected at 1.4GGIHz.," In addition, of our sources are associated with mid-IR detections at $\micron$ and are detected at 1.4GHz."
+ Optical spectra are currently available for 13 of the 16 variable sources., Optical spectra are currently available for 13 of the 16 variable sources.
+ Object2-251.0 is a broad-line AGN., Object is a broad-line AGN.
+ Two others show some evidence of harboring an AGN., Two others show some evidence of harboring an AGN.
+ Object4-970.1. a low redshift spiral. has emission line ratios indicative of a LINER/Sevlert 2 and4-254.0 reveals weak NeV emission.," Object, a low redshift spiral, has emission line ratios indicative of a LINER/Seyfert 2 and reveals weak NeV emission."
+" Based on the V/V,,,, test. there is evidence for evolution within our sample of variable nuclei in the sense that more variables are [found at higher reclshilts."," Based on the $_{max}$ test, there is evidence for evolution within our sample of variable nuclei in the sense that more variables are found at higher redshifts."
+ We estimate the huninosity function of variable AGN in the redshift range 0.4«z«1.1., We estimate the luminosity function of variable AGN in the redshift range $<$ $<$ 1.1.
+ This LF is a lower-limit since our variability selection technique. employing only 2 epochs for detection. could miss —20 of variable nuclei.," This LF is a lower-limit since our variability selection technique, employing only 2 epochs for detection, could miss $\sim$ of variable nuclei."
+ The nuclear magnitudes for our variables extend (o Mj7-15.0 with a median magnitude of Mj2-16.3., The nuclear magnitudes for our variables extend to $_B$$\simeq$ -15.0 with a median magnitude of $_B$ =-16.3.
+ Computing our LF with total integrated. galaxy magnitudes. we compare to the LF of Ulvestad Ho (2001) for local SevIert galaxies.," Computing our LF with total integrated galaxy magnitudes, we compare to the LF of Ulvestad Ho (2001) for local Seyfert galaxies."
+ Hf our, If our
+The QSO 2237|0305. sometimes known as Lluchra’s lens (lluchra ct al.,"The QSO 2237+0305, sometimes known as Huchra's lens (Huchra et al."
+ 1985). is perhaps the most remarkable eravitational lens vet. discovered.," 1985), is perhaps the most remarkable gravitational lens yet discovered."
+ Lt comprises a foreground barred Sb galaxy (220.0394) whose nucleus is surrounded by four images of a radio-quiet QSO (221.695)., It comprises a foreground barred Sb galaxy (z=0.0394) whose nucleus is surrounded by four images of a radio-quiet QSO (z=1.695).
+ Ground based spectroscopic observations have verified that all four images are similar QSO's at the same redshift (Xdam et al., Ground based spectroscopic observations have verified that all four images are similar QSO's at the same redshift (Adam et al.
+ 1989)., 1989).
+ These images are separated by O.7-0.9° [rom cach other and the nucleus., These images are separated by 0.7-0.9” from each other and the nucleus.
+ Broad banc monitoring has shown that significant microlensing events occur (eg., Broad band monitoring has shown that significant microlensing events occur (eg.
+ Irwin et al., Irwin et al.
+ 1989: Corrigan et al., 1989; Corrigan et al.
+ 1991)., 1991).
+ Since the optical depth to microlensing is of order unity at cach of the image positions (eg., Since the optical depth to microlensing is of order unity at each of the image positions (eg.
+ Went Falco 1988: Schneider ct al., Kent Falco 1988; Schneider et al.
+ 1988: Schmidt. Webster Lewis 1998). the magnification effects on the source can be considered as a network of caustics moving across the source plane.," 1988; Schmidt, Webster Lewis 1998), the magnification effects on the source can be considered as a network of caustics moving across the source plane."
+ 2237|0305. provides a unique opportunity to use microlensing events to study the gcometrical structure of the background QSO for two good reasons., Q2237+0305 provides a unique opportunity to use microlensing events to study the geometrical structure of the background QSO for two good reasons.
+ Firstly. the relative closeness of the lensing galaxy means that the time delay between the four quasar images is less than a cay.," Firstly, the relative closeness of the lensing galaxy means that the time delay between the four quasar images is less than a day."
+ Thus it is easy to separate microlensing events from intrinsic variations of the QSO., Thus it is easy to separate microlensing events from intrinsic variations of the QSO.
+ Also. for this QSO the time-scale of microlensing events is typically 30-50 clays. reduced by a factor of 10 from typical time-scales of perhaps a vear.," Also, for this QSO the time-scale of microlensing events is typically 30-50 days, reduced by a factor of 10 from typical time-scales of perhaps a year."
+ Detailed modelling (eg., Detailed modelling (eg.
+ Wambsganss Paczvnski 1991: Fluke Webster. 1999) shows that microlensing events resulting from the source crossing a caustic will show spectral variations which depend on the relative sizes of cillerent emission regions at clillerent wavelengths., Wambsganss Paczynski 1991; Fluke Webster 1999) shows that microlensing events resulting from the source crossing a caustic will show spectral variations which depend on the relative sizes of different emission regions at different wavelengths.
+ Ciround based observations have confirmed differential amplification of the emission region., Ground based observations have confirmed differential amplification of the emission region.
+ Lewis ct al. (, Lewis et al. (
+1998) determine the ratios of emission line equivalent widths relative to one image.,1998) determine the ratios of emission line equivalent widths relative to one image.
+ Phey show (7) the ratios remain fairly constant for one image from line to line. suggesting that the sizes of the emission regions for the lines are not greatly. cüllerent. (77) that the ratios vary from image to image for a single epoch by a factor of —2.5. and (777) that the ratio for a single image varies as a function of time. Le. as a result of a microlensing event.," They show $(i)$ the ratios remain fairly constant for one image from line to line, suggesting that the sizes of the emission regions for the lines are not greatly different, $(ii)$ that the ratios vary from image to image for a single epoch by a factor of $\sim 2.5$, and $(iii)$ that the ratio for a single image varies as a function of time, i.e. as a result of a microlensing event."
+ These results are consistent with earlier. results. of Fillipenko (1989) who measured a ~25% illerence in the width of the Mell lines between the A and D images., These results are consistent with earlier results of Fillipenko (1989) who measured a $\sim25\%$ difference in the width of the MgII lines between the A and B images.
+ A OSO is generally. described as a continuum. source, A QSO is generally described as a continuum source
+5.3 Onthe ,the boundary holonomy as the argument.
+"""Liouville type” representation of theeffective action startwith actions"," The boundary $SL(2,R)$ holonomy can be expressed in terms of the geodesic length A= 2 which in the case under consideration is proportional to $eEs$."
+ for the charged particle ofmass im Z—» exp((—mL(C) +ieC) (C, Hence the Schwinger parameter is proportional to the geodesic boundary length.
+)0.40) where (C)is the socalled spin[actor expressed intermsof thetrajectories, Given these arguments let us remind that the boundary lengths $l_i$ are the natural coordinates on the Teichmueller phase space.
+ in CP and W(C) is theWilson loopalong thecontour C. Thepotential need[or the , It is useful to consider the coherent state representation on the Teichmueller phase space in the Kahler polarization (see \cite{tesh} for the review).
+Liouvillemodeto appear is(he restoration of llowever more convenient Selwinger parameter.Inthe externalsell-dual fieldtheone-loop effectiveaction readsas that£ cH. Let emphasize that the effective where we assume = us actionin QED isnot “correlator” holomorphic(heobject Liouville-tvpehencetheorywe would liketo integrated.represent prod," The Liouville conformal block has an interesting interpretation as the wave function in the quantum Teichmueller theory in the coherent representation which is eigenfunction of the length operator \cite{tesh}
+ (q) = (q) Here $L$ is the length variable in the Teichmueller theory while simultaneously it parameterizes the weights in the Liouville theory via identification Q/2 + Hence we see the clear analogy between the Liouville representation of the ${\cal N}=2$ SYM in the $\Omega$ -background and the QED in the constant external field."
+uct(5.41))as achiralkind of in (hatis of the conformal blocksover the intermediate conformal wei, The integration over the dimension of the intermediate state in the conformal block corresponds to the integration over the Schwinger parameter in QED.
+ghts.Therelore the question is weather the integral over(he Schwinger parameter canbe theory.(reatedas the integral, Let us emphasize once again that the holomorphic prepotential in ${\cal N}=2$ theory corresponds to the conformal block in Liouville theory while the QED effective action plays the role of the correlator.
+ overthe intermediate conformal dimensionsin the Liouville the First AdSnotethat |0.the 0].. Schwingertheotherparameter canbe representedin terms of the wave ΠΙΟΓΙΟof the SL(2.H) 2-dimensiona, Moreover we have mentioned that $s$ has the interpretation of the zero mode of the Liouville field therefore to some extent the quantized values of $s_k$ in the imaginary part calculation correspond to the positions of Liouville ZZ branes \cite{zz} which are localized at this direction.
+lYM theoryonthe disc whichvields (he solutionstothe Ad5»gravity. aud. depends on, In this Letter we have made the first step towards the investigation of the ${\cal N}=2$ SYM vacuum response functions in the graviphoton background focusing on the
+We use the EDisCS dataset to describe galaxies in distant clusters.,We use the EDisCS dataset to describe galaxies in distant clusters.
+ Figure 7 shows the distribution of the local density in this sample and the limits of the 4 density bins., Figure \ref{ediscs_histo} shows the distribution of the local density in this sample and the limits of the 4 density bins.
+" In clusters at high-z, galaxies can assume local density values in a range of almost 2.3 dex."," In clusters at high-z, galaxies can assume local density values in a range of almost 2.3 dex."
+" In Figure 8 we show the mass functions of galaxies in the different bins, compared two by two."," In Figure \ref{ediscs} we show the mass functions of galaxies in the different bins, compared two by two."
+" As for WINGS, we find that there is a clear dependence of the mass function on the local density, in the sense that low density regions have proportionally more low mass galaxies."," As for WINGS, we find that there is a clear dependence of the mass function on the local density, in the sense that low density regions have proportionally more low mass galaxies."
+ The K-S test is able to refuse the null hypothesis of similarity of the populations in the majority of cases., The K-S test is able to refuse the null hypothesis of similarity of the populations in the majority of cases.
+" Again, as for WINGS, also in distant clusters the mass functions in the highest density bin (D4, and perhaps D3) flattens out at low galaxy masses, Adopting a higher mass limit, we find that for logM./Mc> the K-S results are inconclusive, suggesting again that in clusters the effects of the local density on the high mass end shape of the mass function are not visible."," Again, as for WINGS, also in distant clusters the mass functions in the highest density bin (D4, and perhaps D3) flattens out at low galaxy masses, Adopting a higher mass limit, we find that for $\log
+M_{\ast}/M_{\odot}\geq 10.4$ the K-S results are inconclusive, suggesting again that in clusters the effects of the local density on the high mass end shape of the mass function are not visible."
+" Using the ICBS cluster sample (plots not shown), above the mass limit logM../Mc>10.5 we find that the local density range spanned is ~ 2.5 dex, very similar to that of EDisCS."," Using the ICBS cluster sample (plots not shown), above the mass limit $ \log M_{\ast}/M_{\odot}>10.5$ we find that the local density range spanned is $\sim$ 2.5 dex, very similar to that of EDisCS."
+" Moreover, in agreement with the EDisCS findings, we find no dependence at such high masses."," Moreover, in agreement with the EDisCS findings, we find no dependence at such high masses."
+ Our results show that at both redshifts and in all environments there is a dependence of the mass function on the local density., Our results show that at both redshifts and in all environments there is a dependence of the mass function on the local density.
+" Even if we can not perform any inter-sample comparison, since densities have been computed using different criteria, in the following we can qualitatively compare our results coming from the different surveys, to detect if a common trend does exist."," Even if we can not perform any inter-sample comparison, since densities have been computed using different criteria, in the following we can qualitatively compare our results coming from the different surveys, to detect if a common trend does exist."
+" In general, the lower the density, the higher is (proportionally) the number of low mass galaxies, indicating that low-mass galaxies are more common in the “sparsest” regions."," In general, the lower the density, the higher is (proportionally) the number of low mass galaxies, indicating that low-mass galaxies are more common in the “sparsest” regions."
+" shows the cumulative distributions of PM2GC,"," \\ref{cumulate} shows the cumulative distributions of PM2GC,"
+"MIPS22633 looks quite extreme and unusual, showing a very strong ~ aabsorption complex, while being the weakest silicate absorption source in this sample.","MIPS22633 looks quite extreme and unusual, showing a very strong $\sim$ absorption complex, while being the weakest silicate absorption source in this sample."
+" However, we do find a local analog, IRAS23515-2917, which has an even more extreme aabsorption coupled with a relatively shallow ssilicate absorption."," However, we do find a local analog, IRAS23515-2917, which has an even more extreme absorption coupled with a relatively shallow silicate absorption."
+" Lastly our strongest silicate absorption source, MIPS8392 seems to have a mid-IR SED remarkably like that of the hyperluminous galaxy F00183-7111"," Lastly our strongest silicate absorption source, MIPS8392 seems to have a mid-IR SED remarkably like that of the hyperluminous galaxy F00183-7111"
+line of sight through two distinct polarizing regions whose magnetic fields are aligned with {he same orientation. aud a line of sight through two polarizing regions whose magnetic lield orientation on the skv are aligned perpendicular to each other.,"line of sight through two distinct polarizing regions whose magnetic fields are aligned with the same orientation, and a line of sight through two polarizing regions whose magnetic field orientation on the sky are aligned perpendicular to each other."
+ In the first case. as the initially unpolarized background starlisht passes through the first region a weak linear polarization (along the orientation of the magnetic field) is imposed.," In the first case, as the initially unpolarized background starlight passes through the first region a weak linear polarization (along the orientation of the magnetic field) is imposed."
+ The light passes through the second polarizing region which further polarizes the background starlight along (he same orientation., The light passes through the second polarizing region which further polarizes the background starlight along the same orientation.
+ In (he second case. (he unpolarized background starlight again gains polarization [rom the first polarizing region.," In the second case, the unpolarized background starlight again gains polarization from the first polarizing region."
+ However. the second region the starlight and the observer sees much weaker or unpolarized starlight.," However, the second region the starlight and the observer sees much weaker or unpolarized starlight."
+ Because of the possible depolarizing elfect of multiple polarizing regions. a single star may be a poor probe of the full nature of the magnetic field along a line of sight.," Because of the possible depolarizing effect of multiple polarizing regions, a single star may be a poor probe of the full nature of the magnetic field along a line of sight."
+ Therefore. this work uses ensembles of stars (distributed with distance) to characterize each simulated sky. direction.," Therefore, this work uses ensembles of stars (distributed with distance) to characterize each simulated sky direction."
+ The Sun is assumed to be located 3.5 kpc from the Galactic center and in the Galactic midplane., The Sun is assumed to be located 8.5 kpc from the Galactic center and in the Galactic midplane.
+ For each simulated star. the entire line-of-sight [rom the star to the Sun was divided into LO parsee blocks.," For each simulated star, the entire line-of-sight from the star to the Sun was divided into 10 parsec blocks."
+ This size was chosen to be smaller than the tvpical variations found in any of the magnetic and dust models., This size was chosen to be smaller than the typical variations found in any of the magnetic and dust models.
+ At each block. the local dust density was calculated. the magnetic field direction from the models (as seen from the Sun's position) was projected onto the sky. and the changes in Stokes parameters were calculated. by integrating through that block.," At each block, the local dust density was calculated, the magnetic field direction from the models (as seen from the Sun's position) was projected onto the sky, and the changes in Stokes parameters were calculated by integrating through that block."
+ The changes in the Stokes parameters. Irom (1985).. are described by the equations:," The changes in the Stokes parameters, from \citet{LD85}, are described by the equations:"
+G10.6. C19.61. and G20.08 using the models of Myersetal.(1996). in the same way as described in Paper 1.,"G10.6, G19.61, and G20.08 using the models of \citet{Myers96} in the same way as described in Paper 1."
+ From our rms uncertainties in line temperatures and the spectral resolution of our observations. we determined (he uncertainties associated with our derived infall velocities.," From our rms uncertainties in line temperatures and the spectral resolution of our observations, we determined the uncertainties associated with our derived infall velocities."
+ We used these uncertainties to limit the regions in which we claim infall detections., We used these uncertainties to limit the regions in which we claim infall detections.
+ Only if the derived infall velocity is more than three times our uncertainty [or that spectrum do we claim to detect infall in that pixel., Only if the derived infall velocity is more than three times our uncertainty for that spectrum do we claim to detect infall in that pixel.
+ Towards 619.61 and 620.08. robust detections were made in one and three pixels (respectivelv) suggesting that the infall in both of these sources is unresolved.," Towards G19.61 and G20.08, robust detections were made in one and three pixels (respectively) suggesting that the infall in both of these sources is unresolved."
+ Towards GIO.47 and G10.6 (both at approximately 6 kpc). robust infall detections were made in seven Nyequist sampled pixels per map. suggesting (hat the large scale infall in these regions is extencdec.," Towards G10.47 and G10.6 (both at approximately 6 kpc), robust infall detections were made in seven Nyquist sampled pixels per map, suggesting that the large scale infall in these regions is extended."
+ Along the northern. blue shifted. outflow lobe in G10.47. the line profiles again become double peaked.," Along the northern, blue shifted, outflow lobe in G10.47, the $^+$ line profiles again become double peaked."
+ It is quite possible that there is a second instance of a collapsing region here since SiO is mareinally detected (at the 3o level) in the middle of this northern extension., It is quite possible that there is a second instance of a collapsing region here since SiO is marginally detected (at the $\sigma$ level) in the middle of this northern extension.
+ The double peaked profiles do lead to infall velocities (in three pixels) which are greater than 3o: however. we do not include this region in our infall analvsis for a number of reasons.," The double peaked profiles do lead to infall velocities (in three pixels) which are greater than $\sigma$; however, we do not include this region in our infall analysis for a number of reasons."
+ The infall signature is not continuous from the central region to this northern region. which suggests (hat there are (wo distinct regions of infall.," The infall signature is not continuous from the central region to this northern region, which suggests that there are two distinct regions of infall."
+ This discontinuity would confuse conclusions about the central infall region., This discontinuity would confuse conclusions about the central infall region.
+ Similarly. we are interested in the infall activity surrounding ΤΗ regions. where we know massive stars are forming.," Similarly, we are interested in the infall activity surrounding HII regions, where we know massive stars are forming."
+ Neither the mid infrared nor radio maps of this region (Pascuceietal.(2004) and. Wood&Churchwell(1939).. respectivelv) extend. [ar enough to the north to determine whether there are. HII regions in (his northern extension.," Neither the mid infrared nor radio maps of this region \citet{Pascucci04} and \citet{WC89}, respectively) extend far enough to the north to determine whether there are HII regions in this northern extension."
+ Archival SCUBA maps of this region at the same resolution as our observations (DiFrancescoοἱal.2008) show emission which has the same extent and morphology as our observations., Archival SCUBA maps of this region at the same resolution as our observations \citep{scuba08} show emission which has the same extent and morphology as our $^+$ observations.
+ For each region in whieh infall was observed. we determined an average infall velocity bv taking the mean of the infall velocities determined [or each position in each map using ihe models of Myersetal.(1996).," For each region in which infall was observed, we determined an average infall velocity by taking the mean of the infall velocities determined for each position in each map using the models of \citet{Myers96}."
+.. This averaging resulted in the velocities shown in Table 2.., This averaging resulted in the velocities shown in Table \ref{tab:hco_prop}.
+ We chose to use this method over creating an averaged spectrum and determining infall velocities [rom it., We chose to use this method over creating an averaged spectrum and determining infall velocities from it.
+ The latter method resulted in unreasonably high infall rates (ie. 5 km ! in G10.47). which give very large mass infall rates (0.22 M. ') and total infall masses (~2000 AL.) assuming constant infall rates over the lifetime of the outflow.," The latter method resulted in unreasonably high infall rates (i.e. 5 km $^{-1}$ in G10.47), which give very large mass infall rates (0.22 $_{\odot}$ $^{-1}$ ) and total infall masses $\sim 2000$ $_{\odot}$ ) assuming constant infall rates over the lifetime of the outflow."
+ Averaging the spectra and then determining the infall velocity increased. (he infall velocity by increasing the full width at half maximum (FWIIM) of the line., Averaging the spectra and then determining the infall velocity increased the infall velocity by increasing the full width at half maximum (FWHM) of the line.
+ Since the infall velocity varies as the FWIIM squared. (seeEquation9ofMyersetal.1996).. this broadening leads to an overestimate of the infall velocitw.," Since the infall velocity varies as the FWHM squared \citep[see Equation 9 of][]{Myers96}, this broadening leads to an overestimate of the infall velocity."
+ Thus. for our unresolved. sources. the caleulated infall velocities may be upper limits. since there could be unresolved velocity gradients within the JCNIT beam.," Thus, for our unresolved sources, the calculated infall velocities may be upper limits, since there could be unresolved velocity gradients within the JCMT beam."
+ Higher resolution imaging of these sources could decrease our caleulated infall, Higher resolution imaging of these sources could decrease our calculated infall
+"2010),, while the red dashed lines represent LFs from total luminosity (counting all escaped photons without a flux cut).",", while the red dashed lines represent LFs from total luminosity (counting all escaped photons without a flux cut)."
+" While the un-filtered LFs seem to agree with observations of Gronwalletal.(2007) and Ouchietal. (2008), the filtered ones are significantly off."," While the un-filtered LFs seem to agree with observations of \cite{Gronwall07} and \cite{Ouchi08}, the filtered ones are significantly off."
+ The difference comes from the reduction of and fes; due to the flux cut., The difference comes from the reduction of $\La$ and $\fesc$ due to the flux cut.
+" Moreover, the dispersion[yy in the luminosity — halo mass relation at different redshift may cause a large scatter in the LFs."," Moreover, the dispersion in the luminosity – halo mass relation at different redshift may cause a large scatter in the LFs."
+" This plot suggests that the current simulation in this work is not suitable to study a large galaxy population and its statistical properties, because there are too few observable LAEs."," This plot suggests that the current simulation in this work is not suitable to study a large galaxy population and its statistical properties, because there are too few observable LAEs."
+" Moreover, as discussed earlier, the predicted Lya properties may be affected by a number of numerical and physical limitations of our model."," Moreover, as discussed earlier, the predicted $\lya$ properties may be affected by a number of numerical and physical limitations of our model."
+" For example, the one-phase model currently used in the present work may underestimates the density of cold hydrogen gas, and hence underestimates the Lya flux."," For example, the one-phase model currently used in the present work may underestimates the density of cold hydrogen gas, and hence underestimates the $\lya$ flux."
+" We will study the Lyo LFs at different redshift in a forthcoming paper with the improved ART? which incorporates a two-phase ISM model, and a general simulation with mean overdensity in a larger volume (Yajima et al, in preparation)."," We will study the $\lya$ LFs at different redshift in a forthcoming paper with the improved $\art$ which incorporates a two-phase ISM model, and a general simulation with mean overdensity in a larger volume (Yajima et al, in preparation)."
+" As demonstrated above, our model is able to explain a number of observed properties of LAEs at different redshift."," As demonstrated above, our model is able to explain a number of observed properties of LAEs at different redshift."
+" However, we should point out that our current simulations suffer from a number of major limitations which may affect the predicted Lya properties."," However, we should point out that our current simulations suffer from a number of major limitations which may affect the predicted $\lya$ properties."
+" Finally, we stress once again that caution should be taken when comparing directly the results from our calculations to data from a given survey, because, as discussed above, the observed Lyo properties depend sensitively on a number of factors, including galaxy properties, viewing angle, model parameters, and observational threshold."," Finally, we stress once again that caution should be taken when comparing directly the results from our calculations to data from a given survey, because, as discussed above, the observed $\lya$ properties depend sensitively on a number of factors, including galaxy properties, viewing angle, model parameters, and observational threshold."
+" To summarize, we have investigated the Lya properties of progenitors of a local L* galaxy by combining cosmological hydrodynamic simulations with three-dimensional radiative transfer calculations using the new ART? code."," To summarize, we have investigated the $\lya$ properties of progenitors of a local $L^{*}$ galaxy by combining cosmological hydrodynamic simulations with three-dimensional radiative transfer calculations using the new $\art$ code."
+ Our cosmological simulation follows the formation and evolution of a Milky Way-size galaxy and its substructures from redshift z—127 to z=0., Our cosmological simulation follows the formation and evolution of a Milky Way-size galaxy and its substructures from redshift $z = 127$ to $z = 0$.
+" It includes important physics of dark matter, gas dynamics, star formation, black hole growth, and feedback processes, and has high spatial and mass resolutions to resolve a MW-like galaxy at z=0 and its progenitors at higher redshifts."," It includes important physics of dark matter, gas dynamics, star formation, black hole growth, and feedback processes, and has high spatial and mass resolutions to resolve a MW-like galaxy at z=0 and its progenitors at higher redshifts."
+" Our radiative transfer couples Lya line, ionization of neutral hydrogen, and multi-wavelength continuum radiative transfer, which enables self-consistent and accurate calculation of the Lya propertiesain galaxies."," Our radiative transfer couples $\lya$ line, ionization of neutral hydrogen, and multi-wavelength continuum radiative transfer, which enables a self-consistent and accurate calculation of the $\lya$ propertiesin galaxies."
+" We find that the main progenitor of the MW galaxy is Lya bright at high redshift, with the emergent Lya luminosity close to the observed characteristic Ly, of LAEs at z~2— 6."," We find that the main progenitor of the MW galaxy is $\lya$ bright at high redshift, with the emergent $\lya$ luminosity close to the observed characteristic $L_{\lya}^{*}$ of LAEs at $z \sim 2 - 6$ ."
+ Most of the fainter galaxies in the, Most of the fainter galaxies in the
+of LaRosaetal.(1996).,of \citet{LaRosa}.
+. The first term. proporGonal to (22—2) represents the mean acceleration due to the offset in the rates of the two tvpes of reflection.," The first term, proportional to $(R_+ - R_-)$ represents the mean acceleration due to the offset in the rates of the two types of reflection."
+ The second. term. proportional to (RR+R)=R. the total reflection rate. represents the coherent term.," The second term, proportional to $(R_+ + R_-) = R$, the total reflection rate, represents the coherent term."
+ This expression gives a full description of the mean acceleration of charges bv the non-relativistic STFA process., This expression gives a full description of the mean acceleration of charges by the non-relativistic STFA process.
+ To evaluate it in a particular plasma requires the determination of the head-on and catel-up reflection rates Ro and ἰτ , To evaluate it in a particular plasma requires the determination of the head-on and catch-up reflection rates $R_+$ and $R_-$ .
+In order to obtain 2 and HR... consider (he path a charge takes to encounter a nirror.," 	In order to obtain $R_+$ and $R_-$, consider the path a charge takes to encounter a mirror."
+ Note that there is a distinction between an encounter and a mirroring because of the pitch angle condition [or reflection., Note that there is a distinction between an encounter and a mirroring because of the pitch angle condition for reflection.
+ We take the fraction of encounters which reflect to be F and assume (hat this fraction is the same for both head-on and catch-up encounters: FF., We take the fraction of encounters which reflect to be $F$ and assume that this fraction is the same for both head-on and catch-up encounters: $F = F_+ = F_-$ .
+ This assumption is often (taken in the regime where 02c0. (LaRosaetal.1996)..., This assumption is often taken in the regime where $v \gg v_c$ \citep{LaRosa}.
+ While {his assumption is not strictly true. the effects of relaxing it are negligible.," While this assumption is not strictly true, the effects of relaxing it are negligible."
+ In Appendix D we repeat (his calculation without assuming F=Εν, In Appendix B we repeat this calculation without assuming $F_+ = F_-$.
+ There is a well defined mean distance between encounters. Ay as well as a relative velocity between the particle and the compression Unccpmoss ," There is a well defined mean distance between encounters, $\lambda_T$ as well as a relative velocity between the particle and the compression $v_{\pm} = v_\parallel \pm v_c$."
+The small dilference in the head-on and catch-up speeds is responsible for the olfset in rates., The small difference in the head-on and catch-up speeds is responsible for the offset in rates.
+" For each tvpe of encounter. the mean separation is ολ and the rate of reflections of any sort is specified by the general relation A=Fo,4;À. wilh e, the relative velocity. between (he particle and compression. vielding The offset in rates is (hus and we can rewrite (he average acceleration rate as with the associated acceleration “ime scale το=E/(4) "," For each type of encounter, the mean separation is $2 \lambda_T$ and the rate of reflections of any sort is specified by the general relation $R = F v_{rel} \lambda$, with $v_{rel}$ the relative velocity between the particle and compression, yielding The offset in rates is thus and we can rewrite the average acceleration rate as with the associated acceleration time scale $\tau_S = E/\left(\frac {dE}{dt}\right)_S $."
+"We thus see that the acceleration due to (he offset in head-onand catch-up reflection rates ancl the acceleration from the coherent term in c,are equal.", We thus see that the acceleration due to the offset in head-onand catch-up reflection rates and the acceleration from the coherent term in $v_c$are equal.
+"the high-redshift normalizations we see that the shape of their f-parameters at high redshifts agrees remarkably well with our results, even though our treatment is somewhat more simplified.","the high-redshift normalizations we see that the shape of their $f$ -parameters at high redshifts agrees remarkably well with our results, even though our treatment is somewhat more simplified."
+ The sharply peaked dashed curve at z~1000 in Fig., The sharply peaked dashed curve at $z\sim 1000$ in Fig.
+ D| shows the Thomson visibility function (?)., \ref{fig2} shows the Thomson visibility function \citep{1970Ap&SS...7....3S}.
+" More specifically, the quantity Y—zg(z)dlnz gives the probability that the CMB photon last scattered in the logarithmic redshift interval dlnz centered on redshift z."," More specifically, the quantity $\mathcal{V}\sim z\, g(z){\rm d}\ln z$ gives the probability that the CMB photon last scattered in the logarithmic redshift interval ${\rm d}\ln z$ centered on redshift $z$."
+" For the CMB calculations the values of the f-parameter matter only for the redshift range over which the visibility function is large, and in this range our results agree with the ? calculations to better than 5% accuracy."," For the CMB calculations the values of the $f$ -parameter matter only for the redshift range over which the visibility function is large, and in this range our results agree with the \citet{2009PhRvD..80d3526S} calculations to better than $5\%$ accuracy."
+ We also tested the results for the r-channel and found very good agreement., We also tested the results for the $\tau$ -channel and found very good agreement.
+" Thus, our somewhat more simplistic calculation of the f-parameters compared to ? seems to be justified."," Thus, our somewhat more simplistic calculation of the $f$ -parameters compared to \citet{2009PhRvD..80d3526S} seems to be justified."
+ The other lefthand panels of Fig., The other lefthand panels of Fig.
+" (from top to bottom) show the fraction of free electrons, matter temperature, and the Thomson visibility function as a function of redshift for two DM particle masses, 10 GeV and 100 GeV, along with annihilation cross sections {10,100} times, and {100,1000} times the standard thermal cross section (a74v)sta=3x10726 cm?s-!."," \ref{fig2} (from top to bottom) show the fraction of free electrons, matter temperature, and the Thomson visibility function as a function of redshift for two DM particle masses, $10$ GeV and $100$ GeV, along with annihilation cross sections $\{10,100\}$ times, and $\{100,1000\}$ times the standard thermal cross section $\cs_{{\rm std}}=3\times10^{-26}$ $^{3}$ $^{-1}$."
+ See the legend for the line definitions., See the legend for the line definitions.
+" The dotted lines show the standard behavior of ionization fraction, matter temperature and visibility function without any additional energy input from DM annihilation."," The dotted lines show the standard behavior of ionization fraction, matter temperature and visibility function without any additional energy input from DM annihilation."
+ In the panel for the matter temperature we also show the line for the temperature of the CMB., In the panel for the matter temperature we also show the line for the temperature of the CMB.
+" Because the annihilation power scales in proportion to (cAv)n2ympyος(Av)/mpw, we would expect the case with mpm=10 GeV, and (cav)=lO(cAU)sta (100(c Av)sa) give comparable results to the case mpm=100 GeV and (σαυ)=100(o4v)sta (1000(cAU)sa), as long as the f-parameter does not depend strongly on mpw."," Because the annihilation power scales in proportion to $\cs n_{{\rm DM}}^2m_{{\rm DM}}\propto \cs/m_{{\rm DM}}$, we would expect the case with $m_{{\rm DM}}=10$ GeV, and $\cs=10\cs_{\rm std}$ $100\cs_{\rm std}$ ) give comparable results to the case $m_{{\rm DM}}=100$ GeV and $\cs=100\cs_{\rm std}$ $1000\cs_{\rm std}$ ), as long as the $f$ -parameter does not depend strongly on $m_{{\rm DM}}$."
+" This is indeed approximately so, as can be seen from Fig. D]."," This is indeed approximately so, as can be seen from Fig. \ref{fig2}."
+" Even though the lines for the ionization fraction, matter temperature, and CMB visibility function are easily separated at lower redshifts, this simple scaling with (c4v)/mpw seems to hold very well near the peak of the visibility function, and so one would not expect to be able to clearly distinguish the models with the same value of (c4v)/mpw via CMB measurementsalond!]."," Even though the lines for the ionization fraction, matter temperature, and CMB visibility function are easily separated at lower redshifts, this simple scaling with $\cs/m_{{\rm DM}}$ seems to hold very well near the peak of the visibility function, and so one would not expect to be able to clearly distinguish the models with the same value of $\cs/m_{{\rm DM}}$ via CMB measurements."
+" Also, the additional annihilation boost from the structure formation is not expected to influence the CMB measurements."," Also, the additional annihilation boost from the structure formation is not expected to influence the CMB measurements."
+ Only in the very extreme case where instead of ? mass-concentration relation we use a simple power-law extrapolation down to the very low halo masses one is able to cause an additional significantly high low-redshift peak in the visibility function., Only in the very extreme case where instead of \citet{2008MNRAS.391.1940M} mass-concentration relation we use a simple power-law extrapolation down to the very low halo masses one is able to cause an additional significantly high low-redshift peak in the visibility function.
+" However, for any realistic mass-concentration relation, along with annihilation cross sections (c4v) that do not violate the CMB data, the contribution from the structure formation to the CMB signal is completely negligible."," However, for any realistic mass-concentration relation, along with annihilation cross sections $\cs$ that do not violate the CMB data, the contribution from the structure formation to the CMB signal is completely negligible."
+" Also, the contribution to the low-redshift ionization fraction is very mild, and thusthe annihilating DM models which are compatible with CMB measurements could only play a marginal role in reionizing the low-z Universe."," Also, the contribution to the low-redshift ionization fraction is very mild, and thusthe annihilating DM models which are compatible with CMB measurements could only play a marginal role in reionizing the $z$ Universe."
+ T'hese results support our similar findings previously reported in ?.., These results support our similar findings previously reported in \citet{2009A&A...505..999H}.
+" In the righthand panels of Fig. Bl,"," In the righthand panels of Fig. \ref{fig2},"
+" we show the temperature-temperature (TT), temperature-E-mode polarization (TE), and E-mode-E-mode (EE) angular power spectra."," we show the temperature–temperature (TT), temperature–E-mode polarization (TE), and E-mode–E-mode (EE) angular power spectra."
+ The points with error bars show the WMAP seven-year measurements (?).., The points with error bars show the WMAP seven-year measurements \citep{2011ApJS..192...16L}.
+ The dotted lines in all of the panels show theoretical predictions for the concordance ACDM model (?)., The dotted lines in all of the panels show theoretical predictions for the concordance $\Lambda$ CDM model \citep{2011ApJS..192...18K}.
+" As we might already expect, remembering the visibility function behavior as shown in the lowest lefthand panel, the models with the same values for (c4v)/mpw are hardly distinguishable."," As we might already expect, remembering the visibility function behavior as shown in the lowest lefthand panel, the models with the same values for $\cs/m_{{\rm DM}}$ are hardly distinguishable."
+" This is because the redshift dependence of the f- parameter over the width of the visibility function is small, and so only the average value around z~1100 really matters, in agreement with the statements of ?.."," This is because the redshift dependence of the $f$ -parameter over the width of the visibility function is small, and so only the average value around $z\sim 1100$ really matters, in agreement with the statements of \citet{2009PhRvD..80d3526S}."
+" As a final note in this section we point out that, if one increases (c4v) to high enough values and adds the structure formation boost, so that the low-redshift ionization fraction starts to get significantly approach one, the original RECFAST code, along with its DVERK solver, is not properly able to deal with the underlying set of stiff ODEs and simply breaks down."," As a final note in this section we point out that, if one increases $\cs$ to high enough values and adds the structure formation boost, so that the low-redshift ionization fraction starts to get significantly approach one, the original RECFAST code, along with its DVERK solver, is not properly able to deal with the underlying set of stiff ODEs and simply breaks down."
+" To be able to numerically treat these cases, one can use the significantly improved (?), along with its stiff ODE solver."," To be able to numerically treat these cases, one can use the significantly improved \citep{2010MNRAS.tmp.1876C}, along with its stiff ODE solver."
+" However, in codd]our calculations the original RECFAST works fine, since the typically allowed values of (c4v)/mpw compatible with the CMB data are low enough."," However, in our calculations the original RECFAST works fine, since the typically allowed values of $\cs/m_{{\rm DM}}$ compatible with the CMB data are low enough."
+ Now we have all the ingredients available to calculate current and future CMB constraints on annihilating DM., Now we have all the ingredients available to calculate current and future CMB constraints on annihilating DM.
+" We notice that the CMB constraints for the decaying DM are not competitive (e.g. when compared to the bounds available from the low-redshift gamma-ray measurements) due to ο:ppm scaling when compared to the ος02;M scaling for the annihilating DM, and thus are not being discussed in this work."," We notice that the CMB constraints for the decaying DM are not competitive (e.g. when compared to the bounds available from the low-redshift gamma-ray measurements) due to $\propto \rho_{DM}$ scaling when compared to the $\propto \rho_{DM}^2$ scaling for the annihilating DM, and thus are not being discussed in this work."
+" For the currently existing CMB data, we used the latest power spectra, i.e. seven-year spectra, from the WMAP space mission (?).."," For the currently existing CMB data, we used the latest power spectra, i.e. seven-year spectra, from the WMAP space mission \citep{2011ApJS..192...16L}. ."
+ In reality the current CMB bounds can be somewhat tightened if one includes results, In reality the current CMB bounds can be somewhat tightened if one includes results
+ttomoera is simular to that from ΠΟAA.,tomogram is similar to that from HeII.
+. Strong emission sites are scattered around the disk rmi although no emiüssiou appears along the stream., Strong emission sites are scattered around the disk rim although no emission appears along the stream.
+ Civen the velocity resolution of the data. it is difficult to sav with certainty whether the strong emission sites in the tommoerams occur at velocities significautly different frou that expected at the disk vim.," Given the velocity resolution of the data, it is difficult to say with certainty whether the strong emission sites in the tomograms occur at velocities significantly different from that expected at the disk rim."
+ The most reliable maps in this regard would be those from the first VLT dataset for ΠΟΠ aand WellAA., The most reliable maps in this regard would be those from the first VLT dataset for HeII and HeII.
+. In these maps. there is ciission du the stream reeion.," In these maps, there is emission in the stream region."
+ The former also shows cussions sites close to the civcularization velocity., The former also shows emissions sites close to the circularization velocity.
+ These maps. along with that for CTIAA.. all lint that the strezuu may be overflowing the disk.," These maps, along with that for CIII, all hint that the stream may be overflowing the disk."
+ We nieht. alternatively. attempt to explain the velocity of the ciission being iu an area we would associate with disk material in terms of a thick rim model.," We might, alternatively, attempt to explain the velocity of the emission being in an area we would associate with disk material in terms of a thick rim model."
+ If the disk were puffed up bv N-vav irradiation. it would be natural to expect high ionization Hues to appear near the rini.," If the disk were puffed up by X-ray irradiation, it would be natural to expect high ionization lines to appear near the rim."
+ ILowwever. coutemporancous N-aray observations in Paper HI show dips at a wide range of phases.," However, contemporaneous X-ray observations in Paper III show dips at a wide range of phases."
+ Tudeed. siuilaur to the results of Dounet-Didaudetal.(2001)... only the ranges o=0.2 0.3 and o=0.15 0.55 show a of dipping activity.," Indeed, similar to the results of \cite{bonnet01}, only the ranges $\phi=0.2$ $0.3$ and $\phi=0.45$ $0.55$ show a of dipping activity."
+ Given azimuthal disk svuunuetry (or some close approximation) why would one particular phase be singled out for cussion?, Given azimuthal disk symmetry (or some close approximation) why would one particular phase be singled out for emission?
+ Iu the overflowing stream model. this can be attributed to the different physical conditions that exist along the stream and in cach region of the disk with which it interacts with increasing ionization as the central star is approached.," In the overflowing stream model, this can be attributed to the different physical conditions that exist along the stream and in each region of the disk with which it interacts with increasing ionization as the central star is approached."
+ We can compare the same Doppler toimiogriuns with tracks generated usine differeut choices for the 4./ pairs allowed. by the observed eclipses., We can compare the same Doppler tomograms with tracks generated using different choices for the $q$ $i$ pairs allowed by the observed eclipses.
+ This is useful as it is often possible to coustrain further the acceptable range of q (aud hence /) to those cousistent with identifiable catures. in particular the stream.," This is useful as it is often possible to constrain further the acceptable range of $q$ (and hence $i$ ) to those consistent with identifiable features, in particular the stream."
+ We would expect the streai emission to arise at velocities between the balistic and IWeplerian value along the stream trajectory., We would expect the stream emission to arise at velocities between the balistic and Keplerian value along the stream trajectory.
+ The onmoegrani formed from the VLTL observations of the cH liue are shown in Figure 10. with markines ecucrated for a selection of Af values., The tomogram formed from the VLT1 observations of the HeII line are shown in Figure \ref{fig:qdopcomp} with markings generated for a selection of $M_{2}$ values.
+" For M4=1.35M... these appear to avour values of A» at the high eud of the range cousidered in Paper I and the observations are most compatible with he assumption of a main sequence secondary (Af,= ολων ¢=$5.57. Model 3)."," For $M_{1}=1.35M_{\odot}$, these appear to favour values of $M_{2}$ at the high end of the range considered in Paper I and the observations are most compatible with the assumption of a main sequence secondary $M_{2}=0.46M_{\odot}$ , $i=75.5^{\circ}$, Model 3)."
+ ὄνοι(2006) recentlv xoposed a lower limit ἀπ>2.1.., \citet{ozel06} recently proposed a lower limit $M_{1}>2.1M_{\odot}$.
+ However. in that case the Model 3 ballistic trajectory barely passes through he strong ciuission region.," However, in that case the Model 3 ballistic trajectory barely passes through the strong emission region."
+ To achieve the same degree of aerecmeut as for A4=1.35M.. we would require Aly~OSLAM... (plotted as Model D. approximately nore lassive than a iain sequence companion.," To achieve the same degree of agreement as for $M_{1}=1.35M_{\odot}$, we would require $M_{2}\approx0.84M_{\odot}$ (plotted as Model 4), approximately more massive than a main sequence companion."
+ Schenker&Wing(2002) showed how an overmassive secondarv can arise as the result of mass trauster stripping away he secoudary’s euvelope to reveal a helium rich core., \citet{schenker02} showed how an overmassive secondary can arise as the result of mass transfer stripping away the secondary's envelope to reveal a helium rich core.
+ However. the example they eive only las a maxi mass excess of around and then only for Mo«OAL... sienificantly siunaller iu both parameters than we would require.," However, the example they give only has a maximum mass excess of around and then only for $M_{2}<0.1M_{\odot}$, significantly smaller in both parameters than we would require."
+ These difficulties led us to retain the assumption of a main sequence sccoudary with L.35AL.. primary as the choice for the markines m Figures 8 and 9 above., These difficulties led us to retain the assumption of a main sequence secondary with $1.35M_{\odot}$ primary as the choice for the markings in Figures \ref{fig:opttoms} and \ref{fig:hsttoms} above.
+ Stepping back frou the details. we can attempt to identify the conuuon features present in the dataset. focusing on the highest quality lines.," Stepping back from the details, we can attempt to identify the common features present in the dataset, focusing on the highest quality lines."
+ Deeiuniug with the tomogram. these appear to slow two components across several lines;," Beginning with the tomograms, these appear to show two components across several lines."
+ Firstly a riug of e1aission is present that can likely be associated with the accretiou disk (or à coronal laver above it)., Firstly a ring of emission is present that can likely be associated with the accretion disk (or a coronal layer above it).
+ Secoudly. excess enission is usually present ou the left hand side of the tomogram.," Secondly, excess emission is usually present on the left hand side of the tomogram."
+ Iu TellAA.. this is in the upper left quadrant and consisteut with the aceretion stream and/or npact poiut. with the implied disk radius consistent with tidal tyuncation.," In HeII, this is in the upper left quadrant and consistent with the accretion stream and/or impact point, with the implied disk radius consistent with tidal truncation."
+ In the higher ionization resonance lues. most obviously OV and AV. this is preferentially lower iu the tomogram. below the ballistic stream.," In the higher ionization resonance lines, most obviously OV and NV, this is preferentially lower in the tomogram, below the ballistic stream."
+ This sugeests enission from material carried downstream in the disk from the initial inpact., This suggests emission from material carried downstream in the disk from the initial stream-impact.
+ ΠΟΠ aappears as a hvbrid of the two extremes., HeII appears as a hybrid of the two extremes.
+ The velocities of the high excitation lines are lower than expected from a purely ballistic stream overflow. but appear higher than expected from material at the disk rim: they are intermediate between dEKeplerau velocities at the disk riu aud the circularization radius. sugecsting that some peuetration of the disk is occurring.," The velocities of the high excitation lines are lower than expected from a purely ballistic stream overflow, but appear higher than expected from material at the disk rim; they are intermediate between Keplerian velocities at the disk rim and the circularization radius, suggesting that some penetration of the disk is occurring."
+ Liehteurves provide support for this trend., Lightcurves provide support for this trend.
+ ΠΟAA.. which appears to originate at the streau-mupact point in the tomoerams. is deeply eclipsed as expected. whereas the UV lines including ΠΟAA.. which appear to originate from dowustream of the stroaninupact point. show no eclipses.," HeII, which appears to originate at the stream-impact point in the tomograms, is deeply eclipsed as expected, whereas the UV lines including HeII, which appear to originate from downstream of the stream-impact point, show no eclipses."
+ Comparison of the stream position in the Doppler tomograims to models supports the hwpothesis that the secondary is a main sequence star and the system has a lass ratio of 0.31., Comparison of the stream position in the Doppler tomograms to models supports the hypothesis that the secondary is a main sequence star and the system has a mass ratio of 0.34.
+ Examining the trailed spectra. the OV line provides the clearest Savave commpoucut. aud appears to show least disk enission.," Examining the trailed spectra, the OV line provides the clearest S-wave component, and appears to show least disk emission."
+ Its Savave can be clearly ποσα in trailed spectra of other lines when oue allows for their iiultiplet structure where appropriate., Its S-wave can be clearly seen in trailed spectra of other lines when one allows for their multiplet structure where appropriate.
+ Most surprising is that the same S-wave appears to trace out in the Baliner lines., Most surprising is that the same S-wave appears to trace out in the Balmer lines.
+ While it is not required that eas with the same velocity be co-spatial. if it is this poiuts to a two-phase mediuu.," While it is not required that gas with the same velocity be co-spatial, if it is this points to a two-phase medium."
+ Cool material will produce low ionization absorption nes when backlit bv hotter wuderling material. whereas the hot compoucut will produce hieh ionization emission lines. as observed.," Cool material will produce low ionization absorption lines when backlit by hotter underlying material, whereas the hot component will produce high ionization emission lines, as observed."
+ The presence of this compoucut iu IIo at hieh velocities at both phases 0.0 and 0.5 in the VLT2 trailed spectrum indicates that absorption cannot simply be by the disk run but must be bv material above the disk. as the riu will uot be backlit at o=0.5.," The presence of this component in $\alpha$ at high velocities at both phases 0.0 and 0.5 in the VLT2 trailed spectrum indicates that absorption cannot simply be by the disk rim but must be by material above the disk, as the rim will not be backlit at $\phi=0.5$."
+ The picture that thus cmerees is that the accretion stream impacts the disk and some material overflows or penetrates dt. albeit with velocities. closer to the disk rin than to a purely ballistic overflow.," The picture that thus emerges is that the accretion stream impacts the disk and some material overflows or penetrates it, albeit with velocities closer to the disk rim than to a purely ballistic overflow."
+ The overflowing iuaterial forms a two-phase imediun., The overflowing material forms a two-phase medium.
+ The densest chumps remain relatively cool auc produce low ionization absorption of the brighter background disk.," The densest clumps remain relatively cool and produce low ionization absorption of the brighter background disk,"
+"0.0418. On,=0.238 and ον=0.762.",", $\Omega_{\rm m}=0.238$ and $\Omega_\Lambda=0.762$."
+ In addition we assume that 1.3Loar. which is assumed in most of the current simulations.," In addition we assume that $T_s \gg T_{\rm CMB}$, which is assumed in most of the current simulations."
+ The cosmological 21 em maps (075) are simulated. using the algorithm (2).., The cosmological 21 cm maps $\delta T_b$ ) are simulated using the algorithm \citep{thomas09}.
+ is a fast algorithm to simulate the underlying cosmological 21 em signal from the EoR. It is implemented using an N-body/SPH simulation in conjunction with a |-D radiative transfer code under the assumption of spherical symmetry of the ionized bubbles., is a fast algorithm to simulate the underlying cosmological 21 cm signal from the EoR. It is implemented using an N-body/SPH simulation in conjunction with a 1-D radiative transfer code under the assumption of spherical symmetry of the ionized bubbles.
+ The basic steps of the algorithm are as follows: first. a catalogue of ID ionization. profiles of all atomic hydrogen and helium species and the temperature profile that surround the source is calculated for different types of ionizing sources with varying masses. luminosities at different redshifts.," The basic steps of the algorithm are as follows: first, a catalogue of 1D ionization profiles of all atomic hydrogen and helium species and the temperature profile that surround the source is calculated for different types of ionizing sources with varying masses, luminosities at different redshifts."
+ Subsequently. photon rates emanating from dark matter haloes. identified in the N-body simulation. are calculated semi—analytically.," Subsequently, photon rates emanating from dark matter haloes, identified in the N-body simulation, are calculated semi--analytically."
+ Finally. given the spectrum. luminosity and the density around the source. a spherical ionization bubble is embedded around the source. whose radial profile is selected from the catalogue as generated above.," Finally, given the spectrum, luminosity and the density around the source, a spherical ionization bubble is embedded around the source, whose radial profile is selected from the catalogue as generated above."
+ For more details we refer to ?.., For more details we refer to \citet{thomas09}.
+ For the purpose of this paper we use the 27; data cube (2D slices along the frequency/redshift direction) of the cosmological 2| em signal for the “Stars” patchy reionization model (see2).., For the purpose of this paper we use the $\delta T_b$ data cube (2D slices along the frequency/redshift direction) of the cosmological 21 cm signal for the `Stars' patchy reionization model \citep[see][]{thomas09}.
+ The data cube consists of 850 slices in the frequency range from 115MllIz to 200MlIz with 0.1.MlIz step (corresponding to redshift between 6 and 11.5)., The data cube consists of 850 slices in the frequency range from $115~{\rm MHz}$ to $200~{\rm MHz}$ with $0.1~{\rm MHz}$ step (corresponding to redshift between 6 and 11.5).
+ Slices have a size of 100h+ comoving Mpe and are defined on a 512° grid., Slices have a size of $100~\rm{h^{-1}}$ comoving Mpc and are defined on a $512^2$ grid.
+ An example of a random line of sight through simulated 21 em data cube. with angular and frequency resolution matching that of LOFAR. is shown in Fig. 7..," An example of a random line of sight through simulated 21 cm data cube, with angular and frequency resolution matching that of LOFAR, is shown in Fig. \ref{fig:losEOR}."
+ In order to produce the dirty maps of the diffuse emission. we calculate the 2D Fourier transform of the data for each correlation on a fine grid of 1.2arcmin.," In order to produce the dirty maps of the diffuse emission, we calculate the 2D Fourier transform of the data for each correlation on a fine grid of $1.2~{\rm arcmin}$."
+ We assume that there are 24 stations in the LOFAR array that are used for the observations., We assume that there are 24 stations in the LOFAR array that are used for the observations.
+ We then use a bilinear interpolation to estimate the values of the visibilities at the ive points that correspond to the points sampled by the interferometer pairs of the core., We then use a bilinear interpolation to estimate the values of the visibilities at the $uvw$ points that correspond to the points sampled by the interferometer pairs of the core.
+ This is done for Jhrs of synthesis. 10 sec integration and for the whole frequency range between [15 MHz and 180 MHz. using a step of 0.5 MHz.," This is done for 4hrs of synthesis, 10 sec integration and for the whole frequency range between 115 MHz and 180 MHz, using a step of 0.5 MHz."
+ The above procedure is implemented as a parallel algorithm in the CHOPCHOP pipeline (see Labropoulos et al. submitted).," The above procedure is implemented as a parallel algorithm in the CHOPCHOP pipeline (see Labropoulos et al, )."
+ In order to sample the large structure of the foregrounds at scales between 5 and 10 degrees we need interferometer spacing between 6.5 and 13 meters., In order to sample the large structure of the foregrounds at scales between 5 and 10 degrees we need interferometer spacing between 6.5 and 13 meters.
+ Thus the PSF acts as a high-pass spatial filter., Thus the PSF acts as a high-pass spatial filter.
+ Figure 8. shows ‘dirty’ maps of the simulated Galactic synchrotron emission B)) observed with the core stations of the LOFAR telescope at 195MlIz., Figure \ref{fig:dmap} shows `dirty' maps of the simulated Galactic synchrotron emission ) observed with the core stations of the LOFAR telescope at $138~{\rm MHz}$.
+ The total and polarized intensity maps are shown in Fig. 8., The total and polarized intensity maps are shown in Fig. \ref{fig:dmap}.
+.a b. while the polarization angle is presented in Fig. 8.," .a b, while the polarization angle is presented in Fig. \ref{fig:dmap}."
+.c., .c.
+ Note that the large scale structures of the emission are missing as the smallest baseline length is approximately 50m., Note that the large scale structures of the emission are missing as the smallest baseline length is approximately $50~{\rm m}$.
+ One of the major challenges of the EoR experiments is the extraction of the EoR signal from the astrophysical foregrounds., One of the major challenges of the EoR experiments is the extraction of the EoR signal from the astrophysical foregrounds.
+ The extraction is usually formed in total intensity along the frequency direction due to the following characteristics: Thus. the fluctuating EoR signal can be extracted from the foregrounds by fitting the smooth component of the foregrounds out (e.g.seefig.12in?)..," The extraction is usually formed in total intensity along the frequency direction due to the following characteristics: Thus, the fluctuating EoR signal can be extracted from the foregrounds by fitting the smooth component of the foregrounds out \citep[e.g. see fig. 12 in][]{jelic08}."
+ All current EoR radio interferometric arrays have an instrumentally polarized response. which needs to be calibrated.," All current EoR radio interferometric arrays have an instrumentally polarized response, which needs to be calibrated."
+ If the calibration is imperfect. some part of the polarized signal is transferred into a total intensity and vice versa (hereafter leakages’).," If the calibration is imperfect, some part of the polarized signal is transferred into a total intensity and vice versa (hereafter `leakages')."
+ As a result. the extraction of the EoR signal is more demanding.," As a result, the extraction of the EoR signal is more demanding."
+ Moreover. the polarized signal could have similar frequency fluctuations as the cosmological signal and as such could possibly severely contaminate it.," Moreover, the polarized signal could have similar frequency fluctuations as the cosmological signal and as such could possibly severely contaminate it."
+" Thus. to reliably detect the cosmological signal it is essential to minimize the “leakages” and to observe in the regions with the weak polarized foreground emission,"," Thus, to reliably detect the cosmological signal it is essential to minimize the `leakages' and to observe in the regions with the weak polarized foreground emission."
+ We illustrate this through an example for the LOFAR telescope. but the problem is common to all current and planned EoR radio arrays.," We illustrate this through an example for the LOFAR telescope, but the problem is common to all current and planned EoR radio arrays."
+ The ‘leakages’ of the total and polarized signal are produced by two effects: the geometry of the LOFAR array and the cross-talk between the two dipoles in one LOFAR antenna., The `leakages' of the total and polarized signal are produced by two effects: the geometry of the LOFAR array and the cross-talk between the two dipoles in one LOFAR antenna.
+ The cross-talk. a leakage in the electronies that can cause the power from one dipole to be detected with other. is small compared to the geometric effects (Labropoulos et al. submitted) and we will ignore it for the purpose of this paper.," The cross-talk, a leakage in the electronics that can cause the power from one dipole to be detected with other, is small compared to the geometric effects (Labropoulos et al, ) and we will ignore it for the purpose of this paper."
+ The geometry of the LOFAR telescope is such that the array antennae are fixed to the ground., The geometry of the LOFAR telescope is such that the array antennae are fixed to the ground.
+ Therefore. the sources are tracked only by beam-forming and not by steering the antennae mechanically towards the desired. direction.," Therefore, the sources are tracked only by beam-forming and not by steering the antennae mechanically towards the desired direction."
+ This implies that. depending on the position of the source on the sky. a non- (except at the zenith) projection of the two orthogonal dipoles is visible by the source.," This implies that, depending on the position of the source on the sky, a non-orthogonal (except at the zenith) projection of the two orthogonal dipoles is visible by the source."
+ This projection further changes as the source is tracked over time., This projection further changes as the source is tracked over time.
+" Thus. the observed Stokes brightness of the source. S.,5. is given by (2): where Mis a Mueller matrix that quantities the distortions of a true source brightness S=(/.Q.C.V) based on above projection effect."," Thus, the observed Stokes brightness of the source, ${\bf S}_{obs}$, is given by \citep{carozzi09}: where ${\bf M}$ is a Mueller matrix that quantifies the distortions of a true source brightness ${\bf S}={\bf (}I,Q,U,V{\bf)}$ based on above geometry-projection effect."
+ The Mueller matrix is defined as: with the assumption of a coplanar array., The Mueller matrix is defined as: with the assumption of a coplanar array.
+ Note that (7.0.à) are direction cosines and that the level of geometry-projection ‘leakage’ therefore varies across the map.," Note that $(l,m,n)$ are direction cosines and that the level of geometry-projection `leakage' therefore varies across the map."
+" The calculated leakages. for the LOFAR telescope observing at 138Mllz a 55"" pateh of the sky around the zenith. vary between — across the image."," The calculated leakages, for the LOFAR telescope observing at $138~{\rm MHz}$ a $5^\circ \times5^\circ$ patch of the sky around the zenith, vary between – across the image."
+ For the sky model we use the total and polarized intensity maps of Galactic emission (model C)., For the sky model we use the total and polarized intensity maps of Galactic emission (model C).
+ Further. an instant imaging is assumed. i.e. the sky is not tracked over time.," Further, an instant imaging is assumed, i.e. the sky is not tracked over time."
+ Note that the leakages are tiny around the center of the image. but they increase towards the edges.," Note that the leakages are tiny around the center of the image, but they increase towards the edges."
+important potential applications.,important potential applications.
+ For example. by combining radial-velocity anc astrometric observations of S0-2 one could measure the mass. the distance. and the (hree components of motion of Ser A* (Salim&Gould1999).. which is believed to lie almost exactly at the Galactic center.," For example, by combining radial-velocity and astrometric observations of S0-2 one could measure the mass, the distance, and the three components of motion of Sgr A* \citep{visbin}, which is believed to lie almost exactly at the Galactic center."
+" IHowever. the most remarkable characteristic of the Ser À*/S0-2 ""visual binary” is that il exists al all (Ghezetal.2003)."," However, the most remarkable characteristic of the Sgr A*/S0-2 “visual binary” is that it exists at all \citep{ghez03}."
+. 50-2 cannot have lormec in place: the gas cloud from which it would have condensed could never have been larger than ~(ΙΑ.)?AU or it would iive been Udally disrupted.," S0-2 cannot have formed in place: the gas cloud from which it would have condensed could never have been larger than $\sim (m_1/M_\odot)^{1/3}\,\au$ or it would have been tidally disrupted."
+ Here n is the mass of 50-2., Here $m_1$ is the mass of S0-2.
+ Hence. it must have formed at a nuch larger distance [rom Ser A* and was either originally on a highly eccentric orbit that took it within ~125AU of Ser À* or was perturbed onto such an orbit.," Hence, it must have formed at a much larger distance from Sgr A* and was either originally on a highly eccentric orbit that took it within $\sim 125\,\au$ of Sgr A* or was perturbed onto such an orbit."
+ Then. it would have iwl (o have been perturbed again at peribothronu to bring ils eccentricity from (1—€)«1 to (1—6)~0.13.," Then, it would have had to have been perturbed again at peribothron to bring its eccentricity from $(1-e)\ll 1$ to $(1-e)\sim 0.13$."
+ It is this last step that is problematic., It is this last step that is problematic.
+ Spectra al 2; taken by Ghezetal.(2003). exhibit Dr and IHel absorption lines and lack a CO bandheacl., Spectra at $\mu$ m taken by \citet{ghez03} exhibit $\gamma$ and HeI absorption lines and lack a CO bandhead.
+" These features suggest Chat50-2 is à 05-10 dwarl star and so is massive. nm,e15M.. and voung. <10 Myr."," These features suggest thatS0-2 is a O8-B0 dwarf star and so is massive, $m_1\sim 15\,M_\odot$, and young, $\la 10\,$ Myr."
+ The vouth of the star puts a strong limit on ihe length of time available to scatter the star into its current orbit since formation., The youth of the star puts a strong limit on the length of time available to scatter the star into its current orbit since formation.
+ The dvnamical friction time for such a star in an orbit of semi-major axis à~104AU where il might plausibly form. would be fay~300Myr(0/10!AU)L1 even assuming that the stellar profile continued as a power law down to very small radii. as suggested may be the case by Alexander(1999).," The dynamical friction time for such a star in an orbit of semi-major axis $a\sim 10^4\,\au$ where it might plausibly form, would be $t_{\rm df}\sim 300\,{\rm Myr}\,(a/10^4\,\au)^{-1/4}$, even assuming that the stellar profile continued as a power law down to very small radii, as suggested may be the case by \citet{alexander}."
+. Since several dvnamical [rietion times would be required to achieve the orbit of S0-2. this appearsimpossible.," Since several dynamical friction times would be required to achieve the orbit of S0-2, this appears."
+ In this paper we propose that the orbit of SO-2 was a result of the disruption of a massive-star binary in the tidal field of the massive black hole at the Galactic center., In this paper we propose that the orbit of SO-2 was a result of the disruption of a massive-star binary in the tidal field of the massive black hole at the Galactic center.
+ When (he binary disrupts. one star is ejected with energy greater Chan its incoming energy. leaving the other star. SO-2. in a bound orbit with the black hole.," When the binary disrupts, one star is ejected with energy greater than its incoming energy, leaving the other star, SO-2, in a bound orbit with the black hole."
+ In 2.. we show analvtically that for this mechanism (o work. the former binary companion of S0-2 must have been very massive. of order 100 M...," In \ref{sec:anal}, we show analytically that for this mechanism to work, the former binary companion of S0-2 must have been very massive, of order $100\,M_\odot$ ."
+ In J)3.. we confirm this estimate bv numerical simulation. showing (hat the primary star must be m3ο>GOAL.," In \ref{sec:numer}, we confirm this estimate by numerical simulation, showing that the primary star must be $m_2\ga 60\,M_\odot$."
+" In 4. we diseuss how the detection of additional short-period (2<100 vr) ""Iow-eccentricity (6< 0.95) companions to Ser A* could be used to probe very high-mass star formation in the inner pc of the Galactic center."," In \ref{sec:discuss}, we discuss how the detection of additional short-period $P\la 100\,$ yr) “low”-eccentricity $e\la 0.95$ ) companions to Sgr A* could be used to probe very high-mass star formation in the inner pc of the Galactic center."
+ In Appendix A.. we discuss the mechanisms bywhich such massive binaries might be fed to Ser A*.," In Appendix \ref{asec:feed}, , we discuss the mechanisms bywhich such massive binaries might be fed to Sgr A*."
+hand side.,hand side.
+ The average angular has been definedas: having used the axial symmetry of the halo models described in the previous section., The average angular has been defined: having used the axial symmetry of the halo models described in the previous section.
+" Note that v. depends on (£,0,0) through R because of Eqs.(22)) and (23))."," Note that $v_c$ depends on $(\xi, \theta, l)$ through $R$ because of \ref{eq: runo}) ) and \ref{eq: rdue}) )."
+" The gravitomanetic corrections we are investigating are the terms Jj,,, and Iza.) that are of order vc/c."," The gravitomanetic corrections we are investigating are the terms $I_{1a,b}$ and $I_{2a,b}$ that are of order $v_c/c$."
+" It is now straightforward to compute the deflection angle for a spiral galaxy considered as a two component system, i.e. an exponential disk and a dark halo (modelled with the BBS or the PSS profile)."," It is now straightforward to compute the deflection angle for a spiral galaxy considered as a two component system, i.e. an exponential disk and a dark halo (modelled with the BBS or the PSS profile)."
+ The results for the two different modelling are presented and discussed in the next section., The results for the two different modelling are presented and discussed in the next section.
+" To study the detectability of the gravitomagnetic corrections, we have to determine the images positions in order to investigate whether they are significantly modified by these higher order terms."," To study the detectability of the gravitomagnetic corrections, we have to determine the images positions in order to investigate whether they are significantly modified by these higher order terms."
+ This could be a difficult task since the lens equations are highly nonlinear and the results depend also on the source position., This could be a difficult task since the lens equations are highly nonlinear and the results depend also on the source position.
+" However, the images positions are related to the deflection angle so that, if this latter is not modified significantly by gravitomagnetic corrections, we have no detectable effects on the images angular separation which is the easiest quantity to measure."," However, the images positions are related to the deflection angle so that, if this latter is not modified significantly by gravitomagnetic corrections, we have no detectable effects on the images angular separation which is the easiest quantity to measure."
+ We can thus investigate whether the higher order terms introduce appreciable corrections to the total deflection angle., We can thus investigate whether the higher order terms introduce appreciable corrections to the total deflection angle.
+" To this aim, we introduce the:: being a=(o+o2)?1/2 the modulus of the deflection angle."," To this aim, we introduce the: being $\alpha = (\alpha_1^2 + \alpha_2^2)^{1/2}$ the modulus of the deflection angle."
+" Note that this quantity is now function of (£,0) and not only of € since the gravitomagnetic corrections break down the radial symmetry of the deflection angle (Sereno20"," Note that this quantity is now function of $(\xi, \theta)$ and not only of $\xi$ since the gravitomagnetic corrections break down the radial symmetry of the deflection angle \cite{S02,SCnoi02}."
+"02;Sereno&Cardone2002).. In Eqs.(32)) and (33)), we consider only the dependence on the disk mass Ma since the scaling relations described in Sect.22 allow us to fully characterize the model (i.e. disk and halo parameters) as function of this latter quantity."," In \ref{eq: deltai}) ) and \ref{eq: deltamod}) ), we consider only the dependence on the disk mass $M_d$ since the scaling relations described in 2 allow us to fully characterize the model (i.e. disk and halo parameters) as function of this latter quantity."
+" For the PSS model, the ordering parameter is the disk magnitude, but there is a —oone correspondence between M; and Mi so that we will use Ma to compare the results for the BBS and the PSS models."," For the PSS model, the ordering parameter is the disk magnitude, but there is a one correspondence between $M_I$ and $M_d$ so that we will use $M_d$ to compare the results for the BBS and the PSS models."
+" To estimate A; and A we have to fix the values of (£,0)."," To estimate $\Delta_i$ and $\Delta$ we have to fix the values of $(\xi, \theta)$."
+" Since the deflection angle is roughly proportional to the projected mass inside £, we expect the corrections to be most important for larger values of €."," Since the deflection angle is roughly proportional to the projected mass inside $\xi$, we expect the corrections to be most important for larger values of $\xi$."
+" On the other hand, the deflection angle also depends on the angular coordinate 0."," On the other hand, the deflection angle also depends on the angular coordinate $\theta$."
+" We have decided to compute A; and A for three values of £, namely €/ryir=0.5,1.0,1.5, while we fix 0=1/4 to speed up the"," We have decided to compute $\Delta_i$ and $\Delta$ for three values of $\xi$, namely $\xi/r_{vir} = 0.5, 1.0, 1.5$, while we fix $\theta = \pi/4$ to speed up the."
+" Finally, we have also to choose the orientation of the disk plane with respect to the lens one."," Finally, we have also to choose the orientation of the disk plane with respect to the lens one."
+ We arbitrarily fix 6=0 and y=7/4., We arbitrarily fix $\beta = 0$ and $\gamma = \pi/4$.
+" It is reasonable to expect that choosing different values of (8,7) does not change significantly our results."," It is reasonable to expect that choosing different values of $(\beta, \gamma)$ does not change significantly our results."
+" When computing A; and A, we have included the disk contribution to the classical deflection angle, but we have neglected the gravitomagnetic corrections for the disk."," When computing $\Delta_i$ and $\Delta$, we have included the disk contribution to the classical deflection angle, but we have neglected the gravitomagnetic corrections for the disk."
+" In fact, we have checked that the classical deflection angle of the disk is of the same order of magnitude as the gravitomagnetic corrections for the halo at the positions where we calculate A; and A."," In fact, we have checked that the classical deflection angle of the disk is of the same order of magnitude as the gravitomagnetic corrections for the halo at the positions where we calculate $\Delta_i$ and $\Delta$."
+" Hence, the gravitomagnetic corrections for the disk will be much smaller than those for the halo and can be safelyglected""."," Hence, the gravitomagnetic corrections for the disk will be much smaller than those for the halo and can be safely."
+". We stress, however, that the disk plays a key role in our analysis."," We stress, however, that the disk plays a key role in our analysis."
+" First, it is the disk mass Ma which, through the scaling relations discussed in Sect.22, determines the halo parameters and hence the gravitomagnetic corrections."," First, it is the disk mass $M_d$ which, through the scaling relations discussed in 2, determines the halo parameters and hence the gravitomagnetic corrections."
+" Second, the disk potential contributes to the total circular velocity entering Eq.(31))."," Second, the disk potential contributes to the total circular velocity entering \ref{eq: omegamedio}) )."
+ It is thus important to include the, It is thus important to include the
+of equations along the post-Newtouian Dedekind sequence for arbitrary «wy and ay. meaning that the limit presented here can be tacked on continuously.,"of equations along the post-Newtonian Dedekind sequence for arbitrary $w_1$ and $w_2$, meaning that the limit presented here can be tacked on continuously."
+ The axially svuuuetric post-Newtomian solutions we have eenerated. depeud on two parameters or oue if we require that the solution be continuously connected to the PN Dedekind ‘ellipsoids’ with the velocity field(1)., The axially symmetric post-Newtonian solutions we have generated depend on two parameters or one if we require that the solution be continuously connected to the PN Dedekind `ellipsoids' with the velocity field.
+. The solutions are not uniforiulv rotating in ecueral., The solutions are not uniformly rotating in general.
+ If we add this constraint. then requiring that the fourvelocity be shear-free tells us that must hold.," If we add this constraint, then requiring that the four-velocity be shear-free tells us that must hold."
+ We now show that with this additional coustraiut. the solutiou is indeed the PN MNaclaurin solution (there wclemonstrating that the shear free condition is not only necessary. but also sufficient for uniform rotation iu our case).," We now show that with this additional constraint, the solution is indeed the PN Maclaurin solution (thereby demonstrating that the shear free condition is not only necessary, but also sufficient for uniform rotation in our case)."
+ Let us first note that upon taking iuto account the results above uid iu particular Qo>Qu. the componcuts of the velocity become This is precisely the form of the velocity for the post-Newtoniau Maclaurin spheroids. as cau bo found in ? (στα) equation (3). where 9 is a coustaut containing a Newtonian and post-Newtoman contribution. ((28) of that paper.," Let us first note that upon taking into account the results above and in particular $Q_2\to-Q_1$, the components of the velocity become This is precisely the form of the velocity for the post-Newtonian Maclaurin spheroids, as can be found in \citet{Chandrasekhar67_C67} (C67a) equation (3), where $\Omega$ is a constant containing a Newtonian and post-Newtonian contribution, (28) of that paper."
+ Next we note. that for a given equation of state. an anlally sviumetric. stationary and wuiformily rotating fluid is described by two parameters.," Next we note, that for a given equation of state, an axially symmetric, stationary and uniformly rotating fluid is described by two parameters."
+ For our purposes. we can take them to be a;αι which we prescribe using(1)... aud the value for e4. which we leave undetermined.," For our purposes, we can take them to be $a_3/a_1$, which we prescribe using, and the value for $a_1$, which we leave undetermined."
+ Oue has two additional degrees of freedom. which amount to the mapping between a Newtoniau and post-Newtoniui solution and is a matter of convention (cf.ο).," One has two additional degrees of freedom, which amount to the mapping between a Newtonian and post-Newtonian solution and is a matter of convention \citep[cf.][]{Bardeen71}."
+ For example. oue can write the coordinate volume of the star to be where 6 is some relativistic parameter. aux henchovse to have the PN coutribution vanish. Vy=0.," For example, one can write the coordinate volume of the star to be where $\delta$ is some relativistic parameter, and then to have the PN contribution vanish, $V_1=0$."
+ This is the choice tha was mace in CES and C67b aud also in Chaucdrasckhars origina xer on the PN Alaclaurin spheroids στα., This is the choice that was made in CE78 and C67b and also in Chandrasekhar's original paper on the PN Maclaurin spheroids C67a.
+" We ave followed this coveutiou in the current paper. uakiug it casy to compare our results to those of στα,"," We have followed this covention in the current paper, making it easy to compare our results to those of C67a."
+" The second degree of freedom one has was oft unspecified iu much of στα, though Table I ists values with the choice SM=0,5 Tf we introduce the new coordinate and make use of(13).. aud(17).. then the bounding surface (607) in CET8) is given by the-- equation for .the surface z2524 ασUsing12 a5. WhichH holds at the Newtonian-H level aud cau thus be inserted into the PN terii above. one secs that the term with. ez2 vanishes. aud one £uds that the equation is ideutical to (12) of C67a if hold."," The second degree of freedom one has was left unspecified in much of C67a, though Table I lists values with the choice $S_1^\text{M}=S_3^\text{M}=0$ If we introduce the new coordinate and make use of, and, then the bounding surface (47) in CE78) is given by Using the equation for the surface $\varpi^2/a_1^2=1-x_3^2/a_3^2$ , which holds at the Newtonian level and can thus be inserted into the PN term above, one sees that the term with $x_1^2$ vanishes and one finds that the equation is identical to (42) of C67a if hold."
+ As mentioned in that paper. Sol=0 nmav be chosen without loss of," As mentioned in that paper, $S_3^\text{M}=0$ may be chosen without loss of"
+"xoader than the feature due to a DLA. which falls oll as AA, gather than the AA,* scaling for IGM absorption. where AA,, is the dilference between a given wavelength and hat of Lya in the GRB-Llrame (Miralda-","broader than the feature due to a DLA, which falls off as ${\Delta
+\lambda_{\alpha}}^{-2}$ rather than the ${\Delta
+\lambda_{\alpha}}^{-1}$ scaling for IGM absorption, where ${\Delta
+\lambda_{\alpha}}$ is the difference between a given wavelength and that of $\alpha$ in the GRB-frame \citep{miralda98}."
+IEscude1998).. Figure 1 illustrates the effect. of the column density of the DLA (op panel) and the size of the LULL region (bottom panel) on the absorption feature., Figure \ref{fig:params} illustrates the effect of the column density of the DLA (top panel) and the size of the HII region (bottom panel) on the absorption feature.
+" Here we adopt he same simplistic parameterisation as in Equation (2)): an LIL region of size A, surrounded. by a homogeneously neutral LGAL with neutral fraction arg.", Here we adopt the same simplistic parameterisation as in Equation \ref{eqn:tau2}) ): an HII region of size $R_b$ surrounded by a homogeneously neutral IGM with neutral fraction $\bar{x}_H$.
+ The dashed. lines in the top panel represent an IGM. with wy= 0. and the solid lines represent the case with rq=0.5 and ἐν=10 Alpe.," The dashed lines in the top panel represent an IGM with $\bar{x}_H = 0$ , and the solid lines represent the case with $\bar{x}_H = 0.5$ and $R_b = 10$ Mpc."
+ ICM. absorption produces a wider feature than that from a DLA., IGM absorption produces a wider feature than that from a DLA.
+ However. it becomes hard to visually discern absorption owing to a neutral LGA for Nat107em7.," However, it becomes hard to visually discern absorption owing to a neutral IGM for $N_{\rm HI} \gtrsim
+10^{21}~\cmeter^{-2}$."
+" ‘This illustration suggests that it is cüllicult to measure wy from GRDB050904. which has Nyyz1075""em etal. 200G6)."," This illustration suggests that it is difficult to measure $\bar{x}_H$ from GRB050904, which has $N_{\rm HI}
+\approx 10^{21.6}~\cmeter^{-2}$ \citep{totani06}."
+. We return to this point in Section 5.., We return to this point in Section \ref{GRB050904}.
+" DLAs with Nyyc107em? are associated with all but one GRB for which Ny, has been measured (Chenοἱal. 2007).", DLAs with $N_{\rm HI} > 10^{19} \; {\rm cm}^{-2}$ are associated with all but one GRB for which $N_{\rm HI}$ has been measured \citep{chen07}.
+. The cumulative cüstribution of Nyy for the current sample of zz30 GRB DLAs scales as Nyy’ between 1077 and 10757em.7. and about half of all GRBs for which INgi has been measured have Nyy;>1021'cm27 (Chen¢tal.2007).," The cumulative distribution of $N_{\rm HI}$ for the current sample of $\approx 30$ GRB DLAs scales as $N_{\rm HI}^{0.3}$ between $10^{18}$ and $10^{21.5} \;{\rm cm}^{-2}$, and about half of all GRBs for which $N_{\rm HI}$ has been measured have $N_{\rm HI} >
+10^{21.5} \; {\rm cm}^{-2}$ \citep{chen07}."
+ llowever. it is not clear how Nyy should scale with redshift.," However, it is not clear how $N_{\rm HI}$ should scale with redshift."
+ —- is plausible that. since the average galaxy becomes Less massive with redshift. z6 galaxies should. on average. have weaker DLAs.," It is plausible that, since the average galaxy becomes less massive with redshift, $z >6$ galaxies should, on average, have weaker DLAs."
+ For galaxies to reionise the Universe. the escape fraction of ionising photons must be appreciable. implving that sight-lines with Nar107em? must The bottom. panel in Figure 1 depicts the elect. of bubble size on the GRB afterglow spectrum.," For galaxies to reionise the Universe, the escape fraction of ionising photons must be appreciable, implying that sight-lines with $N_{\rm HI} \lesssim 10^{18}\; {\rm
+cm}^{-2}$ must The bottom panel in Figure \ref{fig:params} depicts the effect of bubble size on the GRB afterglow spectrum."
+ The LGAL absorption for CiluDs in bubbles with sizes between 1-10 Alpe is comparable., The IGM absorption for GRBs in bubbles with sizes between $1$ $10 ~\Mpc$ is comparable.
+" Llowever. if the CRB sits in a large bubble with /?,~50Alpe. the absorption is strongly diminished."," However, if the GRB sits in a large bubble with $R_b \sim 50 ~\Mpc$, the absorption is strongly diminished."
+ When rg=0.2. roughly half of the skewers from Cl1tDs in our simulations of reionisation sit in LLL regions that are larger than 50 Alpe (Section 3)).," When $\bar{x}_H = 0.2$, roughly half of the skewers from GRBs in our simulations of reionisation sit in HII regions that are larger than $50$ Mpc (Section \ref{reionisation}) )."
+ While the GRB itself does not. ionise the LGAL around it. earlier star formation [rom its host and neighbouring ealaxies can grow a Large LIL region.," While the GRB itself does not ionise the IGM around it, earlier star formation from its host and neighbouring galaxies can grow a large HII region."
+ To model reionisat around a GRB. we employ three. radiative transferion simulations that are. cach 186.AMIpe on ai side.," To model reionisation around a GRB, we employ three radiative transfer simulations that are each $186 \; \Mpc$ on a side."
+ These simulations are post-processed on a 102475. particle N-body Geld., These simulations are post-processed on a $1024^3$ particle N-body field.
+ Uneresolvecl halos are. included: with extended Press-Schechter merger trees., Unresolved halos are included with extended Press-Schechter merger trees.
+ These three simulations are described in detail in MeQuinnetal.(2007b) and. the methodology is described in MeQuinnetal. (2007a)]]. andthey are meant to span the range of plausible morphologies for reionisation by galaxies.," These three simulations are described in detail in \citet{mcquinn07} [and the methodology is described in \citet{mcquinn06b}] ], andthey are meant to span the range of plausible morphologies for reionisation by galaxies."
+ Here is a briefdescription ofthe models for the sources and sinks of ionising photons used in cach of the three simulations:, Here is a brief description of the models for the sources and sinks of ionising photons used in each of the three simulations:
+The only optical. photoclectric photometry comes from the Carlsberg Meridian. Catalog (CAIC) where two lncasurements are eiven as 5=16.18+0.08 munda.,"The only optical, photoelectric photometry comes from the Carlsberg Meridian Catalog (CMC) where two measurements are given as $r'=16.48\pm0.08$ mag."
+ loester&Wilken(2006) eive Z;4g= 102001 and logg=7.9 for the white dwarf. similar to the plotted model that agrees well with theGALEN far- auc near-ultraviolet fluxes.," \citet{koe06} give $T_{\rm eff}=10\,200$ K and $\log\,g=7.9$ for the white dwarf, similar to the plotted model that agrees well with the far- and near-ultraviolet fluxes."
+ If the model used here is sufficieutlv accurate. the putative excess at | 0716 appears at the 1.6o level at all IRAC waveleneths.," If the model used here is sufficiently accurate, the putative excess at $+$ 0746 appears at the $4-6\sigma$ level at all IRAC wavelengths."
+345. TheSpitzer data for this star is analvzed elsewhere. but theAAARZ IBRC fluxes reported here are original aud appear broadly consistent with the infrared and Spitzer IRAC fluxes (Meliset.al.2010:Brinkworthetal. 2010)..303.," The data for this star is analyzed elsewhere, but the IRC fluxes reported here are original and appear broadly consistent with the near-infrared and IRAC fluxes \citep{mel10,bri10}. ."
+ The jou TRAC fiux at 079 indicates a lo excess (Figure 21) while all other photometric data are consistent with photospheric cluission., The $\mu$ m IRAC flux at $-$ 079 indicates a $4\sigma$ excess (Figure \ref{fig2}) ) while all other photometric data are consistent with photospheric emission.
+ If real. this star joins G166-58 as a moetalxich white dwarf with a long wavelength TRAC excess but uo short waveleneth excess: these two objects may form a class of object whose orbiting disks have laree mner holes.," If real, this star joins G166-58 as a metal-rich white dwarf with a long wavelength IRAC excess but no short wavelength excess; these two objects may form a class of object whose orbiting disks have large inner holes."
+ 3303 may have a mareinal photometric excess at the 2o level (Figure 2)). 630., 303 may have a marginal photometric excess at the $2\sigma$ level (Figure \ref{fig2}) ). .
+ Tis metal-polluted white darf slows the clearest mid-infrared excess (Figure 6)) of all the science targets., This metal-polluted white dwarf shows the clearest mid-infrared excess (Figure \ref{fig3}) ) of all the science targets.
+ Ἱνουπίο&Wilken(2006) eive Tig=TTS for this star.," \citet{koe06} give $T_{\rm eff}=10\,100$K for this star."
+"CGD190. The 2MASS AK,=LL96+ 0.12nuuag mcasurement (Fieure 23) for this helinm- and carbon-rich white dwarf appears erroneous at the 2.450 level.", The 2MASS $K_s=14.96\pm0.13$ mag measurement (Figure \ref{fig2}) ) for this helium- and carbon-rich white dwarf appears erroneous at the $2.5\sigma$ level.
+205. The fit to the spectral euerev distribution (SED) shown in Figure 2. ορίου a higher Tig than that found by previous studies: KIS (Vossetal. 2007).. TIS Usoesteretal.2005).," The fit to the spectral energy distribution (SED) shown in Figure \ref{fig2} employs a higher $T_{\rm eff}$ than that found by previous studies; K \citep{vos07}, K \citep{koe05}."
+. ITowever. the TIN inodel better matches all the photometric data. including the GALLEN fluxes.," However, the K model better matches all the photometric data, including the fluxes."
+235. While the native pixel scale IRAC images of this star reveal a siugle source (single Airy disk) with an elongated core. the 076ppixel + mosaics spatially resolve a neighboring source at 270 and position angle 336°. as measuredbydaophot.," While the native pixel scale IRAC images of this star reveal a single source (single Airy disk) with an elongated core, the $0\farcs6$ $^{-1}$ mosaics spatially resolve a neighboring source at $2\farcs0$ and position angle $336\arcdeg$, as measured by."
+ The 2MÁASS J- aud JI-baud nuages of the white dwarf may show a slisht elongation in the expected direction.but the S/N in these images is too low to be certain (especially at Z7-baud).," The 2MASS $J$ - and $H$ -band images of the white dwarf may show a slight elongation in the expected direction, but the S/N in these images is too low to be certain (especially at $H$ -band)."
+ Photometric decouvolutionwas performed with daophot. vielius good results at3.6 and 15 juu: errorsof and? respectively for the white dwarf.," Photometric deconvolutionwas performed with , yielding good results at3.6 and 4.5 $\mu$ m; errorsof and respectively for the white dwarf."
+ While the ucighboring source is likely extragalacticin nature. there is no," While the neighboring source is likely extragalacticin nature, there is no"
+conceivable that f iucreases with the parent mass Afy a nore massive parent welt encourage more quasar activity * haviug a higher fraction of mereers or collisions.,conceivable that $f$ increases with the parent mass $M_0$ – a more massive parent might encourage more quasar activity by having a higher fraction of mergers or collisions.
+ The ollowing heuristic argument shows that one might expect IN(AR.ALco)fdARPAR.AL) to scale approximately as ~ALLO ," The following heuristic argument shows that one might expect $[d{\cal N}(M_1, z_1 | M, z_0) / dM_1] f(M_1, M)$ to scale approximately as $\sim M^{4/3}$."
+Let Ny be the umber X subhalos inside a xweut halo of mass AM., Let $N_h$ be the number of sub–halos inside a parent halo of mass $M$.
+" The rate o00 lisions is given bv Ngeno£TU where ey is the velocity of the subhalos. σι is rei cross-section and £F"" is the volume of the pareut halo."," The rate of collisions is given by $N_h^2 v_h \sigma_h
+/R^3$ where $v_h$ is the velocity of the sub–halos, $\sigma_h$ is their cross-section and $R^3$ is the volume of the parent halo."
+ —Hug NoxM (which cau be obtaiued from eq., Using $N_h \propto M$ (which can be obtained from eq.
+" 11: in the aree AM lait). eyxR (viria heorem) and 2?xAL Bxed overdensity of VM~200 at the redshift of formation). 1e vate of collisious scales with the parent mass as AL!3,m3 oef one ignores the possibility that o; might depeud on le parent mass as well"," \ref{eq:condPS} in the large $M$ limit), $v_h
+\propto \sqrt{M/R}$ (virial theorem) and $R^3 \propto M$ (fixed overdensity of $\sim 200$ at the redshift of formation), the rate of collisions scales with the parent mass as $M^{4/3}$, if one ignores the possibility that $\sigma_h$ might depend on the parent mass as well."
+ A similar scaling of ΑΣ has con observed in simulations of the starburst model for Lyvuuiurbreak objects (IKolatt et al., A similar scaling of $M^{1.3}$ has been observed in simulations of the star-burst model for Lyman-break objects (Kolatt et al.
+ 1999. Weschler et al.," 1999, Weschler et al."
+ 1999)., 1999).
+ To model the cuhanced rate of collisions iuside massive parent halos. we can simply modify model A by using FORM)η Ομ," To model the enhanced rate of collisions inside massive parent halos, we can simply modify model A by using $f(M_1, M) = (M/M_1)^{1/3}$."
+ effective bias is given by We fud that the above prescription docs ποῖ sienificautly alter our conclusions following from model A: the ALY? cuhancement of the activation rate inside nassive parent halos turus out to be relatively shallow. and trauslate to a sinall effect in the bias.," The effective bias is given by We find that the above prescription does not significantly alter our conclusions following from model A: the $M^{1/3}$ enhancement of the activation rate inside massive parent halos turns out to be relatively shallow, and translate to a small effect in the bias."
+" Finally. we note hat 2, above could in principle depend on Aj aud ALY. a oossibilitv that would require further modeling aud is not o»ursued here."," Finally, we note that $z_1$ above could in principle depend on $M_0$ and $M_1$, a possibility that would require further modeling and is not pursued here."
+ &uother possibility that could uodifv our picture is that onlv a fraction f<1 of the iios harbor DIEs: the duty cvele could then reflect this yaction. rather than the lifetime of quasars.," Another possibility that could modify our picture is that only a fraction $f<1$ of the halos harbor BH's; the duty cycle could then reflect this fraction, rather than the lifetime of quasars."
+ Although here is evidence (e.g. Magorrian et al., Although there is evidence (e.g. Magorrian et al.
+ 1998) that most galaxies harbor a ceutral DIT. this is not necessarily he case at redshifts +=23: the fraction f of galaxies rosting BIVs at 2=23 could. in principle have iierged with the fraction 1f of galaxies without DITs. satistving he local constraint.," 1998) that most galaxies harbor a central BH, this is not necessarily the case at redshifts $z=2-3$: the fraction $f$ of galaxies hosting BH's at $z=2-3$ could, in principle have merged with the fraction $1-f$ of galaxies without BH's, satisfying the local constraint."
+" Using the extended PressSchechter ομαδα (Lacey Cole 1993). oue can compute the rate of merecrs )etween halos of various masses,"," Using the extended Press–Schechter formalism (Lacey Cole 1993), one can compute the rate of mergers between halos of various masses."
+ On galaxymass scales. he merecr rates at 2=3 are comparable to the reciprocal of the IIubble time (ef.," On galaxy–mass scales, the merger rates at $z=2-3$ are comparable to the reciprocal of the Hubble time (cf."
+ Fig., Fig.
+ 5 in Hanau Alenou 2000). inuplvius that a typical galaxy did uot eo hrough uunerous major merecrs between 2=3 and 2=Q 10. that the fraction f cannot be significautly less han unitv at:=23.," 5 in Haiman Menou 2000), implying that a typical galaxy did not go through numerous major mergers between $z=2-3$ and $z=0$, i.e. that the fraction $f$ cannot be significantly less than unity at $z=2-3$."
+ A more detailed: cousideratiou of this issue is bevond the scope of this paper: we simply rote that the lifetimes derived here scale approximately as fo where f is likely of order nity.," A more detailed consideration of this issue is beyond the scope of this paper; we simply note that the lifetimes derived here scale approximately as $1/f$, where $f$ is likely of order unity."
+ The scatter o we assuned around the mean relation between quasar huninositv aud halo mass is motivated by the scatter found enipiricallv for the μμ Ainuec relation (Magorrian et al., The scatter $\sigma$ we assumed around the mean relation between quasar luminosity and halo mass is motivated by the scatter found empirically for the $M_{\rm bh}-M_{\rm bulge}$ relation (Magorrian et al.
+ 1998)., 1998).
+ It is interesting to consider the seusitivitv of our conclusions to an increased σ., It is interesting to consider the sensitivity of our conclusions to an increased $\sigma$.
+" Iu general scatter raises the ΠΡΟ of quasars predicted by our models. by au amount that depends ou the slope of the underline mass function dN11,"," In general, scatter raises the number of quasars predicted by our models, by an amount that depends on the slope of the underlying mass function $dN/dM$."
+ As a result. increasing o raises. aud flattens the predicted LF.," As a result, increasing $\sigma$ raises, and flattens the predicted LF."
+ We find that an increase of 0 from 0.5 to L (an adelitional order of magnitude of scatter) cau be compensated by a steeper Lay relation. typically replacing o with sa0.5.," We find that an increase of $\sigma$ from 0.5 to 1 (an additional order of magnitude of scatter) can be compensated by a steeper $\bar L_M$ relation, typically replacing $\alpha$ with $\approx \alpha-0.5$."
+ As a result of the increase in 6. quasars with a fixed £ are. on average. associated with larger. aud more highly biased halos.," As a result of the increase in $\sigma$, quasars with a fixed $L$ are, on average, associated with larger, and more highly biased halos."
+ Nevertheless. we find that the mean bias 5 of all sources above a fixed dux (cf.," Nevertheless, we find that the mean bias $\bar b$ of all sources above a fixed flux (cf."
+ eq. D).," eq. \ref{eq:bave}) ),"
+ and therefore the correlation leugth rg. is uuchauged bv the increased scatter (at the level of ~3%).," and therefore the correlation length $r_0$, is unchanged by the increased scatter (at the level of $\sim 3\%$ )."
+ The reason for the iusensitivitv of ry to the amplitude of the scatter can be understood as follows., The reason for the insensitivity of $r_0$ to the amplitude of the scatter can be understood as follows.
+ The mean bias 6 of all sources with £o>μμ 1s donated by the bias b(£) of sources near the threshold Lis ," The mean bias $\bar b$ of all sources with $L>L_{\rm
+min}$ is dominated by the bias $b(L)$ of sources near the threshold $L_{\rm
+min}$ ."
+The latter is obtained by averaging b(AL) over Lalos of different masses (cf., The latter is obtained by averaging $b(M)$ over halos of different masses (cf.
+ eq. 1).," eq. \ref{eq:biasL}) ),"
+" aud it is dominated by the bias of the smallest halos within the width of the scatter. Le. of halos with mass A44,zAL/1L07%. where Af defines the mean relation between £44, aud halo nass be. Lag, = LEAL). aud σ quautifics the scatter (see eq. ))."," and it is dominated by the bias of the smallest halos within the width of the scatter, i.e. of halos with mass $M_{\rm
+min} \approx \bar M / 10^\sigma$, where $\bar M$ defines the mean relation between $L_{\rm min}$ and halo mass i.e. $L_{\rm min}$ = $\bar L(\bar M)$, and $\sigma$ quantifies the scatter (see eq. \ref{eq:scatter}) )."
+ As mentioned before. increasing the scatter males the πιος1 Muction flatter. which means to match the observed. τνabundance of halos at a ποια Iuuinositv Layin. oue las to choose a higher M.," As mentioned before, increasing the scatter makes the luminosity function flatter, which means to match the observed abundance of halos at a fixed luminosity $L_{\rm
+min}$, one has to choose a higher $\bar M$."
+" Iu other words. AL scales up with the scatter. and it turus out to scale up approximately as 10. making Mig, aud hence the effective bias roughly independent of scatter."," In other words, $\bar M$ scales up with the scatter, and it turns out to scale up approximately as $10^\sigma$, making $M_{\rm min}$ and hence the effective bias roughly independent of scatter."
+" We note that the relation between quasar Iuninuositv aud halo mass cau. in principle be derived from observations. by measuriug Mya, for the hosts of quasars (e.g. by weak leuxiug. or by finding test particles around quasars. such as nearby satellite galaxies)."," We note that the relation between quasar luminosity and halo mass can, in principle be derived from observations, by measuring $M_{\rm halo}$ for the hosts of quasars (e.g. by weak lensing, or by finding test particles around quasars, such as nearby satellite galaxies)."
+ Tu all of the above. we have assmned that the quasar Lifetime is a single paralcter. independent of the halo mass.," In all of the above, we have assumed that the quasar lifetime is a single parameter, independent of the halo mass."
+" This is not ""reasonable if the Eddineton time. the timescale for the erowth of black hole mass. is indeed the relevant timescale. ο yr."," This is not unreasonable if the Eddington time, the timescale for the growth of black hole mass, is indeed the relevant time--scale, $4\times10^7$ $\epsilon$ /0.1) yr."
+ Implicit in such reasoning is that the active phase of the quasar is coimcident with the plase where the black hole gaius most ofits mass., Implicit in such reasoning is that the active phase of the quasar is coincident with the phase where the black hole gains most of its mass.
+ This is not the only possibility: sce Haehnelt et al. (, This is not the only possibility; see Haehnelt et al. (
+1999) for more diseussious.,1999) for more discussions.
+ One cau attempt to explore how fo depeuds ou halo mass by applying our method to quasars erouped iuto different absolute luminosity ranges. but the intrinsic scatter in the mass-huninosity relation should be kept iu mind.," One can attempt to explore how $t_Q$ depends on halo mass by applying our method to quasars grouped into different absolute luminosity ranges, but the intrinsic scatter in the mass-luminosity relation should be kept in mind."
+ We curphasize. however. since we fit the Iuninositv function and clustering data at the same redshift. there is no need within our formalisur to assume a redshift incdependeut lifetime.," We emphasize, however, since we fit the luminosity function and clustering data at the same redshift, there is no need within our formalism to assume a redshift independent lifetime."
+ Ta fact. performing our exercise as a function of redshift could give interesting coustraiuts ou how fo evolves with redshift.," In fact, performing our exercise as a function of redshift could give interesting constraints on how $t_Q$ evolves with redshift."
+ Iu this paper. we have modeled thequasar luuinosity," In this paper, we have modeled thequasar luminosity"
+‘To solve this problem numerically. we use a pseudo-spectral method with Chebyshey polynomials.,"To solve this problem numerically, we use a pseudo-spectral method with Chebyshev polynomials."
+ We use a Gauss grid. ο)=cosa£N). where IN ds the number of erid points. and the Chebyshey coordinate —1<r<lis related to the radial coordinate by so that rman<r<rua. Where ry=nues and faa ds an outer radius chosen to be [Larger than the outer Lindblad resonance for the r mode.," We use a Gauss--Lobatto grid, $x(i)=\textrm{cos}(\pi i /N)$, where $N$ is the number of grid points, and the Chebyshev coordinate $-1 < x < 1$ is related to the radial coordinate by so that $r_\mathrm{in} < r < r_\mathrm{out}$, where $r_\mathrm{in}=r_\mathrm{ms}$ and $r_\mathrm{out}$ is an outer radius chosen to be larger than the outer Lindblad resonance for the r mode."
+ The representation of the first derivatives in the matrix A is achieved using the Chebyshey collocation derivative matrix. defined bv (Bove2001) where po=pr2 and p;=1 otherwise.," The representation of the first derivatives in the matrix $\mathbf{A}$ is achieved using the Chebyshev collocation derivative matrix, defined by \citep{boydbook}
+ where $p_0=p_N=2$, and $p_j=1$ otherwise."
+ The eencralized eigenvalues ancl eigenfunctions of A are then calculated numerically. using IDL's eigenvalue solver. textrmELCENPROBLEAL which uses a QL decomposition. and is based on LAVPACI routines.," The generalized eigenvalues and eigenfunctions of $\mathbf{A}$ are then calculated numerically, using IDL's eigenvalue solver, IGENPROBLEM, which uses a QR decomposition, and is based on LAPACK routines."
+ Phe boundary conditions used are the following: We allow for the possibility that o is complex. in which case its imaginary part is the growth rate of the disturbance.," The boundary conditions used are the following: We allow for the possibility that $\omega$ is complex, in which case its imaginary part is the growth rate of the disturbance."
+ In fact. in the absence of non-linear mode couplings. we obtain slowly decaving solutions with Γη)«0 as a result of the oultgoine-wave outer boundary condition.," In fact, in the absence of non-linear mode couplings, we obtain slowly decaying solutions with $\mathrm{Im}(\omega)<0$ as a result of the outgoing-wave outer boundary condition."
+ In Fie., In Fig.
+" 1 owe show the variation of o, with radius for two typical trapped solutions. corresponding to two cdilferent radial mode numbers /=0 and /=1. ("," \ref{freermode} we show the variation of $u_r$ with radius for two typical trapped solutions, corresponding to two different radial mode numbers $l=0$ and $l=1$. ("
+Since we are solving an eigenvalue problem. solutions are multiplied by an arbitrary amplitude.),"Since we are solving an eigenvalue problem, solutions are multiplied by an arbitrary amplitude.)"
+ The complex. frequencies of the modes represented in Fig., The complex frequencies of the modes represented in Fig.
+" 1 are 0.03196.3.688410. ""Land 0.02989—1.184610""i (in units of e /C/AL). respectively."," \ref{freermode}
+ are $0.03196-3.6884\times10^ {-7}\textrm{i}$ and $0.02989-1.1846\times10^ {-7}\textrm{i}$ (in units of $c^3/GM$ ), respectively."
+ In these units. and for the value of e used. the maximum of & is 0.03312.," In these units, and for the value of $a$ used, the maximum of $\kappa$ is 0.03312."
+" ""μονο modes are slightly. clampecl because of the boundary. condition used at the outer radius. which selects the outgoing wave only."," These modes are slightly damped because of the boundary condition used at the outer radius, which selects the outgoing wave only."
+ HE the sound speed is smaller. the frequeney of the modes is closer to the value of the epievclic frequency at its maximum. ancl the trapping region and damping rate are smaller. (sec Tab. 1)).," If the sound speed is smaller, the frequency of the modes is closer to the value of the epicyclic frequency at its maximum, and the trapping region and damping rate are smaller (see Tab. \ref{tablefree}) )."
+ We obtained identical results by solving the two-point bouncdary-value problem using a shooting method., We obtained identical results by solving the two-point boundary-value problem using a shooting method.
+ These inertial modes. can be thought of as waves trappec in a virtual potential. C(r)=ko (Lietal. 2003).. which for a frequency close to the maximum of W is similar to the harmonic oscillator potential.," These inertial modes can be thought of as waves trapped in a virtual potential, $U(r)=-k(r)^2$ \citep{lietal2003}, which for a frequency close to the maximum of $\kappa$ is similar to the harmonic oscillator potential."
+ Lf (lr)«0 waves can propagate. being evanescent in the regions where potential barriers exist.," If $U(r)<0$ waves can propagate, being evanescent in the regions where potential barriers exist."
+ Also. as in quantum mechanics. these trapped inertial waves can escape through the potential barriers and propagate on the other side. as p modes.," Also, as in quantum mechanics, these trapped inertial waves can escape through the potential barriers and propagate on the other side, as p modes."
+ The inertial modes are evanescent. between the inner radius and the first. Lindblad resonance ancl between the second Lindblad: resonance and the vertical onc., The inertial modes are evanescent between the inner radius and the first Lindblad resonance and between the second Lindblad resonance and the vertical one.
+ The, The
+Columns (14) and (15) are the rms noise values for the 1612 and 1720 MlIIz. transitions.,Columns (14) and (15) are the rms noise values for the 1612 and 1720 MHz transitions.
+ Blank fields in these two columns indicate that the rms noise values could not be estimated because of severe radio frequency interference (REI)., Blank fields in these two columns indicate that the rms noise values could not be estimated because of severe radio frequency interference (RFI).
+ Table 5 lists the parameters derived. Irom the OLL absorption lines., Table 5 lists the parameters derived from the OH absorption lines.
+ The values in this table are derived using the baseline-subtracted spectra., The values in this table are derived using the baseline-subtracted spectra.
+ Column (1) is the IRAS names of the galaxies., Column (1) is the IRAS names of the galaxies.
+ Column (2) is the (total on-source integration times., Column (2) is the total on-source integration times.
+ Column (3) gives the rms noise levels. and column (4) lists the heliocentric velocities at the center of the 1667.359 ΣΠ absorption lines.," Column (3) gives the rms noise levels, and column (4) lists the heliocentric velocities at the center of the 1667.359 MHz absorption lines."
+ Columns (5) and (9) are the full velocity widths at half maxinum of the 1667 and 1665 ΔΙ lines respectively. columns (6) and (10) are their peak optical depths. and columns (7) and (11) are their integrated. optical depths.," Columns (5) and (9) are the full velocity widths at half maximum of the 1667 and 1665 MHz lines respectively, columns (6) and (10) are their peak optical depths, and columns (7) and (11) are their integrated optical depths."
+ Column (8) is the column densities of the 1667 MlIz lines divided by the excitation temperature using the following equation [rom Turner(1973). ancl assuming a covering factor of unitv: where [τας is the integrated optical depth of the 1667 MIIz absorption feature given in column (7).," Column (8) is the column densities of the 1667 MHz lines divided by the excitation temperature using the following equation from \citet{Tur73}
+ and assuming a covering factor of unity: where $\int{\tau dv}$ is the integrated optical depth of the 1667 MHz absorption feature given in column (7)."
+ Column (12) lists the hyperfine ratios. obtainecl bv dividing the integrated optical depths of the (wo main Hines.," Column (12) lists the hyperfine ratios, obtained by dividing the integrated optical depths of the two main lines."
+ Columus (13) and (14) are (he rms noise values [or ihe 1612 and 1720 Ην (transitions., Columns (13) and (14) are the rms noise values for the 1612 and 1720 MHz transitions.
+ Blank fields in these two columns indicate that the rms noise values could not be estimated because of severe REI., Blank fields in these two columns indicate that the rms noise values could not be estimated because of severe RFI.
+ Table 6 lists relevant data for all sources with no detections in the OIL 18 em mainlines., Table 6 lists relevant data for all sources with no detections in the OH 18 cm mainlines.
+ Column (1) is the IIRLÀAS names of the galaxies., Column (1) is the IRAS names of the galaxies.
+ Column (2) lists the rms noise values of the OIL 18 cm main line non-delections., Column (2) lists the rms noise values of the OH 18 cm main line non-detections.
+" Column (3) is the logarithms of predicted OIL luminosities from equation 5. aud column (4) is (he logarithms of the maximum OIL Iuninosities for the 1667 Mllz main line determined using the equation (from Darling&Giovanelli 2000)): where D, is the luminosity distance. σ is the rms noise value listed in column (2). AV=150kis +. eis the speed of light. v,, is the rest frequency of the 1667 MITz transition. and z is the redshift."," Column (3) is the logarithms of predicted OH luminosities from equation 5, and column (4) is the logarithms of the maximum OH luminosities for the 1667 MHz main line determined using the equation (from \citealp{DG00}) ): where $D_{\rm L}$ is the luminosity distance, $\sigma$ is the rms noise value listed in column (2), $\Delta V= 150~ \rm km~s^{-1}$ , $c$ is the speed of light, $\nu_o$ is the rest frequency of the 1667 MHz transition, and $z$ is the redshift."
+ Columns (5) and (6) are the rms noise values of the 1612 and 1720 MIIz iransitions., Columns (5) and (6) are the rms noise values of the 1612 and 1720 MHz transitions.
+ Blank fields in this table indicate that the rms noise values could not be estimated becauseof severe REI., Blank fields in this table indicate that the rms noise values could not be estimated becauseof severe RFI.
+models.,models.
+ The models with more effective DM heating — 1.c. the AH4 model compared to the 100 GeV unboosted case — burn up their original adiabatically contracted DM the most quickly and enter the KH contraction phase the soonest., The models with more effective DM heating – i.e. the AH4 model compared to the $100$ GeV unboosted case – burn up their original adiabatically contracted DM the most quickly and enter the KH contraction phase the soonest.
+" This in turn leads to a smaller final mass, as feedback effects will shut off accretion sooner."," This in turn leads to a smaller final mass, as feedback effects will shut off accretion sooner."
+ When comparing the cases with different concentration parameter we notice that the DS final mass is an increasing function of c., When comparing the cases with different concentration parameter we notice that the DS final mass is an increasing function of $c$.
+" As previously explained an increase in the concentration parameter leads to à longer DS phase, and hence to more mass accreted."," As previously explained an increase in the concentration parameter leads to a longer DS phase, and hence to more mass accreted."
+ For the case of the unboosted 100 GeV WIMP with ο=3.5 the final mass is around 900.... whereas for the 47/4 case it is close to 700...," For the case of the unboosted $100$ GeV WIMP with $c=3.5$ the final mass is around $900 \msun$, whereas for the $AH4$ case it is close to $700 \msun$."
+ For e ranging from 2—5 the DS will have a mass in the 300.—100011. as it reaches the main sequence., For $c$ ranging from $2-5$ the DS will have a mass in the $800\msun-1000\msun$ as it reaches the main sequence.
+€ 155.16(= 08311557) is a radio galaxy at a redshift of 1.240 (Lawrence et al 1996) in the Pearson-Readheacl (1981. m 5 CGllz lux-density-limited: complete1 sample.,"4C+55.16 (= 0831+557) is a radio galaxy at a redshift of 0.240 (Lawrence et al 1996) in the Pearson-Readhead (1981, 1988) 5 GHz flux-density-limited complete sample."
+1 The radio source is compact and powerful (1.1.107οιWz tsed at 5 611)., The radio source is compact and powerful $1.1\times 10^{26}$ $^{-1}$ $^{-1}$ at 5 GHz).
+ 1n the high resolution radio mao obtaine rom ALERLIN and VLBL (Whsyborn et al 1985: Pearson Iteadhead: 1988). highly irregular. small-scale structures iive been resolved.," In the high resolution radio maps obtained from MERLIN and VLBI (Whyborn et al 1985; Pearson Readhead 1988), highly irregular, small-scale structures have been resolved."
+" The radio spectrum of t1ο compac core shows a turnover around 1 Cillz (Component Co in Whyborn ct al 1985). resembling the class of Ciga-hertz 'eaked Spectrum. (GPS) SOULCOS (ο,g.. O'Dea 998) whils he mini double-Iobe with a projected size of (54 kpe at a redshift oft) has a steep. low-frequency spectrum."," The radio spectrum of the compact core shows a turnover around 1 GHz (Component C in Whyborn et al 1985), resembling the class of Giga-hertz Peaked Spectrum (GPS) sources (e.g., O'Dea 1998) whilst the mini double-lobe with a projected size of (54 kpc at a redshift of 0.24) has a steep, low-frequency spectrum."
+ X sum of the two components resuts in a [la radio spectrum over the Gllz range., A sum of the two components results in a flat radio spectrum over the GHz range.
+ A blue continuum contributes 5t) per cent of the light at ((Aleckman et al 1983) and its optical emission-line spectrum shows a medium ionization state (e.e.. Lleckman ct al 1983 based on the OLLA5007 OLLA3727. ratio: Whyborn ct al 1985: OLHL]A5007/1L72~4. Lawrence et al 1996)," A blue continuum contributes $\sim 50$ per cent of the light at (Heckman et al 1983) and its optical emission-line spectrum shows a medium ionization state (e.g., Heckman et al 1983 based on the $\lambda 5007$ $\lambda 3727$ ratio; Whyborn et al 1985; $\lambda 5007$ $\beta\sim 4$, Lawrence et al 1996)."
+ X-ray emission was detected during the ROSAT All Sky Survey 08349|5534. Brinkmann et al 1995: Laurent-Aluehlei et al 1996: Bade et al 1998).," X-ray emission was detected during the ROSAT All Sky Survey J08349+5534, Brinkmann et al 1995; Laurent-Muehlei et al 1996; Bade et al 1998)."
+ The ROSAT papers assumed that the X-ray emission. originates in an active nucleus residing in 4€155.16., The ROSAT papers assumed that the X-ray emission originates in an active nucleus residing in 4C+55.16.
+ Llowever. as our imaging and spectral study using he ROSAT HII and SCA data show below. the observed X-rays appear to be dominated. by cluster emission surrounding the radio galaxy.," However, as our imaging and spectral study using the ROSAT HRI and ASCA data show below, the observed X-rays appear to be dominated by cluster emission surrounding the radio galaxy."
+ The optical CCD images of 4€|55.16 were obtained by Llutchines. Johnson Pyke (1988) using the Canada- Telescope (CELT) with D and ? filters.," The optical CCD images of 4C+55.16 were obtained by Hutchings, Johnson Pyke (1988) using the Canada-France-Hawaii Telescope (CFHT) with $B$ and $R$ filters."
+through restricting the free parameters to the minimum number (one) and understanding directly how uncertainties propagate through the model. that we retain predictive power. rather than simply reproduce the observations which we introduce as constraints.,"through restricting the free parameters to the minimum number (one) and understanding directly how uncertainties propagate through the model, that we retain predictive power, rather than simply reproduce the observations which we introduce as constraints."
+ Present day late tvpe galaxies come in a variety of sizes and shapes. but appear to obey well established scaling laws.," Present day late type galaxies come in a variety of sizes and shapes, but appear to obey well established scaling laws."
+ To a first approximation. disk galaxies can be described by four observational parameters: the total luminosity of the exponential light surface density. profile ancl the radial scale length. of this profile. the central density. ancl the asymptotic rotation velocity of the dark maer halo.," To a first approximation, disk galaxies can be described by four observational parameters: the total luminosity of the exponential light surface density profile and the radial scale length of this profile, the central density and the asymptotic rotation velocity of the dark matter halo."
+ The first two of these parameters relate directly to the barvonic component. and the last two refer mostlv to the dark matter.," The first two of these parameters relate directly to the baryonic component, and the last two refer mostly to the dark matter."
+ The total luminosity (mass?), The total luminosity (mass?)
+ is tightly. correlated with the asvmptotic rotation velocity through the Tully-Fisher relation (Lully Fisher 1977). the total luminosity also scales with the exponential disk scale length (Flores et al.," is tightly correlated with the asymptotic rotation velocity through the Tully-Fisher relation (Tully Fisher 1977), the total luminosity also scales with the exponential disk scale length (Flores et al."
+ 1993)., 1993).
+ The question addressed in this work is whether the observationally available information on present day ealactic structural parameters and their scaling laws is sullicient to introduce some constraints on the initial conditions of these galaxies. without having to impose some specific dark matter density. profile.," The question addressed in this work is whether the observationally available information on present day galactic structural parameters and their scaling laws is sufficient to introduce some constraints on the initial conditions of these galaxies, without having to impose some specific dark matter density profile."
+ To this end we apply a simple method. of mapping a set of initial conditions in terms of virialized halo structure. barvon fraction and angular momentum into à set of present day galactic structural parameters.," To this end we apply a simple method of mapping a set of initial conditions in terms of virialized halo structure, baryon fraction and angular momentum into a set of present day galactic structural parameters."
+ This methocl is essentially the one introduced by Blumenthal et al. (, This method is essentially the one introduced by Blumenthal et al. (
+1986) ane Flores et al. (,1986) and Flores et al. (
+1993). which we extend here to the case of a halo density profile which is treated as a mocel parameter.,"1993), which we extend here to the case of a halo density profile which is treated as a model parameter."
+ We restrict analysis to late-type and LSB galaxies. as only in these cases is the influence of dark matter substantial at all radii.," We restrict analysis to late-type and LSB galaxies, as only in these cases is the influence of dark matter substantial at all radii."
+ Analysis of earlier type galaxies is necessarily more sensitive to the evolution of the baryonic components., Analysis of earlier type galaxies is necessarily more sensitive to the evolution of the baryonic components.
+ We constrain all the initial conditions through the condition that the resulting ealactic structure follows the observational constraints. taking special note of the centra clensity of the dark halos. which can be derived. from. the recently published rotation curves of LSB galaxies of de Blok et al. (," We constrain all the initial conditions through the condition that the resulting galactic structure follows the observational constraints, taking special note of the central density of the dark halos, which can be derived from the recently published rotation curves of LSB galaxies of de Blok et al. ("
+1996).,1996).
+ Once this is done. we wish to investigate wha relations between the initial conditions result from having --mposed certain observational restrictions.," Once this is done, we wish to investigate what relations between the initial conditions result from having imposed certain observational restrictions."
+ Lt is possible vat when seen in terms of initial conditions. observationa laws such as the Tullv-Fisher relation or the approximate scaling of the disk mass with the square of the scale radius. appear natural.," It is possible that when seen in terms of initial conditions, observational laws such as the Tully-Fisher relation or the approximate scaling of the disk mass with the square of the scale radius, appear natural."
+ Adclitionally. one is interested to see to wha extent dillerent galaxy classes populate clillerent regions in the initial condition space.," Additionally, one is interested to see to what extent different galaxy classes populate different regions in the initial condition space."
+ As no specilie dark halo shape is imposed. we use present αν galactic structure o solve for the halo profile. which we find to diller from hat which results from the Cully selbeonsistent structure formation C'DAL simulations of Navarro. Frenk White (1996). heneclorth NEW.," As no specific dark halo shape is imposed, we use present day galactic structure to solve for the halo profile, which we find to differ from that which results from the fully self-consistent structure formation CDM simulations of Navarro, Frenk White (1996), henceforth NFW."
+ A comparison of this two dark 1alo profiles will be mace. in connection to the well studied dark halo of a particular gas rich cwarf galaxy. DDO 154.," A comparison of this two dark halo profiles will be made, in connection to the well studied dark halo of a particular gas rich dwarf galaxy, DDO 154."
+ Section (2) presents à dimensional analysis to. the ooblem. which is expanded. in. section(3) to a more complete galactic formation scenario. calibrated: using the data on LSB galaxies.," Section (2) presents a dimensional analysis to the problem, which is expanded in section(3) to a more complete galactic formation scenario, calibrated using the data on LSB galaxies."
+ In Section (4) we apply the previous framework to normal late tvpe galaxies. and in section (5) we present the conclusions of this work.," In Section (4) we apply the previous framework to normal late type galaxies, and in section (5) we present the conclusions of this work."
+ Although most of the material in this section can be found elsewhere. it is included here to give the reader an intuitive framework for the problem. and to present a zero order solution to it.," Although most of the material in this section can be found elsewhere, it is included here to give the reader an intuitive framework for the problem, and to present a zero order solution to it."
+" The three galactic structural paranicters which are most reacily accessible to observations in late type galaxies are M, the total disk: mass. Z2. the disk exponential scale radius and Vay. the maximunr rotation velocity."," The three galactic structural parameters which are most readily accessible to observations in late type galaxies are $M_d$, the total disk mass, $R_d$, the disk exponential scale radius and $V_M$ , the maximum rotation velocity."
+" AZ, is derivable from total luminosities once a 11, ratio has been assumed.", $M_d$ is derivable from total luminosities once a $M/L$ ratio has been assumed.
+ Uncertainties in the value of this ratio are generally of less than a factor of 2 for a given system., Uncertainties in the value of this ratio are generally of less than a factor of 2 for a given system.
+" These three quantities are seen to. be correlated hrough the Tullv-Fisher. CE-E) lav M,oxVy). and an approximate scaling of AZ;xλεν with a close to + and } close to 2."," These three quantities are seen to be correlated through the Tully-Fisher (T-F) law $M_d \propto V_M ^\alpha$ , and an approximate scaling of $M_d \propto R_d ^\beta$, with $\alpha$ close to 4 and $\beta$ close to 2."
+ In the following section we present data from he literature vielding a=3.5.92.33.," In the following section we present data from the literature yielding $\alpha =3.5, \beta = 2.33$."
+ Lt is perhaps easier to understand these correlations if one translates the observable quantities into initial conditions for the galaxy. or example. a total mass M. total potential energy. WW. xwvon fraction £'. and a total angular momentum. L.," It is perhaps easier to understand these correlations if one translates the observable quantities into initial conditions for the galaxy, for example, a total mass $M$ , total potential energy $W$, baryon fraction $F$, and a total angular momentum, $L$."
+ Once AM and P are fixed. the total barvon content is given by £A.," Once $M$ and $F$ are fixed, the total baryon content is given by $FM$."
+ lt is customary to refer not to L. but to the cimensionless quantity where C is Newton's constant. so it shall be treated thus in this work.," It is customary to refer not to $L$, but to the dimensionless quantity where $G$ is Newton's constant, so it shall be treated thus in this work."
+ On purely dimensional erouncs. we can establish the following relations between the initial conditions and the observed. parameters: at fixed f° and A the T-E law translates into a relation between AZ and M.," On purely dimensional grounds, we can establish the following relations between the initial conditions and the observed parameters: at fixed $F$ and $\lambda$ the T-F law translates into a relation between $M$ and $W$."
+ Substituting Wo for Z2; from the above relations one obtains: ic. for a galaxy population having constant £ and A. which could. represent a fixed. galaxv-tvpe sample. one expects a connection between the νι index. ancl the exponent in the Aes.Ay relation.," Substituting $W$ for $R_d$ from the above relations one obtains: i.e. for a galaxy population having constant $F$ and $\lambda$, which could represent a fixed galaxy-type sample, one expects a connection between the T-F index, and the exponent in the $M_d vs. R_d$ relation."
+ For a=4 equation (2) gives 3=2.quite in agreement with observations.," For $\alpha =4$ equation (2) gives $\beta =2$,quite in agreement with observations."
+ Using theabove relations between theinitial conditions ancl the final resultto translate the observational relations one obtains: Eliminating AZ and V. from the above two relations for j—onf(n 2) we get:, Using theabove relations between theinitial conditions and the final resultto translate the observational relations one obtains: Eliminating $M$ and $W$ from the above two relations for $\beta=\alpha /(\alpha -2) $ we get:
+that is also reproduced by our simulation results.,that is also reproduced by our simulation results.
+" On the other hand, the correlation between present-day stellar mass and star formation rate seen in our simulations seems to be not nearly as well-defined as in the observational data."," On the other hand, the correlation between present-day stellar mass and star formation rate seen in our simulations seems to be not nearly as well-defined as in the observational data."
+ This is related to the fact that we do not find a good correlation between the present stellar and gaseous masses; many simulated satellites have comparable stellar masses but differ in their gas fractions by huge factors., This is related to the fact that we do not find a good correlation between the present stellar and gaseous masses; many simulated satellites have comparable stellar masses but differ in their gas fractions by huge factors.
+" Gas-rich and completely gas-depleted satellites coexist in the same total and stellar mass regime, rendering a tight correlation with the star formation rate unlikely."," Gas-rich and completely gas-depleted satellites coexist in the same total and stellar mass regime, rendering a tight correlation with the star formation rate unlikely."
+ But perhaps the most significant discrepancy between the simulation results and observations lies in the inferred mass-to-light ratios., But perhaps the most significant discrepancy between the simulation results and observations lies in the inferred mass-to-light ratios.
+ The mass-to-light ratios of the simulated galaxies are off by about a factor of 5 when compared at face value to the observational estimates., The mass-to-light ratios of the simulated galaxies are off by about a factor of $5$ when compared at face value to the observational estimates.
+" This means that they are either too massive, or too faint for their mass."," This means that they are either too massive, or too faint for their mass."
+ The discrepancy could also be caused by a systematic underestimate of the total satellite masses in the observations., The discrepancy could also be caused by a systematic underestimate of the total satellite masses in the observations.
+" Due to the difficulty of reliably determining the ‘outer edge’ of the dark matter halo of an orbiting satellite, this possibility cannot be easily excluded."," Due to the difficulty of reliably determining the `outer edge' of the dark matter halo of an orbiting satellite, this possibility cannot be easily excluded."
+" In summary, we find that the current generation of cosmological hydrodynamic simulations is able to explain many properties of the observed satellite population surprisingly well."," In summary, we find that the current generation of cosmological hydrodynamic simulations is able to explain many properties of the observed satellite population surprisingly well."
+ We have shown that different feedback, We have shown that different feedback
+Μη—0. relation along with other dynamical black hole mass measurements compiled by Gültekinetal. in Figure 12..,$M_{BH} - \sigma_{*}$ relation along with other dynamical black hole mass measurements compiled by \citet{Gult} in Figure \ref{Msigma}.
+ We plot separate points for our black hole mass measurements under two different assumptions., We plot separate points for our black hole mass measurements under two different assumptions.
+" Because there is some ambiguity about the appropriate way to measure the stellar velocity dispersion in the bulge in a system such as NGC 6240, we plot our measured black hole mass with several different o, values from the literature as well as with our measured o, very close to the black hole."," Because there is some ambiguity about the appropriate way to measure the stellar velocity dispersion in the bulge in a system such as NGC 6240, we plot our measured black hole mass with several different $\sigma_{*}$ values from the literature as well as with our measured $\sigma_{*}$ very close to the black hole."
+" NGC 6240 appears to lie well within the scatter of the Mpgg—σε relation, which may suggest that the black hole mass and the bulge velocity dispersion grow simultaneously and at similar rates during a major merger."," NGC 6240 appears to lie well within the scatter of the $M_{BH} - \sigma_{*}$ relation, which may suggest that the black hole mass and the bulge velocity dispersion grow simultaneously and at similar rates during a major merger."
+" NGC 6240 is a late-stage merger, but these data could also suggest that a system doesn't evolve along the Μη—o, relation until the very end stages of a merger, perhaps during nuclear coalescense."," NGC 6240 is a late-stage merger, but these data could also suggest that a system doesn't evolve along the $M_{BH} - \sigma_{*}$ relation until the very end stages of a merger, perhaps during nuclear coalescense."
+" To test this, more systems at this and later stages of merging will have to be studied."," To test this, more systems at this and later stages of merging will have to be studied."
+" This is interesting to compare to galaxy merger simulations; for example, Dasyraetal.(2006) conclude that, if the accretion efficiency stays constant, the Mpg—σι relation should be maintained from midway between the first encounter and coalescence through to final relaxation."," This is interesting to compare to galaxy merger simulations; for example, \citet{Dasyra06} conclude that, if the accretion efficiency stays constant, the $M_{BH} - \sigma_{*}$ relation should be maintained from midway between the first encounter and coalescence through to final relaxation."
+ New higher-resolution simulations are being performed by Stickley et al. , New higher-resolution simulations are being performed by Stickley et al. (
+which show that the stellar velocity dispersion (instays prep)low in the nuclear regions much longer than in the rest of the system; results will follow young and old stellar populations separately to help understand the dynamics measured from specific lines.,in prep) which show that the stellar velocity dispersion stays low in the nuclear regions much longer than in the rest of the system; results will follow young and old stellar populations separately to help understand the dynamics measured from specific lines.
+" Black holes are an important ingredient in galaxy evolution, as seen by the tight correlations between black hole mass and host galaxy properties."," Black holes are an important ingredient in galaxy evolution, as seen by the tight correlations between black hole mass and host galaxy properties."
+" Hydrodynamic simulations of galaxy mergers have suggested that gas-rich mergers, which can drive galaxy evolution, can provide fuel for black hole accretion."," Hydrodynamic simulations of galaxy mergers have suggested that gas-rich mergers, which can drive galaxy evolution, can provide fuel for black hole accretion."
+" In turn, these accreting black holes can radiate enough energy to affect the surrounding galaxy."," In turn, these accreting black holes can radiate enough energy to affect the surrounding galaxy."
+" In light of these discussions, it is particularly interesting to study systems that appear to bein the middle of such an evolutionary event."," In light of these discussions, it is particularly interesting to study systems that appear to be in the middle of such an evolutionary event."
+ NGC 6240 represents a nearby galaxy such as this: a merging gas-rich system with two actively accreting black holes., NGC 6240 represents a nearby galaxy such as this: a merging gas-rich system with two actively accreting black holes.
+" However, such systems are notoriously difficult to study because of their unrelaxed dynamics and their dusty cores."," However, such systems are notoriously difficult to study because of their unrelaxed dynamics and their dusty cores."
+" We have presented high-spatial resolution kinematics within the sphere of influence of the black hole in the south nucleus of NGC 6240, a nearby late-stage merger, made possible by laser guide star adaptive optics."," We have presented high-spatial resolution kinematics within the sphere of influence of the black hole in the south nucleus of NGC 6240, a nearby late-stage merger, made possible by laser guide star adaptive optics."
+" For this test case, we have explored two possible methods for measuring black hole mass in such a system and compared their results and assumptions."," For this test case, we have explored two possible methods for measuring black hole mass in such a system and compared their results and assumptions."
+" We have utilized the JAM modeling technique made public by Cappellari(2008), demonstrating that it is possible to complete such an analysis on a late-stage merger."," We have utilized the JAM modeling technique made public by \citet{JAM}, demonstrating that it is possible to complete such an analysis on a late-stage merger."
+" We point out that this technique likely overestimates the black hole mass by assuming that all measured velocity dispersion is due to a relaxed system (ignoring intervening unrelaxed material, e.g. tidal tails), and therefore report an upper limit of 2.0+10? M."," We point out that this technique likely overestimates the black hole mass by assuming that all measured velocity dispersion is due to a relaxed system (ignoring intervening unrelaxed material, e.g. tidal tails), and therefore report an upper limit of $2.0 \pm 0.2 \times 10^9 M_{\sun}$ ."
+" To provide a lower limit to the black hole mass, we explore the opposite assumption: that all velocity dispersion is caused by intervening material, and that the young stars sit in a thin disk around the black hole."," To provide a lower limit to the black hole mass, we explore the opposite assumption: that all velocity dispersion is caused by intervening material, and that the young stars sit in a thin disk around the black hole."
+ This model of Keplerian rotation around a black hole plus smooth spheroidal mass profile suggests a black hole which is at least 8.7+0.3x10°Mo., This model of Keplerian rotation around a black hole plus smooth spheroidal mass profile suggests a black hole which is at least $8.7 \pm 0.3 \times 10^8 M_{\sun}$.
+" We find that these two techniques provide measurements that are roughly consistent, and that follow the biases implied by their intrinsic assumptions."," We find that these two techniques provide measurements that are roughly consistent, and that follow the biases implied by their intrinsic assumptions."
+" 'To determine which set of assumptions is more reliable would require a detailed study of high-resolution galaxy simulations, which is beyond the scope of this paper."," To determine which set of assumptions is more reliable would require a detailed study of high-resolution galaxy simulations, which is beyond the scope of this paper."
+" Still we are encouraged that, to within a factor of about two, both measurements agree."," Still we are encouraged that, to within a factor of about two, both measurements agree."
+" While we cannot make generalizations on how all black hole-galaxy coevolution must proceed, it seems that in this case, the black hole and the host galaxy parameters grow together along the Mgy—o. scaling relation, instead of one preceeding the other."," While we cannot make generalizations on how all black hole-galaxy coevolution must proceed, it seems that in this case, the black hole and the host galaxy parameters grow together along the $M_{BH} - \sigma_{*}$ scaling relation, instead of one preceeding the other."
+" We are beginning an observing campaign to study other local merging galaxies using the same techniques, to investigate a larger sample size."," We are beginning an observing campaign to study other local merging galaxies using the same techniques, to investigate a larger sample size."
+" We enthusiastically thank the staff of the W. M. Keck Observatory and its AO team, for their dedication and hard work."," We enthusiastically thank the staff of the W. M. Keck Observatory and its AO team, for their dedication and hard work."
+" Data presented herein were obtained at the W. M. Keck Observatory, which is operated as a scientific partnership among the California Institute of Technology, the University of California, and the"," Data presented herein were obtained at the W. M. Keck Observatory, which is operated as a scientific partnership among the California Institute of Technology, the University of California, and the"
+Even with these limitations. we believe that the tov model provides strong support towards internal heating due to early star formation. allowing the gas {ο more easily stripped due to the larger scale heights and lower densitv of warm and hot eas.,"Even with these limitations, we believe that the toy model provides strong support towards internal heating due to early star formation, allowing the gas to more easily stripped due to the larger scale heights and lower density of warm and hot gas."
+ In particular we were able to reproduce the fraction of dwarf galaxies (hat retain. assuming that the dwarl ealaxies infall al a redshift of z=10 consistent with Luxetal.(2010)., In particular we were able to reproduce the fraction of dwarf galaxies that retain assuming that the dwarf galaxies infall at a redshift of $z=10$ consistent with \citet{Lux2010}.
+. The covering fraction of warm gas stripped from the dwarls is expected to be low. with (he contrails containing approximately (he same amount of warn gas as exists inside the cold streams that cover ~2% of the projected area of galaxies (Faucher-Giguere&eres2010).," The covering fraction of warm gas stripped from the dwarfs is expected to be low, with the contrails containing approximately the same amount of warm gas as exists inside the cold streams that cover $\sim2\%$ of the projected area of galaxies \citep{Faucher-Giguere2010}."
+. This gas breaking away in small clamps of warm material will form a warm component of the hot halo which is stabilised against heat conduction from the halo bv cooling before falling onto the Galaxy as a warm rain.," This gas breaking away in small clumps of warm material will form a warm component of the hot halo which is stabilised against heat conduction from the halo by cooling \citep{Vieser2007,Vieser2007a} before falling onto the Galaxy as a warm rain."
+ Even with the majority of cwarl galaxies entering al early redshifts. the large Gimescale requirecl to complete large pericentre orbits (13 Gvr for ry=270 kpc) means that gas from any late infalling as those from z=3 or 2—=I. will only be falling onto the disk of the Galaxy today.," Even with the majority of dwarf galaxies entering at early redshifts, the large timescale required to complete large pericentre orbits $13$ Gyr for $r_p=270$ kpc) means that gas from any late infalling dwarfs---such as those from $z=3$ or $z=1$ —will only be falling onto the disk of the Galaxy today."
+ A search lor this gas in Πα around local galaxies is now being undertaken using the \larvlancl- Tunable Filter (Veilleuxetal.2010) and the Grantecan Osiris Tunable Filter (Cepaetal.2003) where the expected emission measure of the gas (E0.1 ° pe) may be just large enough to be detectable., A search for this gas in $\alpha$ around local galaxies is now being undertaken using the Maryland-Magellan Tunable Filter \citep{Veilleux2010} and the Grantecan Osiris Tunable Filter \citep{Cepa2003} where the expected emission measure of the gas $\la 0.1$ $^{-6}$ pc) may be just large enough to be detectable.
+ ALN. is supported by an Australian Postgraduate Award., M.N. is supported by an Australian Postgraduate Award.
+ J.D.-II. is supported by a Federation Fellowship from the Australian Research Council., J.B.-H. is supported by a Federation Fellowship from the Australian Research Council.
+ The authors wish to thank Svivain Veilleux. Ken Freeman. Brent Tully Doug Lin [ον their help and comments throughout this work.," The authors wish to thank Sylvain Veilleux, Ken Freeman, Brent Tully Doug Lin for their help and comments throughout this work."
+padiative treatment.,radiative treatment.
+ The colors refer to the Ilo mass raction., The colors refer to the $_2$ mass fraction.
+ Molecular hydrogen is dissociated in the presence of ionizing radiation (mostly in the Lvman-Werner band). and this effect. increases. with decreasing redshift. due to he higher number of available photons.," Molecular hydrogen is dissociated in the presence of ionizing radiation (mostly in the Lyman-Werner band), and this effect increases with decreasing redshift, due to the higher number of available photons."
+ As a consequence. his lowers the star formation process. since El» provides he Largest part of the sub-107 cooling.," As a consequence, this lowers the star formation process, since $_2$ provides the largest part of the $10^4\rm \, K$ cooling."
+ 1t is evident rom the plots that without radiativeIx. feedback Ll. can reach levels 107 and boost carly star formation., It is evident from the plots that without radiative feedback $_2$ can reach levels $\gtrsim 10^{-2}$ and boost early star formation.
+ In presence of radiation from stellar. sources. molecular ractions decrease of several orders. of magnitude down osLo 1o," In presence of radiation from stellar sources, molecular fractions decrease of several orders of magnitude down to $\lesssim 10^{-4}$."
+Cphis elect is stronger for low-density gas. where molecules re-formi more slowly.," This effect is stronger for low-density gas, where molecules re-form more slowly."
+ Such behaviour is clear from the both cases shown. at redshift z=10.61 and 2=6.14.," Such behaviour is clear from the both cases shown, at redshift $z=10.61$ and $z=6.14$."
+ In the latter case. a wider temperature spread. due to thermal heating of the infalling gas at low density. appears as a consequence of the ongoing structure formation.," In the latter case, a wider temperature spread, due to thermal heating of the infalling gas at low density, appears as a consequence of the ongoing structure formation."
+" Additionally. also SN-heated gas at 10)10""I is evident with extremely low molecular fractions (5.10 y"," Additionally, also SN-heated gas at $10^5-10^6\,\rm K$ is evident with extremely low molecular fractions $\lesssim 10^{-16}$ )."
+ The lack of very cold gas in the more realistic case with both chemistry and ICE suggests that low-temperature cooling bv metals (e.g.7) could be an important mechanism to sustain star formation after the first generation of stars ().., The lack of very cold gas in the more realistic case with both chemistry and RT suggests that low-temperature cooling by metals \cite[e.g.][]{Maio2007} could be an important mechanism to sustain star formation after the first generation of stars \cite[][]{Maio2010}.
+ To better underline the impacts of radiative elfects on early chemistry. we also show in Fig.," To better underline the impacts of radiative effects on early chemistry, we also show in Fig."
+" 9 the redshift evolution of the mean and maximum number fractions of the two main molecules. E» and LD. and of the basic species for the different channels of molecular formation. and LL,."," \ref{fig:comp_evolution} the redshift evolution of the mean and maximum number fractions of the two main molecules, $_2$ and HD, and of the basic species for the different channels of molecular formation, $^-$ and $_2^+$."
+" As mentioned previously. after carly star formation sets in. around redshift +~12. the LH» fractions drops by several orders of magnitude (rom average values of ~10.7 down lo 10 1""). as ionizing radiation starts to propagate in"," As mentioned previously, after early star formation sets in, around redshift $z \sim 12$, the $_2$ fractions drops by several orders of magnitude (from average values of $\sim 10^{-5}$ down to $\lesssim 10^{-10}$ ), as ionizing radiation starts to propagate in"
+" (—15”) UL,=7Thkms!Mpe.1) 15"" pau (Z4<30). (Tq<30k) ","$L_{\rm IR} > 10^{11}
+L_{\sun}$ $\sim
+8-15\arcsec$ $H_o=75\kps
+\Mpc^{-1}$ $15\arcsec$ $\mu$ $\Td \la 30 \K$ $\Td \la 30 \K$ "
+brightuess plausibly arises from a reductiou in the coutributiou to the system light from a scattered component that depeucds on binary phase. not on AN. Perhaps the bigges surprise of tlie recent. monitoring has been the [act iat the star has uot. faded to anywhere Hear {ie low brightuess levels predicted. by extat| moclels.,"brightness plausibly arises from a reduction in the contribution to the system light from a scattered component that depends on binary phase, not on $\Delta$ X. Perhaps the biggest surprise of the recent monitoring has been the fact that the star has not faded to anywhere near the low brightness levels predicted by extant models."
+ Also. as the 2000-2010 data clearly slow. we continue to see substantial variatiou ou tle orbital period with a peak brightuess o “abou [= L6.70.," Also, as the 2009-2010 data clearly show, we continue to see substantial variation on the orbital period with a peak brightness of about I = 16.70."
+ As Figs., As Figs.
+"Lo and 5. show. this variatiolint e scatteed liel component has appareutly existed as a ""Που to the system brightuess marked |w the inflection point at AN=+41 iu the ight ¢uve lor many years (see Secton 3.1).","\ref{fig3} and \ref{fig2a} show, this variation in the scattered light component has apparently existed as a “floor"" to the system brightness marked by the inflection point at $\Delta$ X=+1 in the light curve for many years (see Section 3.1)."
+ If auvting. ie “floor” was a a brighter level iu the past (i.e. pre-2005) than it is today.," If anything, the “floor"" was at a brighter level in the past (i.e. pre-2005) than it is today."
+" The behavior o“this οςonuponent is uo consistent. with eiher a scatte‘ine halo arouud the individal stars or fo""wald scaδις: (at leas the simplest versicis of such models). since they would precict a dependence of the scatte'ec Πο on AN. not on orjtal pliase."," The behavior of this component is not consistent with either a scattering halo around the individual stars or forward scattering (at least the simplest versions of such models), since they would predict a dependence of the scattered light on $\Delta$ X, not on orbital phase."
+ The risiug brightuess of tle ---uUlection poiitat AXN= +1 of the ligh Curves seen on FigoOs., The rising brightness of the inflection point at $\Delta$ X=+1 of the light curves seen on Figs.
+ 1. and 5 shows that the scattered ligh component is not a ftnetion of AN. Iideed. wha the recent cata show is that alter all dlrec starlieht is removecl fi‘oun the syste here is a persistent laree-amplitucle periodic variation tlat 'emialus.," \ref{fig3} and \ref{fig2a} shows that the scattered light component is not a function of $\Delta$ X. Indeed, what the recent data show is that after all direct starlight is removed from the system, there is a persistent large-amplitude periodic variation that remains."
+ Iu Fig., In Fig.
+ο) d Wes how the plas ata [rou 2009-2010. which illustrates that variation. at da sixtli-order polsnomia fit. (," \ref{fig3a} we show the phased data from 2009-2010, which illustrates that variation, and a sixth-order polynomial fit. ("
+We do 1 other to fit the small. but real. reversal in briglitness iear yhase 0.5 which is due to star B. since is of 1O consequence to tlius discussion.),"We do not bother to fit the small, but real, reversal in brightness near phase 0.5 which is due to star B, since it is of no consequence to this discussion.)"
+ This rather stbstautial contribution to he system light as Clea‘ly been present siuce 1995 and may have been even brighter in the past., This rather substantial contribution to the system light has clearly been present since 1995 and may have been even brighter in the past.
+ It is sullicienly wieht tlat it contributes at about the level near a binary phase of 1.0 even when star A is ally visible., It is sufficiently bright that it contributes at about the level near a binary phase of 1.0 even when star A is fully visible.
+ It is clear that once star A is fully obscurecl we see a good deal of scattered Lielt (AgoletWinnetal.2006:Herbst2008:Silvia&Agol2008) but he geomet‘y ol the scattering has remained uncertain.," It is clear that once star A is fully obscured we see a good deal of scattered light \citep{agol04,Ham05,w06,Her08,sa08} but the geometry of the scattering has remained uncertain."
+ There are at least [our possile scenarios. aud they are not mutually exchsive.," There are at least four possible scenarios, and they are not mutually exclusive."
+ There may be a seattering Lalo around. eacl star arisitD>oO ron matte “within the magnetos2169/60. perhaps coulined to accretion columus (Winetal.2006).," There may be a scattering halo around each star arising from matter within the magnetosphere, perhaps confined to accretion columns \citep{w06}."
+. There may be lorward scattering ol light (rom each star reflectiug off the presumed uzzy edge of he occultiug clisk (or ring) (Silvia&Agol2008)., There may be forward scattering of light from each star reflecting off the presumed fuzzy edge of the occulting disk (or ring) \citep{sa08}.
+. There must be some back scattered light [ro1i star A conug to us trom the [ar side of the disk to account for the reflection eflec observed in its velocity during eclipse (Herbsetal. 2008)., There must be some back scattered light from star A coming to us from the far side of the disk to account for the reflection effect observed in its velocity during eclipse \citep{Her08}.
+. Aud. there could be general scattering from mo'e," And, there could be general scattering from more"
+falling into it.,falling into it.
+ It would thus be unlikely to observe them at these very particular phases., It would thus be unlikely to observe them at these very particular phases.
+ Nevertheless such examples could be found in WASP-12. 18 and 19 (eg. ?)).," Nevertheless such examples could be found in WASP-12, 18 and 19 (eg. \citet{Hellier:2011p10960}) )."
+ ? deseribe how planets initially placed on mildly inclined or aligned orbits. are less likely to in-fall and more likely to survive until observed.," \citet{Matsumura:2010p8927} describe how planets initially placed on mildly inclined or aligned orbits, are less likely to in-fall and more likely to survive until observed."
+ Nevertheless. in most cases tidal realignment corresponds to the disappearance of the planet.," Nevertheless, in most cases tidal realignment corresponds to the disappearance of the planet."
+ If retrograde planets plunge into their star as they tidally realign. a decreasing. number of hot Jupiters should be observed with time.," If retrograde planets plunge into their star as they tidally realign, a decreasing number of hot Jupiters should be observed with time."
+ No such decreasing trend can be found when considering all the hot Jupiters presented in the literature around stars in the mass range considered for this paper., No such decreasing trend can be found when considering all the hot Jupiters presented in the literature around stars in the mass range considered for this paper.
+ This is at odds with evidence of a trend between semimajor axis and stellar age showing a lack of very short orbits around older stars as presented by ? 2.. ?.. ? (??)..," This is at odds with evidence of a trend between semimajor axis and stellar age showing a lack of very short orbits around older stars as presented by \citet{Jackson:2009p13324} \ref{tab:Param}, \citet{Winn:2010p7311}, \citet{Lai:2011p14503} \citep{Hut:1981p2945,Barker:2009p11693}."
+ ? ? | , \citet{Ogilvie:2004p11737} \citet{Lai:2011p14503} \ref{fig:tracks} 
+and 327. a sound speec of 550-700 |. a galaxy velocity of 1500 and assundug an iutracluster deusitv of »=10+ 7. we derive Alfvéunu speeds between 300  and 600 1 aud a magnetic field strength of ~ 1-3 pC at the location of NGC 1569.,"and $32^{\circ}$, a sound speed of $550$ $700$ $^{-1}$, a galaxy velocity of $1500$ $^{-1}$, and assuming an intracluster density of $n=10^{-4}$ $^{-3}$, we derive Alfvénn speeds between $300$ $^{-1}$ and $600$ ${-1}$ and a magnetic field strength of $\sim 1$ $3$ $\mu$ G at the location of NGC 4569."
+ The derived range of Alfvéóun speeds is thus comparable to that of the sound speed., The derived range of Alfvénn speeds is thus comparable to that of the sound speed.
+ The derived maguetic Ποιά streneth is consistent with typical magnetic field streneths in ealaxy clusters (Ferrari ot al., The derived magnetic field strength is consistent with typical magnetic field strengths in galaxy clusters (Ferrari et al.
+ 2008)., 2008).
+ NGC [388's projected location within the Vireo cluster is close to AD 87 (1.37.=385 kpc)., NGC 4388's projected location within the Virgo cluster is close to M 87 $1.3^{\circ}=385$ kpc).
+ Yoshida et al. (20013) , Yoshida et al. \cite{yoshida}) )
+discovered anu Πα plume which extends up to 35 kpc northeast from the galaxy center., discovered an $\alpha$ plume which extends up to $35$ kpc northeast from the galaxy center.
+ Based on a Effelsbere 1001: telescope detection of the Πο pluie. Volliner Tucltimeier (2003) made dynamical models aud concludes that rami pressure stripping can account for the observe extraplanar Ho and eiissiou.," Based on a Effelsberg 100m telescope detection of the $\alpha$ plume, Vollmer Huchtmeier (2003) made dynamical models and concluded that ram pressure stripping can account for the observed extraplanar $\alpha$ and emission."
+ Oosterloo van Corkou (2005) observed NGC [388 with the WSRT aud discoverec zu tail which extends more han LOO kpe to the north., Oosterloo van Gorkom (2005) observed NGC 4388 with the WSRT and discovered an tail which extends more than $100$ kpc to the north.
+ We clearly detect X-ray cussion associated to the and Πα tail of NGC 1388., We clearly detect X-ray emission associated to the and $\alpha$ tail of NGC 4388.
+ We can only attribute the observed: X-ray cluission fo he first. ~5 (~25 kpc) of the tail which is about 5 times louger., We can only attribute the observed X-ray emission to the first $\sim 5'$ $(\sim 25$ kpc) of the tail which is about 5 times longer.
+ This is because the projected distribution approaches AL 86 and the halo of M 86 dominates the N-ray cuiission in this region., This is because the projected distribution approaches M 86 and the halo of M 86 dominates the X-ray emission in this region.
+ As for NGC 1569 we adjusted a Mach cone to the X- and cussion distribution with the direction being based on the dynamical model (Vollmer IIuchtiieier 2003. Volhuer 2009).," As for NGC 4569 we adjusted a Mach cone to the X-ray and emission distribution with the direction being based on the dynamical model (Vollmer Huchtmeier 2003, Vollmer 2009)."
+ The apparent Mach cone angle is ac26° (Fie. 8))., The apparent Mach cone angle is $\alpha \sim 26^{\circ}$ (Fig. \ref{fig:n4388m}) ).
+ With a normalized 3D model velocity vector (0.09.-0.75.0.66) (Vollmer 2009) the true Mach cone anele is 207.," With a normalized 3D model velocity vector (0.09,-0.75,0.66) (Vollmer 2009) the true Mach cone angle is $20^{\circ}$."
+ This corresponds toa Mach nuiuber of 2.9 and a Reealaxy velocity of ~ 1600-2000 | (with a sound speed of 550-700. ly which is consistent with the estimated galaxy. velocity based on the dyvuamicalmodel (1700 +. Vollmer 2009).," This corresponds to a Mach number of $2.9$ and a galaxy velocity of $\sim 1600$ $2000$ $^{-1}$ (with a sound speed of $550$ $700$ $^{-1}$ ), which is consistent with the estimated galaxy velocity based on the dynamicalmodel $1700$ $^{-1}$, Vollmer 2009)."
+" To obtain an Alfvémnic Mach umuber of 2.9 with c4=1700 + . a lnagenetic fiek streneth of ~23 (iG is needed with au intrachister density of 7=10Sem 3,"," To obtain an Alfvénnic Mach number of $2.9$ with $v_{\rm gal}=1700$ $^{-1}$ , a magnetic field strength of $\sim 3$ $\mu$ G is needed with an intracluster density of $n=10^{-4}$ $^{-3}$."
+" Within the N-rav tail we incasure a temperature of OST keV (10"" K) and a deusity of n4)010Bom 7,"," Within the X-ray tail we measure a temperature of $0.87$ keV $10^7$ K) and a density of $n \sim \eta^{-0.5}\,10^{-3}$ $^{-3}$."
+ The thermal pressure is thus pai;=1.1410.DP 7. whereas rau pressure is about Ls<10£2 ," The thermal pressure is thus $p_{\rm th}=1.4 \times 10^{-12}$ $^{-2}$, whereas ram pressure is about $4$ $8 \times 10^{-12}$ $^{-2}$."
+As in the halo of NGC 1569. the thermal xessure based ou the N-rav observations is siguificautly lower han the estimated ram pressure acting on NGC 1388.," As in the halo of NGC 4569, the thermal pressure based on the X-ray observations is significantly lower than the estimated ram pressure acting on NGC 4388."
+ We made an attempt to calculate the total mass ofthe hot N-rav gas within the whole tail., We made an attempt to calculate the total mass of the hot X-ray gas within the whole tail.
+ Unfortunately. this is uot directly possible due to the dominating X-rav halo of M 86.," Unfortunately, this is not directly possible due to the dominating X-ray halo of M 86."
+ Several assuniptious have to be made: first. a constant surface-briglhtuess along thewhole tail is assumed.," Several assumptions have to be made: first, a constant surface-brightness along thewhole tail is assumed."
+ Next. the flux from the area of detectable N-rav tail is rescaled to the total volume of the tail.," Next, the flux from the area of detectable X-ray tail is rescaled to the total volume of the tail."
+ Since the tailis ~3 times longer than the N-ray. tail. this leads toa 3 times larger vole.," Since the tail is $\sim 3$ times longer than the X-ray tail, this leads to a 3 times larger volume."
+" Given the temperature from the spectral ft. the density of the hot gas isa ~yp""7107 ."," Given the temperature from the spectral fit, the density of the hot gas is $n \sim \eta^{-0.5} 10^{-3}$ $^{-3}$."
+ A total eas nass of ~2&105 AL. is found in the part of the N-vay tail which is uot dominated by M 86s halo (south o the declination 12Ls00)., A total gas mass of $\sim 2 \times 10^8$ $_{\odot}$ is found in the part of the X-ray tail which is not dominated by M 86's halo (south to the declination $12\ 48\ 00$ ).
+ Since our volume estimate is on the high side and since we expect a volume filling actor sunaller than l1. a more realistic mass estimate of he X-ray tail is ~10° AL...," Since our volume estimate is on the high side and since we expect a volume filling factor smaller than 1, a more realistic mass estimate of the X-ray tail is $\sim 10^8$ $_{\odot}$."
+ Within a 2 times longer tail. here might be a few 105 AL. of hot eas associated with he entire tail of NCC [3ss.," Within a 3 times longer tail, there might be a few $10^{8}$ $_{\odot}$ of hot gas associated with the entire tail of NGC 4388."
+ This total mass of hot eas is comparable to the total mass of atomic bydrogen in he tail (3.1«105 ALL: Oosterloo vau Corkom 2005)., This total mass of hot gas is comparable to the total mass of atomic hydrogen in the tail $3.4 \times 10^{8}$ $_{\odot}$; Oosterloo van Gorkom 2005).
+ The total gas mass of the tail (~6«105 AL.) ds much ower than the expected stripped gas mass based on the deficiency <LOPAL: : Voliner Thichtiucier 2007)., The total gas mass of the tail $\sim 6 \times 10^{8}$ $_{\odot}$ ) is much lower than the expected stripped gas mass based on the deficiency $\times$ $^9$ $_{\sun}$; Vollmer Huchtmeier 2007).
+ Possible reasons for this discrepancy are strong ISALICAL musing with a strong decrease of the eas deusity or star formation within the eas tail., Possible reasons for this discrepancy are strong ISM-ICM mixing with a strong decrease of the gas density or star formation within the gas tail.
+" Based on the recoustructed star formation history frou, deep optica spectra Pappalardo et al. (", Based on the reconstructed star formation history from deep optical spectra Pappalardo et al. (
+2010) found that the strippec eas left the galactic disk 200 Myr ago.,2010) found that the stripped gas left the galactic disk $\sim 200$ Myr ago.
+ The eas which has been stripped from the location where the optica spectra were taken is presunmablv now in middle of the tail., The gas which has been stripped from the location where the optical spectra were taken is presumably now in middle of the tail.
+ The eas within the tail thus traveled with a velocity iu the sky plane of ~300 +., The gas within the tail thus traveled with a velocity in the sky plane of $\sim 300$ $^{-1}$.
+ Witla radia velocity of ~300 | the total outflow velocity of the stripped gas is ~ 100-150 +., With a radial velocity of $\sim 300$ $^{-1}$ the total outflow velocity of the stripped gas is $\sim 400$ $450$ $^{-1}$.
+ This 3D coufiguratiou Is consisten with the Volliner Huchbtiueier (2003) model., This 3D configuration is consistent with the Vollmer Huchtmeier (2003) model.
+ The stripping timescale is reasonable in a scenario where NGC [385 as been stripped by a recent passage through the Virgo cluster core. close to AL δε (Volluer 2009).," The stripping timescale is reasonable in a scenario where NGC 4388 has been stripped by a recent passage through the Virgo cluster core, close to M 87 (Vollmer 2009)."
+ Therefore. we think that the tail PaACCs he eutire stripping process.," Therefore, we think that the tail traces the entire stripping process."
+ Caven the simall number of regions within the tail (Yoshida et al., Given the small number of regions within the tail (Yoshida et al.
+ 2001). gas loss via star ornuation seclus not to be significant.," 2004), gas loss via star formation seems not to be significant."
+ Vollmer Huchtineier (2007) sugeested[mI] an alternative explanation for the missing stripped eas: spiral galaxies catering a cluster may rapidly lose the atomic gas which is located bevoud the optical radius., Vollmer Huchtmeier (2007) suggested an alternative explanation for the missing stripped gas: spiral galaxies entering a cluster may rapidly lose the atomic gas which is located beyond the optical radius.
+ This gas is (i) less bound than the eas in the immer disk aud (i) no longer stirred by strong interstellar turbulence aud can casily be heated aud evaporated., This gas is (i) less bound than the gas in the inner disk and (ii) no longer stirred by strong interstellar turbulence and can easily be heated and evaporated.
+ The stripping of the outer gas disk mieht happen before the ealaxy euters the clustercore., The stripping of the outer gas disk might happen before the galaxy enters the clustercore.
+ Iu this case the galaxy already: shows au initial deficiency of 0.1 vefore the stripping event., In this case the galaxy already shows an initial deficiency of $0.4$ before the stripping event.
+ If such au initial deficiency is assuued. the sua of theobserved aud extrapolated N-vav gas mass nüeht account for the whole stripped gas.," If such an initial deficiency is assumed, the sum of theobserved and extrapolated X-ray gas mass might account for the whole stripped gas."
+ Alternatively. a fraction of the gas mass ndüeht be in the form of tenuous ionized eas with temperatures below LO° Is.," Alternatively, a fraction of the gas mass might be in the form of tenuous ionized gas with temperatures below $10^6$ K."
+the opposite.,the opposite.
+ relradius shows (hat (he planets radius is modestly smaller (han planets of similar mass and irradiation level., \\ref{radius} shows that the planet's radius is modestly smaller than planets of similar mass and irradiation level.
+ This suggests (hat the planet is more enriched in heavy. elements (han most other (ransiGue planets., This suggests that the planet is more enriched in heavy elements than most other transiting planets.
+ Given (he high stellar metallicitv. and the connection between stellar metallicity and planetary heavy elements (Guillot οἱ al.," Given the high stellar metallicity, and the connection between stellar metallicity and planetary heavy elements (Guillot et al."
+ 2006. Burrows et al.," 2006, Burrows et al."
+ 2007. Miller Fortney 2011). this is is not surprising.," 2007, Miller Fortney 2011), this is is not surprising."
+ Using the tables of Fortney et ((2007). we estimate a heavy element mass within the planet of at least 30-40 Earth masses.," Using the tables of Fortney et (2007), we estimate a heavy element mass within the planet of at least 30-40 Earth masses."
+ We report the discovery of Ixepler-15b. a hot Jupiter (P=4.94) that is enriched in heavy elements (20% in mass) in orbit around a metal-rich G-tvpe star with IN5—13.16.," We report the discovery of Kepler-15b, a hot Jupiter $P=4.94$ d) that is enriched in heavy elements $\approx 20\%$ in mass) in orbit around a metal-rich G-type star with $_{\rm p}$ =13.76."
+ ]xepler-15b is the first eiant. planet from theAepler mission that we confirmed. with the HET. demonstrating the capability of this facility as an integral part of the ground-based spectroscopic follow-up effort of theAepler mission.," Kepler-15b is the first giant planet from the mission that we confirmed with the HET, demonstrating the capability of this facility as an integral part of the ground-based spectroscopic follow-up effort of the mission."
+ In 2010 we spent a total of 65 hours on theAepler field ancl collected data for 11Kepler candidates., In 2010 we spent a total of 65 hours on the field and collected data for 11 candidates.
+ Besides the planet presented here. we have confirmed several other giant planets aroundAepler stars. e.g. IXepler-17b (Déssert et al.," Besides the planet presented here, we have confirmed several other giant planets around stars, e.g. Kepler-17b (Déssert et al."
+ 2011)., 2011).
+ The remaining WET planet confinmations will be included in a catalog olKepler eia. planets (Caldwell et al., The remaining HET planet confirmations will be included in a catalog of giant planets (Caldwell et al.
+ in prep.)., in prep.).
+ The HET will obtain a major upgrade to its secondary tracker assembly to allow a wider field of view., The HET will obtain a major upgrade to its secondary tracker assembly to allow a wider field of view.
+ This upgrade will start in fall 2011., This upgrade will start in fall 2011.
+ During the telescope downtime the IRS will also undergo a major upgrade. including more efficient optics and fibers. as well as nage slicers. to boost the overall throughput by several factors.," During the telescope downtime the HRS will also undergo a major upgrade, including more efficient optics and fibers, as well as image slicers, to boost the overall throughput by several factors."
+ Once the HET is back on skv in early 2012. these improvements should allow us to use the IIET/IIBS also for the confirmation and validation of low massAepler planet candidates in (he Neptune aud super-Earth range. similar to IXepler-10b (Batalha et 22011).," Once the HET is back on sky in early 2012, these improvements should allow us to use the HET/HRS also for the confirmation and validation of low mass planet candidates in the Neptune and Super-Earth range, similar to Kepler-10b (Batalha et 2011)."
+ Funding for this Discovery mission is provided bv NASA's Science Mission Directorate., Funding for this Discovery mission is provided by NASA's Science Mission Directorate.
+ The ILobby-Eberly. Telescope (LIET) is a joint project o£ the University of Texas at Austin. ihe Pennsylvania State University. Stanlorcd University. Lucdwig-Maximilians-Universitàtt Alünnchen. ancl Georg-August-Universitatt Gótttingen.," The Hobby-Eberly Telescope (HET) is a joint project of the University of Texas at Austin, the Pennsylvania State University, Stanford University, Ludwig-Maximilians-Universitätt Münnchen, and Georg-August-Universitätt Götttingen."
+ The WET is named in honor of its principal benefactors. William P. Hobby and Robert E. Eberly.," The HET is named in honor of its principal benefactors, William P. Hobby and Robert E. Eberly."
+ Dased in part on observations made with the Nordic Optical Telescope. operated on the islaud of La Palma jointly by Denmark. Finland. Iceland. Norway. and Sweden. in the Spanish Observatorio del Roque cle los Muchachos of the Instituto de Astrolisica de Canarias.," Based in part on observations made with the Nordic Optical Telescope, operated on the island of La Palma jointly by Denmark, Finland, Iceland, Norway, and Sweden, in the Spanish Observatorio del Roque de los Muchachos of the Instituto de Astrofisica de Canarias."
+in oplically (hick regions (Boss 1984). ancl so Because Lf=Fix. where F is the radiative Πας. one also finds We seek a solution that decouples the radiative transfer [rom the hvdrodynamies.,"in optically thick regions (Boss 1984), and so Because ${\vec H} = {\vec F} / 4 \pi$, where ${\vec F}$ is the radiative flux, one also finds We seek a solution that decouples the radiative transfer from the hydrodynamics."
+ As in IIubeny (1990). we assume that the test cloud or disk is at rest (v=0).," As in Hubeny (1990), we assume that the test cloud or disk is at rest ${\bf v} = 0$ )."
+ With this assumption. the energy equation becomes where /7j; is a heating term constant in time.," With this assumption, the energy equation becomes where $E_H$ is a heating term constant in time."
+ In steady state. (then. the heating term must equal —L. so that Assuming the spherically svmmetric test cloud has constant. uniform density p and constant. uniform opacity &. then The following racial temperature distribution is consistent with 7j constant in both space (inside the disk) and time," In steady state, then, the heating term must equal $- L$, so that Assuming the spherically symmetric test cloud has constant, uniform density $\rho$ and constant, uniform opacity $\kappa$, then The following radial temperature distribution is consistent with $E_H$ constant in both space (inside the disk) and time"
+We examine the size of the shifts in image position due to substructure. and without any lens modeling. in Figure 12..,"We examine the size of the shifts in image position due to substructure, and without any lens modeling, in Figure \ref{fig:pos}."
+ The differences between using the MOX substructure model and the ALL substructure model seem modest., The differences between using the M04 substructure model and the ALL substructure model seem modest.
+ In both cases. the median change in image position between a lens system with no substructure and one with substructure is 30 mas.," In both cases, the median change in image position between a lens system with no substructure and one with substructure is $\sim$ 30 mas."
+ However. when examining the images separately. we see a significant difference between substructure models.," However, when examining the images separately, we see a significant difference between substructure models."
+ For the ALL substructure model. the astrometric perturbations for {κ are significantly smaller than that of 7.4.," For the ALL substructure model, the astrometric perturbations for $\mathcal I_{-\bigstar}$ are significantly smaller than that of $\mathcal I_{+\bigstar}$."
+ In the MO4 case. these two cases are similarly large.," In the M04 case, these two cases are similarly large."
+ As mentioned previously. we add clumps to the positions where the input macromodel would find images in the absence of substructure.," As mentioned previously, we add clumps to the positions where the input macromodel would find images in the absence of substructure."
+ From Figure 12.. we see that whether the observed Image position is near a clump depends on the image.," From Figure \ref{fig:pos}, we see that whether the observed image position is near a clump depends on the image."
+" For example. in the ALL model. it seems that a significant number of faint images and of {Γκ may remain well aligned with their perturbers. while the observed /_, may have moved a significant distance relative to the scale radius of the clump."," For example, in the ALL model, it seems that a significant number of faint images and of $\mathcal I_{-\bigstar}$ may remain well aligned with their perturbers, while the observed $\mathcal I_{+\bigstar}$ may have moved a significant distance relative to the scale radius of the clump."
+ Radio observations of quasar sources strongly lensed by galaxy lenses provide information on the mass distribution along the line-of-sight and. às such. have been suggested as probes of galaxy substructure and a small-scale test of CDM.," Radio observations of quasar sources strongly lensed by galaxy lenses provide information on the mass distribution along the line-of-sight and, as such, have been suggested as probes of galaxy substructure and a small-scale test of CDM."
+ One simple measure of the subclumps ts the substructure mass fraction at the scale of the Einstein radius in a galaxy lens., One simple measure of the subclumps is the substructure mass fraction at the scale of the Einstein radius in a galaxy lens.
+ Simulations have predicted that this fraction is ως0.566 at z=0., Simulations have predicted that this fraction is $f_{\rm sub}\sim0.5\%$ at $z=0$.
+ Previous comparisons of observations to simulations have suggested that there may be more substructure than predicted in simulations with varying levels of discrepancy (2?2)..," Previous comparisons of observations to simulations have suggested that there may be more substructure than predicted in simulations with varying levels of discrepancy \citep{dalal_kochanek02,maccio_etal06,amara_etal06}."
+ ? pointed out that such magnification perturbations did not behave identically in each image and the behavior depended upon the parity (orientation). of the image., \citet{kochanek_dalal04} pointed out that such magnification perturbations did not behave identically in each image and the behavior depended upon the parity (orientation) of the image.
+ The observed parity dependence they found is such that the brightest negative parity image is demagnified with respect to the magnification predicted by smooth models and has a magnification perturbation smaller than that of the other images (0.4.«O and ὃκ<Oat omnes), The observed parity dependence they found is such that the brightest negative parity image is demagnified with respect to the magnification predicted by smooth models and has a magnification perturbation smaller than that of the other images $\delta_{-\bigstar} < 0$ and $\delta_{-\bigstar} < \delta_{\rm all ~others}$ ).
+ Here. we use the observed parity dependence in order to extract additional information about the substructure distribution.," Here, we use the observed parity dependence in order to extract additional information about the substructure distribution."
+ We create mock observations of strong lens systems with clumpy lens halos and compare the distribution of magnification perturbations between the mock and observed samples using an automated lens modeling code., We create mock observations of strong lens systems with clumpy lens halos and compare the distribution of magnification perturbations between the mock and observed samples using an automated lens modeling code.
+ Overall. our results are consistent with previous results which suggest that there is more substructure in the observations than can be accounted for in dark matter simulations.," Overall, our results are consistent with previous results which suggest that there is more substructure in the observations than can be accounted for in dark matter simulations."
+ Our conclusions are summarized as follows: Given the small sample of observed radio lenses. it is difficult to put robust constraints on the size or significance of the observed parity dependence.," Our conclusions are summarized as follows: Given the small sample of observed radio lenses, it is difficult to put robust constraints on the size or significance of the observed parity dependence."
+ In addition. while we approximate the observed parity dependence using simple substructure models. we cannot exclude more complicated models for the lens halos.," In addition, while we approximate the observed parity dependence using simple substructure models, we cannot exclude more complicated models for the lens halos."
+ We also consider only substructure within the lens halo itself anc do not consider the possible parity dependence effects of the significant numbers of clumps that may be found along the line-of-sight., We also consider only substructure within the lens halo itself and do not consider the possible parity dependence effects of the significant numbers of clumps that may be found along the line-of-sight.
+ Previous studies of the importance of line-of-sight clumps have found conflicting results (222)...," Previous studies of the importance of line-of-sight clumps have found conflicting results \citep{chen_etal03,metcalf05a,metcalf05b}."
+ The large amount of line-of-sight) structure suggests that this may be an effective way of resolving the parity dependence problem., The large amount of line-of-sight structure suggests that this may be an effective way of resolving the parity dependence problem.
+ However. line-of-sight clumps are less effective than substructure in inducing magnification perturbations and the overall effect of line-of-sight clumps ts unclear.," However, line-of-sight clumps are less effective than substructure in inducing magnification perturbations and the overall effect of line-of-sight clumps is unclear."
+ Parity dependence as observed in four-image lenses with radio observations is a strong measure of anomalous fluxes and a constraint on substructure in galaxies., Parity dependence as observed in four-image lenses with radio observations is a strong measure of anomalous fluxes and a constraint on substructure in galaxies.
+ In this paper. we describe the contradictions between the observed parity dependence and previous theoretical estimates and resolve the discrepancies using modeling of mock observations of lenses.," In this paper, we describe the contradictions between the observed parity dependence and previous theoretical estimates and resolve the discrepancies using modeling of mock observations of lenses."
+ We attempt to constrain the kinds of substructure models that will reproduce the observed parity dependence., We attempt to constrain the kinds of substructure models that will reproduce the observed parity dependence.
+ The substructure models necessary to approximate the observed results is consistent with suggestions that there is more substructure in galaxies than in estimates from simulations., The substructure models necessary to approximate the observed results is consistent with suggestions that there is more substructure in galaxies than in estimates from simulations.
+ In addition. we show that it is unlikely that anomalous flux ratios are due solely to observed luminous satellite galaxies.," In addition, we show that it is unlikely that anomalous flux ratios are due solely to observed luminous satellite galaxies."
+ Future studies of CDM substructure using strong gravitational lenses may require larger samples of observed lenses as well as more physically realistic simulations of dark matter and baryons., Future studies of CDM substructure using strong gravitational lenses may require larger samples of observed lenses as well as more physically realistic simulations of dark matter and baryons.
+tacitly make is that both the primary and secondary stars under consideration are on the Main Sequence.,tacitly make is that both the primary and secondary stars under consideration are on the Main Sequence.
+" This may seem obvious, but we note that many known binaries with a hot O, B type dwarf star have a cool, G, K or M-type giant secondary (see e.g. ?))."," This may seem obvious, but we note that many known binaries with a hot O, B type dwarf star have a cool, G, K or M-type giant secondary (see e.g. \citealt{ginestet_2002}) )."
+" In our sample, HR 2577, an unresolved spectroscopic binary with a K2II companion listed in Table 1 is a prime example."," In our sample, HR 2577, an unresolved spectroscopic binary with a K2II companion listed in Table \ref{logtarget} is a prime example."
+" In most such binary stars, the B-type component dominates at optical wavelengths, while the secondary dominates at longer wavelengths, simply because cooler stars are redder."," In most such binary stars, the B-type component dominates at optical wavelengths, while the secondary dominates at longer wavelengths, simply because cooler stars are redder."
+" In contrast, the current study is conducted in the infrared already and the brightest binary component is the previously known B-star whose cataloguedK band magnitide is consistent with the optical brightness, while the secondary is fainter."," In contrast, the current study is conducted in the infrared already and the brightest binary component is the previously known B-star whose catalogued band magnitide is consistent with the optical brightness, while the secondary is fainter."
+" For reference, we have also plotted the mass of a selected number of giant stars against their intrinsicK band magnitudes in Fig. 4.."," For reference, we have also plotted the mass of a selected number of giant stars against their intrinsic band magnitudes in Fig. \ref{sptpfig}."
+" The cool G, K and M-giants are brighter than most B stars up to a spectral type of B3V. This means that for all stars with spectral types of B3 or later in our sample, we can be fairly certain that the companion is not a giant as the primary star, the B star, is the brightest."," The cool G, K and M-giants are brighter than most B stars up to a spectral type of B3V. This means that for all stars with spectral types of B3 or later in our sample, we can be fairly certain that the companion is not a giant as the primary star, the B star, is the brightest."
+" Indeed, only for BOV-B3V stars there would be a chance that the (fainter) companion is a giant rather than a main sequence dwarf star."," Indeed, only for B0V-B3V stars there would be a chance that the (fainter) companion is a giant rather than a main sequence dwarf star."
+" In such a case, the magnitude difference will be at most 2 magnitudes."," In such a case, the magnitude difference will be at most 2 magnitudes."
+ Inspection of Table 2 however reveals that almost all binaries with a BOV-B3V primary have a magnitude difference larger than 2., Inspection of Table \ref{restarget} however reveals that almost all binaries with a B0V-B3V primary have a magnitude difference larger than 2.
+ We conclude that it is very likely that the secondary stars are main sequence stars and that we are therefore justified in applying our K-mass relation to determine the mass ratio of the systems., We conclude that it is very likely that the secondary stars are main sequence stars and that we are therefore justified in applying our -mass relation to determine the mass ratio of the systems.
+" One correction needs to be made for the Be stars, the free-free emission from the ionized disk increases with wavelength, and can give slight excess emission at the K-band (e.g. ?))."," One correction needs to be made for the Be stars, the free-free emission from the ionized disk increases with wavelength, and can give a slight excess emission at the -band (e.g. \citealt{dougherty_1991}) )."
+" To arrive at the intrinsica magnitude of the star itself, we need to subtract the disk contribution from theK band flux."," To arrive at the intrinsic magnitude of the star itself, we need to subtract the disk contribution from the band flux."
+" This is done in the standard manner by dereddening theV band magnitude (using intrinsic colours for the respective spectral types from Schmidt-Kaler in ?,Chap. 4)) and computing the predictedK band magnitude using theK colours as outlined above.", This is done in the standard manner by dereddening the band magnitude (using intrinsic colours for the respective spectral types from Schmidt-Kaler in \citet[Chap. 4]{Landolt}) ) and computing the predicted band magnitude using the colours as outlined above.
+ The difference between the predicted and observed K-band magnitude is the excess., The difference between the predicted and observed -band magnitude is the excess.
+" In two cases, HR 6712 and HR 5193, the excesses were negative by up to half a magnitude and thus unphysical."," In two cases, HR 6712 and HR 5193, the excesses were negative by up to half a magnitude and thus unphysical."
+ These stars are variable and using the optical colours from SIMBAD rather than from the BSC solved this situation., These stars are variable and using the optical colours from SIMBAD rather than from the BSC solved this situation.
+" The “excesses” computed for the normal B stars averaged 0.01 magnitude with a scatter of 0.1 magnitude, which gives an indication of the uncertainties involved in this exercise."," The “excesses” computed for the normal B stars averaged 0.01 magnitude with a scatter of 0.1 magnitude, which gives an indication of the uncertainties involved in this exercise."
+" As all magnitude differences between the companions are much larger than 0.3 magnitudes, this error does not affect the resulting mass ratio."," As all magnitude differences between the companions are much larger than 0.3 magnitudes, this error does not affect the resulting mass ratio."
+" In fact, only two Be stars have computed excesses which are much larger than 0.3 magnitude, these are HR 2142 (0.6 mag) and HR 2356 (0.8 mag)."," In fact, only two Be stars have computed excesses which are much larger than 0.3 magnitude, these are HR 2142 (0.6 mag) and HR 2356 (0.8 mag)."
+" After correcting for the excess emission, we can then compute the photospheric magnitude difference and the mass ratio q."," After correcting for the excess emission, we can then compute the photospheric magnitude difference and the mass ratio $q$."
+ In Fig., In Fig.
+ 5 the results for the mass ratios are presented in a similar fashion as for the separations in the previous figure., \ref{mratiofig} the results for the mass ratios are presented in a similar fashion as for the separations in the previous figure.
+ Note that the two objects for whose close companions we were unable to retrieve accurate photometry are not plotted., Note that the two objects for whose close companions we were unable to retrieve accurate photometry are not plotted.
+" The left hand plot shows the distribution of mass ratios, which range from 0.07 to 0.9."," The left hand plot shows the distribution of mass ratios, which range from 0.07 to 0.9."
+" Based on the relatively low numbers, it is hard to describe this distribution, but it would seem that the distributions are similar and rather flat for the B and Be stars."," Based on the relatively low numbers, it is hard to describe this distribution, but it would seem that the distributions are similar and rather flat for the B and Be stars."
+ It does not rise steeply to low masses as would be the case if the secondary masses were randomly sampled from the Initial Mass Function (see e.g. ?))., It does not rise steeply to low masses as would be the case if the secondary masses were randomly sampled from the Initial Mass Function (see e.g. \citealt{Wheelwright_2010}) ).
+ The cumulative plot in the right hand panel confirms that there is not a significant difference between the two distributions., The cumulative plot in the right hand panel confirms that there is not a significant difference between the two distributions.
+ The KS test returns a probability that the samples are drawn from different populations can be rejected at50%.., The KS test returns a probability that the samples are drawn from different populations can be rejected at.
+" As with the separations, if the mass ratios would have been very different, the current sample is large enough to highlight this."," As with the separations, if the mass ratios would have been very different, the current sample is large enough to highlight this."
+ It therefore seems likely that the mass ratios of both the B and Be type binaries too are sampled from, It therefore seems likely that the mass ratios of both the B and Be type binaries too are sampled from
+carbonate dust in astronomical objects. ib is important (o discuss (he relation between such powder spectra and the small particle spectra based on our optical constants.,"carbonate dust in astronomical objects, it is important to discuss the relation between such powder spectra and the small particle spectra based on our optical constants."
+ By ‘powder spectra or “PE speclra. we understand the transmission spectra of sub-micron sized ground powder embecded in polvethylene (PE). transformed inlo absorption cross sections.," By `powder spectra' or `PE spectra', we understand the transmission spectra of sub-micron sized ground powder embedded in polyethylene (PE), transformed into absorption cross sections."
+ Powder spectra generally depend on circumstances such as particle ---- and shape. orientation. agelomeration degree. and the relractive index of the surrounding medium ).," Powder spectra generally depend on various circumstances such as particle size and shape, orientation, agglomeration degree, and the refractive index of the surrounding medium \\cite{BH83}) )."
+ rel PECDEshowsboththepowderspeetrumpublishedby Nen —measuredonasampleof unsedimentedealeitegrains(dash- dolled—line)—-—andofasampleofealcitegrainswhichhadpr to s," \\ref{f:PE_CDE} shows both the powder spectrum published by Kemper et (2002a) -- measured on a sample of unsedimented calcite grains (dash-dotted-line) – and of a sample of calcite grains which had previously been subjected to sedimentation in a liquid (acetone), in order to remove large grains (solid line)."
+horter wavelengths with respect to the data published by temper et ((2002a). although both have been measured at room temperature.," The peak of the latter spectrum is shifted by about $\mu$ m to shorter wavelengths with respect to the data published by Kemper et (2002a), although both have been measured at room temperature."
+"particle p.CC,,,,, aud coiipute it considering the gravitational interactiou between tle VD aud ouly the elements listed in ££y,,. while the Ἑρω component is computed separately for cach p particle by the direct iuteraction with the elemoeuts listed in ZZ,","particle $p \in C_{group}$ and compute it considering the gravitational interaction between the VB and only the elements listed in $IL_{far}$, while the ${\bf F}_{near}$ component is computed separately for each $p$ particle by the direct interaction with the elements listed in $IL_{near}$."
+" Moreover. ρω coutaius a restricted προ of elements in comparison with the Fy,,. list. so we expect a uct gain in performauce even if 7;€Ty."," Moreover, ${\bf F}_{near}$ contains a restricted number of elements in comparison with the ${\bf F}_{far}$ list, so we expect a net gain in performance even if $ T_{l} \le T_{f}$."
+" The gain that is possible depends ou several paranneters CN,ig. the size of the Coron, and the radiis). whose rauges of variation are coustrained by the ιαππα allowed value of the overall error of the method. as we will describe in the following sections."," The gain that is possible depends on several parameters $N_{crit}$, the size of the $C_{group}$ and the ), whose ranges of variation are constrained by the maximum allowed value of the overall error of the method, as we will describe in the following sections."
+ Before showing the performance of our WD99 procedure in an N-body simulation of the large scale structure of the universe. it is inportaut that we perform an error analysis of the procedure itself.," Before showing the performance of our WD99 procedure in an N-body simulation of the large scale structure of the universe, it is important that we perform an error analysis of the procedure itself."
+" Cousidering that the cumulative error. when sinulatious for the LSS studies are run using the original DIT αἰσοια, is lower than 1[6].. fixing the opening criterion 0= 1. we will give some constraint concerning the size of Cu,. the IN.í; value. aud theradius needed to have ueelieible cumulative error."," Considering that the cumulative error, when simulations for the LSS studies are run using the original BH algorithm, is lower than $1\%$, fixing the opening criterion $\theta=1$ , we will give some constraint concerning the size of $C_{group}$, the $N_{crit}$ value, and the needed to have negligible cumulative error."
+ The following sub-sectious discuss the two main sources of error., The following sub-sections discuss the two main sources of error.
+" The first error source is that WD99 uses the Sphere criterion aud the VD to create au interaction ist £L,=ILy aud WD99 applies the LLy-4 to all bodies pCορουμ.", The first error source is that WD99 uses the Sphere criterion and the VB to create an interaction list $IL_g \equiv IL_{VB}$ and WD99 applies the $IL_{VB}$ to all bodies $p \in C_{group}$.
+" This approximation could introduce an errorgp iu the force value ou the p particle if the ZL,. created using the original DII algoritlin. aud the ZZy:5 have a differcuce in the clements ercater then J 4."," This approximation could introduce an error in the force value on the $p$ particle if the $IL_p$, created using the original BH algorithm, and the $IL_{VB}$ have a difference in the elements greater then $1\%$ ."
+" As we found with our ests. in order to decrease this difference it is necessary to limit the size of Cyan, flat is equivalent o fixiug adevel of the tree structure: cells above the critical level cannot form a eroupiug cell."," As we found with our tests, in order to decrease this difference it is necessary to limit the size of $C_{group}$ that is equivalent to fixing a of the tree structure: cells above the critical level cannot form a grouping cell."
+" The user has to fix the critical level considering the density of local bodies iu the box where voces are arranged: the critical level iust be chosen im such a wav as to make the difference vetwoen Z£, aud LLy-p negligible (uo more than 1% of the clements).", The user has to fix the critical level considering the density of local bodies in the box where bodies are arranged: the critical level must be chosen in such a way as to make the difference between $IL_p$ and $IL_{VB}$ negligible (no more than $1\%$ of the elements).
+ It seems reasonable for a LSS simulation iu a 50 Mpe box with more thin 2 million particles to fix the critical level as he sixth level of the tree., It seems reasonable for a LSS simulation in a 50 Mpc box with more than 2 million particles to fix the critical level as the sixth level of the tree.
+ The uext section shows in detail the obtaimed results., The next section shows in detail the obtained results.
+" In any case. the cumulative error is very siiall considering the merease in accuracy compared with the original BITalgoritlin. due to the incision of all pCC,,,,,iu the interaction list."," In any case, the cumulative error is very small considering the increase in accuracy compared with the original BHalgorithm, due to the inclusion of all $p \in C_{group}$in the interaction list."
+of the dSph against the tical forces of the massive host galaxy or of the galaxy cluster or group in which the dSph resides.,of the dSph against the tidal forces of the massive host galaxy or of the galaxy cluster or group in which the dSph resides.
+ This would explain why only a handful of dSphs in the Local Group show clear signs of an ongoing interaction despite being close satellites of either the Milky Way or M31 (Johnston.Spergel.Lernquist1995:Mateoetal.1996: 2007).," This would explain why only a handful of dSphs in the Local Group show clear signs of an ongoing interaction despite being close satellites of either the Milky Way or M31 \citep{jsh95,m96,mi06,s07,l07}."
+. Alany dSphs and dl£s still contain an interstellar medium and some even host low-level star formation (Blitz2005:Liskeretal.2006:Youngct 2007).. showing that supernova explosions are not very efficient at expelling gas (Marcolini.Brighenti.D'Ircole.|2003:Valcke.," Many dSphs and dEs still contain an interstellar medium and some even host low-level star formation \citep{br00,ggh03,cogw03,de03,b05,l06,y07}, showing that supernova explosions are not very efficient at expelling gas \citep{mbe03,vdd08}."
+DeRijcke.Dejonghe| 2005).. Aiüght elliptical. galaxies. or Es. and cles follow the same photometric and kinematic scaling relations (Graham2005:SmithCastellietal. 2008).," Bright elliptical galaxies, or Es, and dEs follow the same photometric and kinematic scaling relations \citep{gg03,mg05,de05,smc08}."
+. In the luminosity interval 24magMc14 mag the parameters of the Sérrsic profile follow simple power-laws as a function of luminosity and carly ancl late type galaxies trace parallel Tullv-Fisher relations (DeRijekeetal...2007).," In the luminosity interval $-24\,\,{\rm mag} \lesssim {\rm M}_V \lesssim
+-14$ mag the parameters of the Sérrsic profile follow simple power-laws as a function of luminosity and early and late type galaxies trace parallel Tully-Fisher relations \citep{de07}."
+. From this wealth of data a picture of (dwarf) galaxy formation emerges that suggests an underlying unity in the physics driving the formation and evolution of stellar systems. with the environment plaving a role that is in many situations subordinate to that of internal processes.," From this wealth of data a picture of (dwarf) galaxy formation emerges that suggests an underlying unity in the physics driving the formation and evolution of stellar systems, with the environment playing a role that is in many situations subordinate to that of internal processes."
+ More specifically. numerical simulations and semi-analvtic models. of galaxy formation within a ACDAL cosmology can account lor the observed. scaling relations when taking into account supernova feedback in galactic gravitational potential wells steepening with galaxy mass (Carraroetal.2001:Nagashima&Yoshii2004:lüjceke.Dejonghe 2008).," More specifically, numerical simulations and semi-analytic models of galaxy formation within a $\Lambda$ CDM cosmology can account for the observed scaling relations when taking into account supernova feedback in galactic gravitational potential wells steepening with galaxy mass \citep{c01,ny04,ri05,m06,vdd08}."
+. The goal of this paper is to investigate whether the photometric scaling relations traced by cles ancl Es persist own to the dSphs (M14 mag)., The goal of this paper is to investigate whether the photometric scaling relations traced by dEs and Es persist down to the dSphs ${\rm M}_V \gtrsim -14$ mag).
+ We also check [or possible environmental influences. other than the obvious ensitv-morphology relation (i.e. the fact that dls/dSphs ave found. predominantly. in. high-density cnvironments).," We also check for possible environmental influences, other than the obvious density-morphology relation (i.e. the fact that dEs/dSphs are found predominantly in high-density environments)."
+ We present new cata based on our LIST/ACS imaging of dSphs/dEs in the Perseus Cluster and combine these with 6ata of earlv-tygalaxies in the Antlia. Fornax. and Virgo Clusters ancl the peLocal Group (see section 2)).," We present new data based on our HST/ACS imaging of dSphs/dEs in the Perseus Cluster and combine these with data of early-typegalaxies in the Antlia, Fornax, and Virgo Clusters and the Local Group (see section \ref{data}) )."
+ The resulting 1;hotometric scaling laws are presented in section 3..., The resulting photometric scaling laws are presented in section \ref{dare}.
+ We 6iscuss the results in section 4.., We discuss the results in section \ref{disc}. .
+ Phe Perseus Cluster (Abell 426) is one of the richest nearby ealaxy clusters. with a redshift ο = 5366 km (Struble&Rood 1999).. and at a distance D = 72 Alpe (as given by NED).," The Perseus Cluster (Abell 426) is one of the richest nearby galaxy clusters, with a redshift $v$ = 5366 km $^{-1}$ \citep{Stublerood99}, and at a distance $D$ = 72 Mpc (as given by NED)."
+" Due to its low Galactic latitude (6 = 18"") it has not been studied in as much detail as other nearby clusters such as Fornax. Virgo and Coma."," Due to its low Galactic latitude $b$ $\approx$ $-13^{\circ}$ ) it has not been studied in as much detail as other nearby clusters such as Fornax, Virgo and Coma."
+ We have obtained high resolution Xdvanced Camera lor Surveys (ACS) WEC imaging in the F555W and FSI4W bands of five fields in the Perseus Cluster core. in the immediate vicinity of NGCI275 and NGC1272. the cluster’s brightest members. obtained in 2005 (program. GO 10201).," We have obtained high resolution Advanced Camera for Surveys (ACS) WFC imaging in the F555W and F814W bands of five fields in the Perseus Cluster core, in the immediate vicinity of NGC1275 and NGC1272, the cluster's brightest members, obtained in 2005 (program GO 10201)."
+" The scale of the images is 0.05” pixel+ with a [eld of view of 202"" providing a total survey area of ~ 57 arcemin."," The scale of the images is $''$ $^{-1}$, with a field of view of $202''
+\times 202''$, providing a total survey area of $\sim$ 57 $^{2}$."
+2 The positions on the sky of these five pointings are presented in Fig. l.., The positions on the sky of these five pointings are presented in Fig. \ref{pointings}.
+ Exposure times were 2368 and 2260 seconds for the F555W and ESI4NV bands. respectively.," Exposure times were 2368 and 2260 seconds for the F555W and F814W bands, respectively."
+ Ehe fields were chosen to cover the most likely cluster dSphs and clés identified. fron eround-based imagery hy Consclice.Gallagher.Wyse(2003)., The fields were chosen to cover the most likely cluster dSphs and dEs identified from ground-based imagery by \cite{cgw03}.
+. For some of these. there is spectroscopic confirmation of their cluster membership (Penny&Conseliee2008).," For some of these, there is spectroscopic confirmation of their cluster membership \citep{pc08}."
+. For the others. we use morphological criteria to decide cluster membership.," For the others, we use morphological criteria to decide cluster membership."
+ The CAS system for quantifying compactness. asvniumctry. and clumpiness/smoothness (Conselice2003b) proves very useful to reject e.g. background spiral galaxies based on a smoothness criterion and background bright cllipticals based on a Compactness criterion (sec Pennyetal.(9005111.," The CAS system for quantifying compactness, asymmetry, and clumpiness/smoothness \citep{c03} proves very useful to reject e.g. background spiral galaxies based on a smoothness criterion and background bright ellipticals based on a compactness criterion (see \cite{pe08}) )."
+ The Perseus dataset straddles the dlz-dSph transition at My148 mage and is therefore essential to the discussion that follows., The Perseus dataset straddles the dE-dSph transition at $_V \sim -14$ mag and is therefore essential to the discussion that follows.
+ A detailed account of our photometric analysis of the USTΑΟ images of Perseus dSphs and of the properties of the individual galaxies will be reported elsewhere. (Penny. Conselice. De Rijeke. Held. in prep.)," A detailed account of our photometric analysis of the HST/ACS images of Perseus dSphs and of the properties of the individual galaxies will be reported elsewhere (Penny, Conselice, De Rijcke, Held, in prep.)."
+ The photometric data. including resolved. photometryfor surface brightness profiles. of the Local Group dSphs that are identified as Milkv Way satellites are collected from. Grebel.Gallagher.Harbeck:(2003) and Irwinitviow (1995).. adopting the distances listed in Cirebel.Ga (2003)..," The photometric data, including resolved photometryfor surface brightness profiles, of the Local Group dSphs that are identified as Milky Way satellites are collected from \cite{ggh03} and \cite{ih95}, adopting the distances listed in \cite{ggh03}. ."
+ Data of the M31 dSph satellites iJa aken from Peletier(1993)... Caldwell: (1999).. Grebel.Gaagher.Llarbeck|. (2003)... AleConnachic&Irwin. (2006)... AleConnachic.Arimoto.Lewin (2007).. and Zucker (2007).," Data of the M31 dSph satellites is taken from \cite{p93}, \cite{c99}, \cite{ggh03}, \cite{mi06}, \cite{mai07}, and \cite{z07}."
+. Data of three Local Group dSphs that are not linked o à giant host galaxy. the “Pucana dSph. DDO2IO.. ane IxIx25. come from Saviane.Held.Piotto(1996)... Gallagher.Larbeck—. (2003).. and. AleConnachiej i).," Data of three Local Group dSphs that are not linked to a giant host galaxy, the Tucana dSph, DDO210, and KKR25, come from \cite{sv96}, \cite{ggh03}, and \cite{mi06}."
+ DeRijekeetal...(2005). (D05) and. Mieskectal.(207) provide photometric data on the earlv-tvpe dwarf galaxy population of the Fornax cluster., \cite{de05} (D05) and \cite{mi07} provide photometric data on the early-type dwarf galaxy population of the Fornax cluster.
+ Half of the D05 sample consists of clés from the Νέας5044. and NCGCB5980 groups., Half of the D05 sample consists of dEs from the NGC5044 and NGC5989 groups.
+ The data of the dSphs and clés in the Antlia eluster are taken from SmithCastellietal.(2008)., The data of the dSphs and dEs in the Antlia cluster are taken from \cite{smc08}.
+. Data data for the giant elliptical and for Coma cles is taken [rom Crahans&Guzmán.(2003). (C03)., Data data for the giant elliptical and for Coma dEs is taken from \cite{gg03} (GG03).
+ This sample of early-type galaxies comprises dwarf spheroidals. with 14magMoXS mag. dwarl ellipticals. with 19magZzAL14 mag. and bright ellipticals. with MiX19 mag.," This sample of early-type galaxies comprises dwarf spheroidals, with $-14\,\,{\rm mag} \lesssim {\rm M}_V \lesssim -8$ mag, dwarf ellipticals, with $-19\,\,{\rm mag} \lesssim {\rm M}_V \lesssim
+-14$ mag, and bright ellipticals, with ${\rm M}_V \lesssim
+-19$ mag."
+ We plot the positions of the sample galaxies in diagrams of V-band absolute magnitude vs. half-light radius It. (in kpc). vs. the Sérrsic exponent n of the best fittingSérrsic profile. vs.(2 the central V-band surface brightness of the best fitting Sérrsic profile. and vs. V. Ecolour.," We plot the positions of the sample galaxies in diagrams of V-band absolute magnitude vs. half-light radius $_{\rm e}$ (in kpc), vs. the Sérrsic exponent $n$ of the best fittingSérrsic profile, vs. the central V-band surface brightness of the best fitting Sérrsic profile, and vs. $-$ I colour."
+ These diagrams are shown in Fig. 2.., These diagrams are shown in Fig. \ref{LVVI}. .
+ The datasets are plotted using dilferent symbols that are explained within each panel of the figure., The datasets are plotted using different symbols that are explained within each panel of the figure.
+ Phe presence of a dataset in a given panel depends solelyon the availability of the required. data., The presence of a dataset in a given panel depends solelyon the availability of the required data.
+empirical colour-temperature relations are preferred over the heoretical ones (see Paper LV’).,empirical colour-temperature relations are preferred over the theoretical ones (see Paper IV).
+ Following the procedure described in refcalibracion.. we have calibrated the elfective temperatures xedicted by ALO onto the svstem of SINC (see Fig. 4)).," Following the procedure described in \\ref{calibracion}, we have calibrated the effective temperatures predicted by ALO onto the system of SKC (see Fig. \ref{compALOSKC}) )."
+ For emperatures derived from the 2V relations. a statistically significant olfset of 26 Ix was found. which has been applied o correct and. bootstrap the predicted. cata against. the reference system.," For temperatures derived from the $B-V$ relations, a statistically significant offset of 26 K was found, which has been applied to correct and bootstrap the predicted data against the reference system."
+ On the contrary. no significant deviation mas been found for temperatures derived from the WoA relations.," On the contrary, no significant deviation has been found for temperatures derived from the $V-K$ relations."
+" In addition. the fact that. the bens stancdaux deviations from the fits are very similar (tp,y = 75 In anc Oy w= TSW.) allows us to deduce that the data equality of both predictions is the same."," In addition, the fact that the r.m.s standard deviations from the fits are very similar $_{\rm
+B-V}$ = 75 K and $_{\rm V-K}$ = 78 K.) allows us to deduce that the data quality of both predictions is the same."
+ Thus. final effective. temperatures for cluster stars were calculated by averaging the temperatures derivec [rom Dl and VÁN relations (when only £2V was available. final temperatures were exclusively derivec from this colour).," Thus, final effective temperatures for cluster stars were calculated by averaging the temperatures derived from $B-V$ and $V-K$ relations (when only $B-V$ was available, final temperatures were exclusively derived from this colour)."
+ Input. metallicities are those previously established in Section 6.1.. while gravities are the same as in G93 and W94.," Input metallicities are those previously established in Section \ref{metalcum}, while gravities are the same as in G93 and W94."
+ Concerning the input colours. Vdy values where taken from G93. whilst 2V colours are from Alermilliod et al.," Concerning the input colours, $V-K$ values where taken from G93, whilst $B-V$ colours are from Mermilliod et al."
+ (1997). for Coma ane Hyacdes. N94 for a [ow horizontalbranch stars. and. G93 for the rest of the sample.," \shortcite{MER} for Coma and Hyades, W94 for a few horizontal–branch stars, and G93 for the rest of the sample."
+ The reddening corrections were applied using the color excesses given by C93., The reddening corrections were applied using the color excesses given by G93.
+ ]t is important to note that. due to the validity range of the colour-temperature calibrations. the effective temperatures of the hot stars Ilva vB 104. M5 11-53. M92 1-10. M92 11-23 and. M92. NLI-24. could. not be predicted in this way.," It is important to note that, due to the validity range of the colour-temperature calibrations, the effective temperatures of the hot stars Hya vB 104, M5 II-53, M92 I-10, M92 II-23 and M92 XII-24 could not be predicted in this way."
+ In these cases. temperatures from. WO4 were kept.," In these cases, temperatures from W94 were kept."
+ Surface eravities for most of the cluster stars were originally derived by G93 by fitting the location of the stars in the IIR diagrams to cillerent evolutionary. tracks (see the original paper [or details on the tracks that were used)., Surface gravities for most of the cluster stars were originally derived by G93 by fitting the location of the stars in the HR diagrams to different evolutionary tracks (see the original paper for details on the tracks that were used).
+ In the recent vears. and thanks to the new Llipparcos data. distances and absolute ages of most globular clusters have been substantially modified (see e.g. Carretta et al.," In the recent years, and thanks to the new Hipparcos data, distances and absolute ages of most globular clusters have been substantially modified (see e.g. Carretta et al."
+ 2000 and references therein)., 2000 and references therein).
+ In. this work we have studied what changes should be introduced in the derived surface &ravities to account for the new. usually larger. distances and. therefore. vounger ages for the elobular cluster stars.," In this work we have studied what changes should be introduced in the derived surface gravities to account for the new, usually larger, distances and, therefore, younger ages for the globular cluster stars."
+ Changes in the effective. temperatures. do not allect the derived. gravities since the original values were computed by matching only absolute magnitudes to avoid uncertainties in the colour-Z;4r determinations., Changes in the effective temperatures do not affect the derived gravities since the original values were computed by matching only absolute magnitudes to avoid uncertainties in the $T_{\rm eff}$ determinations.
+ Using the relation z;Mlogg)=Af(log(Al/Al.))|0.442; and applving the analytic formulae by Hurley. Pols ‘Tout (2000) to convert an age dillerence to a mass change for the," Using the relation $\Delta(\log g)=\Delta(\log (M/M_\odot)) + 0.4\Delta M_{\rm
+bol}$ and applying the analytic formulae by Hurley, Pols Tout (2000) to convert an age difference to a mass change for the"
+In an optical interferometer light is collected al two or more apertures ancl brought to a central location where the beams are combined and a Inge pattern produced.,In an optical interferometer light is collected at two or more apertures and brought to a central location where the beams are combined and a fringe pattern produced.
+ For a broadband source of wavelength A the Ginee pattern is limited in extent and appears only when the optical paths through the arms of the interferometer are equalized to within a coherence length (A=A7/ AA).," For a broadband source of wavelength $\lambda$ the fringe pattern is limited in extent and appears only when the optical paths through the arms of the interferometer are equalized to within a coherence length $\Lambda =
+\lambda^2/\Delta\lambda$ )."
+" For a two-apertiure interferometer. neglecting dispersion. the intensity measured at one of the combined beams is given by where V is the lringe contrast or ""visibili. which can be related to the morphology. of the source. and AA is the optical bandwidth of the interferometer assuming a flat optica. bandpass (for PTT AXA= 0.4jun)."," For a two-aperture interferometer, neglecting dispersion, the intensity measured at one of the combined beams is given by where $V$ is the fringe contrast or “visibility”, which can be related to the morphology of the source, and $\Delta\lambda$ is the optical bandwidth of the interferometer assuming a flat optical bandpass (for PTI $\Delta\lambda = 0.4 \mu$ m)."
+ The differential (between arms of the interferometer; optical path 6 is found from geometric considerations to be Bis the baseline (he vector connecting the νο apertures., The differential (between arms of the interferometer) optical path $\delta$ is found from geometric considerations to be $\vec{B}$ is the baseline – the vector connecting the two apertures.
+ sis the unit vector in the source direction. and d is any differential optical path introduced by (he instrument.," $\vec{s}$ is the unit vector in the source direction, and $d$ is any differential optical path introduced by the instrument."
+ For a 100-m baseline interferometer an astrometric precision of 10 µας corresponds to knowing d (o 5 nm. a difficult but not impossible proposition.," For a 100-m baseline interferometer an astrometric precision of 10 $\mu$ as corresponds to knowing $d$ to 5 nm, a difficult but not impossible proposition."
+ The dominant source of measurement error is aümospheric turbulence above (he interferometer. which adds varving amounts of optical path and hence makes the fringes appear to move about rapidly (ποιος1981).," The dominant source of measurement error is atmospheric turbulence above the interferometer, which adds varying amounts of optical path and hence makes the fringes appear to move about rapidly \citep{rod81}."
+. This atmospheric turbulence. which changes over distances of tens of centimeters ancl millisecond Gme-scales. forces one to use very short exposures to maintain [fringe contrast. and hence limits the sensitivity of the instrument.," This atmospheric turbulence, which changes over distances of tens of centimeters and millisecond time-scales, forces one to use very short exposures to maintain fringe contrast, and hence limits the sensitivity of the instrument."
+ It also severely limits the astrometric accuracy of a simple interferometer. al least. over large skv angles.," It also severely limits the astrometric accuracy of a simple interferometer, at least over large sky angles."
+" llowever. (his atmospheric turbulence is correlated. over small angles. and hence it is still possible to obtain hieh precision ""narrow-angle astrometry."," However, this atmospheric turbulence is correlated over small angles, and hence it is still possible to obtain high precision “narrow-angle” astrometry."
+ Dual-star interferometric narrow-angle astrometry (Shao&1994) promises astrometric performance at the 10-100 µας level for pairs of stars separated by 10-60 areseconds. and has been demonstrated at PTI.," Dual-star interferometric narrow-angle astrometry \citep{shao92,col94}
+ promises astrometric performance at the 10-100 $\mu$ as level for pairs of stars separated by 10-60 arcseconds, and has been demonstrated at PTI."
+ Achieving such performance requires, Achieving such performance requires
+shell of compensation.,shell of compensation.
+ In order to verify the effect of such a shell. let us first assume that the average matter density inside the layer is constant. pon. and its thickness is A/?.," In order to verify the effect of such a shell, let us first assume that the average matter density inside the layer is constant, $\rho_\mathrm{sh}$, and its thickness is $\Delta R$."
+ The Local Group. of negligible size. is still located at the edge of the Local Void. as in Figures.," The Local Group, of negligible size, is still located at the edge of the Local Void, as in Figures."
+ Putting now the mass element inside the intersection of the shell with the ΖοΑ. we can see that the differential spurious acceleration will read velocity with q and R=(#.0) defined as before. but nowAR.," Putting now the mass element inside the intersection of the shell with the ZoA, we can see that the differential spurious acceleration will read , with $\bmath{q}$ and $\bmath{R}=(R,0)$ defined as before, but now."
+" The y-component of the acceleration vanishes from symmetry. and along the . direction we havefor with Z.j, different from Z only by the limits of integration in £ and equal to ΠΠΙΗΕ ταd£Rose."," The $y$ -component of the acceleration vanishes from symmetry, and along the $x$ direction we have, with $\mathcal{I}_\mathrm{sh}$ different from $\mathcal{I}$ only by the limits of integration in $\xi$ and equal to _0 _1."
+ This result means that the spurious acceleration of the LG due to the part of the shell of compensation hidden behind the ZoA will also vanish if the density distribution p. is not constant. but depends on the distance from the centre of the LV (as we can divide the shell into infinitesimally thin layers of constant density each).," This result means that the spurious acceleration of the LG due to the part of the shell of compensation hidden behind the ZoA will also vanish if the density distribution $\rho_\mathrm{sh}$ is not constant, but depends on the distance from the centre of the LV (as we can divide the shell into infinitesimally thin layers of constant density each)."
+ The bottom-line is that a compensating shell with radial density distribution will affect the spurious acceleration of the Local Group provided that the latter is located at the inner edge of the shell. as observational constraints On the other hand. if the LG was placed in the interior of the compensating layer. or at its outer edge. the discussed spurious acceleration from the LV would bediminished.. by up to 2/3 in the limit of an infinitesimally thin shell.," The bottom-line is that a compensating shell with radial density distribution will affect the spurious acceleration of the Local Group provided that the latter is located at the inner edge of the shell, as observational constraints On the other hand, if the LG was placed in the interior of the compensating layer, or at its outer edge, the discussed spurious acceleration from the LV would be, by up to $2\slash 3$ in the limit of an infinitesimally thin shell."
+ Having calculated tye amplitude of the spurious acceleration in the model. we can appy our results to observational data.," Having calculated the amplitude of the spurious acceleration in the model, we can apply our results to observational data."
+ As an example. we take the analysis of Malleretal.(2003) of the clustering dipole of the 2-Micron All-Sky Survey (2MASS. Skrutskie 2006). since this is the censest and the deepest all-sky survey to date.," As an example, we take the analysis of \citet{Mall} of the clustering dipole of the 2-Micron All-Sky Survey (2MASS, \citealt{Skr}) ), since this is the densest and the deepest all-sky survey to date."
+" Malleretal.(2003). used the measurements in the near-infrared A; band the Zone of Avoidance as the regionwith [b for/and>>""n""230° vereor/«130 and |b]<12° for/>330 or/«30°,", \citet{Mall} used the measurements in the near-infrared $K_\mathrm{s}$ band and defined the Zone of Avoidance as the regionwith $|b|<7\degr$ for $l>230\degr$ or $l<130\degr$ and $|b|<12\degr$ for $l>330\degr$ or $l<30\degr$.
+ Let us a mean value’:malhrmvad., Let us take a `mean value':.
+. We now use Eq. (2?) , We now use Eq. \ref{eq:v_spurious}) )
+an for consistency with Tullyetal.(2008) we applyὃς Ou=024 and 2=16Alpe. to obtain an approximate value of the additional. spurious velocity of the Local Group. measured if the Local Void is not properly accounted for:," an for consistency with \citet{Tully.etal}, we apply, $\Omega_\mrm=0.24$ and $R=16\,\mathrm{Mpc}$, to obtain an approximate value of the additional, spurious velocity of the Local Group, measured if the Local Void is not properly accounted for:."
+ When compared to the velocity of the LG relative to the CMB reference Tame. roar=6)322232cr+35kms one can see that this effect is of the same order as the error in the measurement of roar.," When compared to the velocity of the LG relative to the CMB reference frame, $v_\mathrm{CMB}=622\pm35\;\kms$, one can see that this effect is of the same order as the error in the measurement of $v_\mathrm{CMB}$."
+ ote also that due to the proportionality of jiu to the radius of the _LV GEV. 1. ," Note also that due to the proportionality of $v_\mathrm{spur}$ to the radius of the LV (Eq. \ref{eq:v_spurious}) ),"
+increasing its diameter to 45Mpc. while preserving sphericity. woul cause! thé spufiods to raise significantly o 6Okms," increasing its diameter to $45\,\mathrm{Mpc}$, while preserving sphericity, would cause the spurious velocity to raise significantly to $\sim60\kms$."
+ However. owing to geometry of the problem. us would not largely affect general conclusions of our analysis (cf.," However, owing to geometry of the problem, this would not largely affect general conclusions of our analysis (cf."
+ subsequent Sections)., subsequent Sections).
+ Moreover. observational constraints poin rather to some degree of elongation of the LV than to a larger size of thes hole structure.," Moreover, observational constraints point rather to some degree of elongation of the LV than to a larger size of the whole structure."
+ We will now shortly address this issue., We will now shortly address this issue.
+ he eμου]ations so far assumed a simplistic model of a spherica void., The calculations so far assumed a simplistic model of a spherical void.
+" But we c""m1 sed example on fig.", But we can see for example on fig.
+ IO of Tullyetal.(2008) tha 18Local Void should be possibly modelled by a more sophisticated structure. like an ellipsoid.," 10 of \citet{Tully.etal} that theLocal Void should be possibly modelled by a more sophisticated structure, like an ellipsoid."
+ Current observational data suggest tha 1ο LV is elongated in the supergalactic SGY=0 plane. roughly coincident with the Galactic plane. hence our wr axis.," Current observational data suggest that the LV is elongated in the supergalactic $\mathrm{SGY}=0$ plane, roughly coincident with the Galactic plane, hence our $x$ axis."
+ One should ét 1Ἠ however that this effect could be a manifestation of the ZoA itself: lack of observed galaxies in this region of the sky may be simply due to obscuration., One should bear in mind however that this effect could be a manifestation of the ZoA itself: lack of observed galaxies in this region of the sky may be simply due to obscuration.
+Nevertheless. in case the is real. for completeness of our analysis let us examine thiselongation possibility.,"Nevertheless, in case the elongation is real, for completeness of our analysis let us examine this possibility."
+ For that purpose we assume that the section of the LV in our ry plane is an ellipse with semiaxes 6=2 anda =&Je. where elongation L.," For that purpose we assume that the section of the LV in our $xy$ plane is an ellipse with semiaxes $b\equiv R$ and $a\equiv\kappa\,R $, where elongation."
+. The major axis of the ellipse is placed along the + axis of the coordinate system., The major axis of the ellipse is placed along the $x$ axis of the coordinate system.
+ As in the spherical case. owing to symmetries of the problem. the only non-vanishing component of the spurious acceleration ‘acting’ on the LG is the . one.," As in the spherical case, owing to symmetries of the problem, the only non-vanishing component of the spurious acceleration `acting' on the LG is the $x$ one."
+" It is easy to check that the relevant formula for g, is now BE, with the integral Z,. given by en "," It is easy to check that the relevant formula for $g_x$ is now R, with the integral $\mathcal{I}_e$ given by _0 _0 ."
+Here.S(47)=1(iμ.μ... is a properly normalised equation. of the ellipse in polar coordinatesj with eccentricityHO.," Here, $\Xi(\varphi)\equiv1\slash\left(\kappa\sqrt{1-e^2 \cos^2\varphi}\right)$ is a properly normalised equation of the ellipse in polar coordinates with eccentricity."
+ The spurious velocity of the LG due to such elongation of the LV will be larger in comparison to the spherical case by a factor , The spurious velocity of the LG due to such elongation of the LV will be larger in comparison to the spherical case by a factor .
+Obviously. for &=1. we have 7—].," Obviously, for $\kappa=1$, we have $\tau=1$."
+ What is important here is that a linearincrease in A results only in a slower than linear raise of the 7 factor: for example #=2 gives 7=4/3 and if Bo= 3. then 7= 3/2. This means that for the observationallyallowed elongation of &zsC2 IZand minor semiaxis b215Alpe . the spurious velocity of the LV would raise by 1/3 to," What is important here is that a linearincrease in $\kappa$ results only in a slower than linear raise of the $\tau$ factor: for example $\kappa=2$ gives $\tau=4\slash 3$ and if $\kappa=3$ , then $\tau=3\slash 2$ This means that for the observationallyallowed elongation of $\kappa\simeq2$ and minor semiaxis $b\simeq15\Mpc$\citep{Tully.etal}, , the spurious velocity of the LV would raise by $1/3$ to"
+thirty-six observations of 22039 ive listed in Table 5.. and they are shown in Figure 5 along with a I&eplerian fit to the data with the orbital parameters listed in Table 2..,"thirty-six observations of 2039 are listed in Table \ref{vel2039}, and they are shown in Figure \ref{hd76700_rv_curve}
+ along with a Keplerian fit to the data with the orbital parameters listed in Table \ref{orbits1}."
+ The resultant wim companion mass for 776700 is 197TE0.017Ajup.. with an orbital semi-major axis e 00.0109480.00 Laan and eccentricity 620.00220.01.," The resultant minimum companion mass for 76700 is $\pm$, with an orbital semi-major axis $a$ $\pm$ au and eccentricity $e$ $\pm$ 0.04."
+. This zero ecceutricitv is consistent with the expectation that a planet with a period of just 23.9712E0.001 dd will almost certainly lic in au orbit which hasbeen tidally cireularised (AlarcyButler 1998)., This zero eccentricity is consistent with the expectation that a planet with a period of just $\pm$ d will almost certainly lie in an orbit which has been tidally circularised \citep{mb98}.
+". The resulting orbital parameters for 776700 place if amongst the plauctary companions with the lowest known minium masses,", The resulting orbital parameters for 76700 place it amongst the planetary companions with the lowest known minimum masses.
+ 77€τοῦ. jolus 119671. 1161HL. 1168716 116375Pepe (Butler2002) iu the group of extra-solu planetary companions with measured mundi masses less than a Saturni mass Mjup)).," 76700, joins 49674, 16141, 168746 46375 \citep{butler02b, mbv00, pepe02} in the group of extra-solar planetary companions with measured minimum masses less than a Saturn mass )."
+ The remaining three extra-solar plaucts all have elliptical orbits. though the ellipticity of the orbit for the colpanion to 330177 (¢=0.2240.17) is not differeut from zero with great significance.," The remaining three extra-solar planets all have elliptical orbits, though the ellipticity of the orbit for the companion to 30177 $e$ $\pm$ 0.17) is not different from zero with great significance."
+ As further data are acquired over the coming vears this parameter will become far better constraiued. and it is possible that this extra-solar planet couk turn out to be in a substantially circular orbit.," As further data are acquired over the coming years this parameter will become far better constrained, and it is possible that this extra-solar planet could turn out to be in a substantially circular orbit."
+ If so this system wouldjoin with the other known nearly circular systems with σαςο planets Iwineg iu orbits between where the Earth aud Jupiter lie in our own Solar System (e RRet. (1208. the outer components of UUMa. 228185 (Jonesetal. 2002a))).," If so this system would join with the other known nearly circular systems with gas-giant planets lying in orbits between where the Earth and Jupiter lie in our own Solar System $\epsilon$ Ret, 4208, the outer components of UMa, 28185 \citep{AAPSIV}) )."
+ These extra-solar planets would seen to indicate that eas-eiauts exist in nearly circular orbits with semi-major axes all tle wav out to. and bevoud that of Jupiter. as confirmed recently by the detection of the outer planet in the CCnc svsteii (Marcyetal.2002a).," These extra-solar planets would seem to indicate that gas-giants exist in nearly circular orbits with semi-major axes all the way out to, and beyond that of Jupiter, as confirmed recently by the detection of the outer planet in the Cnc system \citep{55Cnc}."
+. Of the twenty metal-enriched stars included. as a sub-salple along with our main sample iu late 1999. five have revealed the presence of plauetary companions (the four planets discussed bere 330177. 73526. 76700 and 2039 along with the previously known companion to IIDsS583115 (Butleretal2002a:Alavoretal.|2002)3)).," Of the twenty metal-enriched stars included as a sub-sample along with our main sample in late 1999, five have revealed the presence of planetary companions (the four planets discussed here – 30177, 73526, 76700 and 2039 – along with the previously known companion to 83443 \citep{AAPSV,mayor_83443}) )."
+" This gives us a lower nuit (there may be longer period or lower mass plauets present which we cannot vet detect) to the discovery rate of 52-11 for this ""1netallicity-biassed? sub-sauuple.", This gives us a lower limit (there may be longer period or lower mass planets present which we cannot yet detect) to the discovery rate of $\pm$ for this “metallicity-biassed” sub-sample.
+ This com2!ares with the overall disccun rates estimated for the Nec- Lick and AAPS of «54 Ge of stars surveyed. have pajets in orbits within oof their host stars (Butleretal.20020))) — a cüffereuce which. while uot of ercat statistical significauce. is nof unexpected given that extra-solar plaucts sccm to be being found prefercutially around iietal-euriclied stars (ce.," This compares with the overall discovery rates estimated for the Keck, Lick and AAPS of $\sim$ (ie of stars surveyed have planets in orbits within of their host stars \citep{properties}) ) – a difference which, while not of great statistical significance, is not unexpected given that extra-solar planets seem to be being found preferentially around metal-enriched stars (eg."
+ Rex(2002):Laughlin(2000) and references therein).," \citet{reid02,laughlin00} and references therein)."
+ It is also interesting to 1cte (even if perhaps not statistically significaut) that all four of the stars dm this paper wouk seen to be beginning their evolution off the main sequence (sce Fig 1j)., It is also interesting to note (even if perhaps not statistically significant) that all four of the stars in this paper would seem to be beginning their evolution off the main sequence (see Fig \ref{models}) ).
+ These results would suggest hat the biassiug of plane Survovs toward inetalouricle host stars offer a vcneht in the planet detectioji rate., These results would suggest that the biassing of planet surveys toward metal-enriched host stars offer a benefit in the planet detection rate.
+ However. such au Increase cüscoverv rate last be balanced against the fac hat it will produce an inherentvobiassed sample of extra-solar plajets;," However, such an increased discovery rate must be balanced against the fact that it will produce an inherently biassed sample of extra-solar planets."
+" With the total nuuber of extra-solar planets still wuerie less than a hundred. aud the parameters o his enseuble of planets still poorly placed iu a scheme of extra-solar planctary formation aud evolution. now is no he time for pluiet searches to beein biassiug their large surveys in the cliase for better ""lit rates” at the expense of scientific utility."," With the total number of extra-solar planets still numbering less than a hundred, and the parameters of this ensemble of planets still poorly placed in a scheme of extra-solar planetary formation and evolution, now is not the time for planet searches to begin biassing their large surveys in the chase for better “hit rates” at the expense of scientific utility."
+ The Auglo-Australian. Planet. Search team would like o eratefully acknowledge the support of Dr Brian Bovle. Director of the AAO. the coutiibutions aud assistauce of Drs Stuart Ryder and Debra Fischer. aud the superb echuical support which has been received throughout the xoeraunmue from AAT staff - in particular. R.Patterson. D.Stafford. F.Freeman. S.Lee. J.Posesonu. Colitlev aud J. Stevenson.," The Anglo-Australian Planet Search team would like to gratefully acknowledge the support of Dr Brian Boyle, Director of the AAO, the contributions and assistance of Drs Stuart Ryder and Debra Fischer, and the superb technical support which has been received throughout the programme from AAT staff - in particular R.Patterson, D.Stafford, F.Freeman, S.Lee, J.Pogson, G.Kitley and J. Stevenson."
+ We further acknowledge support by: the xutuers of the Auglo-Australiau Telescope Aerecimment (COTURASAIP): NASA evant NACH5-8299 NSF eraut AST95-20118 (CAWADT): NSF eraut. AST-0988087. (RPB): and Sun Microsystems.," We further acknowledge support by; the partners of the Anglo-Australian Telescope Agreement (CGT,HRAJ,AJP); NASA grant NAG5-8299 NSF grant AST95-20443 (GWM); NSF grant AST-9988087 (RPB); and Sun Microsystems."
+ NSO/Kitt Peak FTS data used here were produced by NSF/NOAO., NSO/Kitt Peak FTS data used here were produced by NSF/NOAO.
+ The formation of the Calactic halo is preseutlv at the ceuter of: an open debate., The formation of the Galactic halo is presently at the center of an open debate.
+ Stetson. et al. ] (1996)), Stetson et al. \cite{s_96}) )
+ state that. apart from a handful of anomalous clusters that may well have been captured from a satellite dw galaxy. there is no strong evidence for a significant spread in age amoue clusters of a eiven mctal abundance. while Clabover et al. (1996))," state that, apart from a handful of anomalous clusters that may well have been captured from a satellite dwarf galaxy, there is no strong evidence for a significant spread in age among clusters of a given metal abundance, while Chaboyer et al. \cite{c_96}) )"
+ support an age spread of: 5 Cyr among the bulk of. the Calactictale eeloular clusters (GOCs) (which is increased to 9 Cyr. ↽∎⋅∙if the vouugestm clusters mEare considered).," support an age spread of 5 Gyr among the bulk of the Galactic globular clusters (GGCs) (which is increased to 9 Gyr, if the youngest clusters are considered)."
+" One of the largest underlying sources of: uncertaiutv. is the heterogeneity of the data used in these studies. which prevents ""large scale” tests."," One of the largest underlying sources of uncertainty is the heterogeneity of the data used in these studies, which prevents “large scale” tests."
+ This is the main reason that prompted our eroup to gather an homogeneous photometric data base of GC. in V. aud 2. as discussed iu Saviane et al. (1997)).," This is the main reason that prompted our group to gather an homogeneous photometric data base of GGC, in $V$ and $I$, as discussed in Saviane et al. \cite{s97}) )."
+" To date. we have observed about ofthe closest GC's (G0.M),x 16) with 111 class telescopes. aud our data set allows us to obtain color-magnitude diagrams (CMDsx) from the RGB tip down to a few magnitudes below the turn-off (TO)."," To date, we have observed about of the closest GGC's $(m-M)_v<16$ ) with 1m class telescopes, and our data set allows us to obtain color-magnitude diagrams (CMDs) from the RGB tip down to a few magnitudes below the turn-off (TO)."
+ Several vouug (or suspect vouug) GCs were iucluded in ow program: Palomar 12 (Catton Ortolani 1988.. hereafter COSS). Ruprecht 106 (Buonanno et al.1990)). Àrp 2 (Buouauno et al. L995a)}).," Several young (or suspect young) GGCs were included in our program: Palomar 12 (Gratton Ortolani \cite{g88}, hereafter GO88), Ruprecht 106 (Buonanno et \cite{b90}) ), Arp 2 (Buonanno et al. \cite{b95a}) ),"
+ Terzan 7 (Buonauno ct al. —.19955)), Terzan 7 (Buonanno et al. \cite{b95b}) )
+ and Palomar 1 (Roseubere et al. 1998..," and Palomar 1 (Rosenberg et al. \cite{r98},"
+ hereafter Paper D., hereafter Paper I).
+ Their age is suspected to sieuificautlv deviate from the general distribution., Their age is suspected to significantly deviate from the general distribution.
+ A precise determination of this deviation within our homogeneous data set is particularly valuable in the general framework of the COC ages aud of ercat muportauce iu order to decide on the aimodels of Calactic formation., A precise determination of this deviation within our homogeneous data set is particularly valuable in the general framework of the GGC ages and of great importance in order to decide on the models of Galactic formation.
+ For this reason. we will present and discuss in separate papers of this series the photometric data for the clusters whose age is significantly different frou the average age of the CCC in our sample.," For this reason, we will present and discuss in separate papers of this series the photometric data for the clusters whose age is significantly different from the average age of the GGC in our sample."
+ Tn Paper I we have discussed the case of Palomar | which resulted to be the voungest CCC in our Galaxy.," In Paper I we have discussed the case of Palomar 1, which resulted to be the youngest GGC in our Galaxy."
+" Iu this paper we conceutrate on Pal 12 (C2113-21 ο=21""1626. 3)=21715: l=3075. b= 1727) discovered bv Warrington Zwicky (1953)) on the Palomar Sky. Survey plates"," In this paper we concentrate on Pal 12 (C2143-214, $\alpha_{\rm 2000}=21^h\,46\fm6$, $\delta_{\rm
+2000}=-21^\circ\,15'$; $=30\fdg5$, $=-47\fdg7$ ) discovered by Harrington Zwicky \cite{h53}) ) on the Palomar Sky Survey plates."
+ .Ivdeed. Pal 12 has Όσοι +ιο first cluster to be classified as vounger than the bulls EÉGOGCS," Indeed, Pal 12 has been the first cluster to be classified as younger than the bulk of GGCs."
+ Iowever. the age «eteriunation di previous PAuclies of its CAID. namely Stetπο ret al. (1989..," However, the age determination in previous studies of its CMD, namely Stetson et al. \cite{s89},"
+ hereafter οκ.89). aud Da Costa Axiiancdi«Xt (1990.. hereafter DAÀ90) was affected by large uucertaiities dn its metal content nore than 0.3 dex in [Fe/TI]).," hereafter S89), and Da Costa Armandroff \cite{d90}, hereafter DA90) was affected by large uncertainties in its metal content (more than 0.3 dex in [Fe/H])."
+ Siice then. new low- aud ligh-resolution spectroscopy have been used iu order to estimate⋅ the metallicity⋅⋅ of. Paloliavy 12 (Armandroft Da Costa ↽↽1991.. AD91I:↽ Brown et a 1997..," Since then, new low- and high-resolution spectroscopy have been used in order to estimate the metallicity of Palomar 12 (Armandroff Da Costa \cite{a91}, AD91; Brown et al. \cite{b97},"
+ BOT)., B97).
+ The observations and data recuction are presented iu Sect., The observations and data reduction are presented in Sect.
+ 2 and the resulting CMD is discussed in Sect. 3.., \ref{obs_reduct} and the resulting CMD is discussed in Sect. \ref{sec_cmd}.
+ The relative ageOo determination Is carried out in Sect. [L.," The relative age determination is carried out in Sect. \ref{age},"
+ and in Sect., and in Sect.
+5- the structural paraneters of the cluster are established., \ref{sec_struct} the structural parameters of the cluster are established.
+ The data were collected ou October 20. 1997 at the ESO Danish 1.5lin telescope equipped with DFOSC.," The data were collected on October 20, 1997 at the ESO Danish 1.54m telescope equipped with DFOSC."
+ The canacra enployed a 20182015 pixels Loral CCD. with a pixel size of 07110 on the sky. for a total field of view of 13/6« 13/6.," The camera employed a $2048\times2048$ pixels Loral CCD, with a pixel size of 40 on the sky, for a total field of view of $13\farcm6\times 13\farcm 6$ ."
+ Llere Ay is the innermost stable orbit of the disk.,	 Here $R_{\rm i}$ is the innermost stable orbit of the disk.
+ Photons emitted by the disk may enter the jet and scatter with the jet electrons (see Fig. 2))., Photons emitted by the disk may enter the jet and scatter with the jet electrons (see Fig. \ref{emissiontype}) ).
+ We also investigate the scattering of CAIB blackbody photons in the jet., We also investigate the scattering of CMB blackbody photons in the jet.
+ The seed CALB photons are selected from. a single-temperature blackbody., The seed CMB photons are selected from a single-temperature blackbody.
+ In this work we consider jets at redshift of z=2., In this work we consider jets at redshift of $z=2$.
+ We note that the results can be eeneralized to higher or lower red-shifts simply by acjusting the temperatures of the injected blackbody photons., We note that the results can be generalized to higher or lower red-shifts simply by adjusting the temperatures of the injected blackbody photons.
+ The cillerential Klein-Nishina cross-section for scattering of polarized electromagnetic waves. with polarization vector e. olf electrons in the electron rest frame is (seeLleitlerJauch&Rohrlich1980.for example). Here. ry is the classical electron radius and The subscript € denotes quantities calculated in the electron rest frame and primed quantities denote values after scattering.," The differential Klein-Nishina cross-section for scattering of polarized electromagnetic waves, with polarization vector $\mathbf{e}$, off electrons in the electron rest frame is \citep[see][for example]{Heitler36, Jauch80}, Here, $r_{0}$ is the classical electron radius and The subscript 'e' denotes quantities calculated in the electron rest frame and primed quantities denote values after scattering."
+ Here. £& is the frequency of the incident. photon. mois the mass of an electron. 6 is the scattering anele (angle between the incident and scattered photon wavevector) and Ó is the azimuthal angle.," Here, $\nu$ is the frequency of the incident photon, $m_{\rm e}$ is the mass of an electron, $\theta$ is the scattering angle (angle between the incident and scattered photon wavevector) and $\phi$ is the azimuthal angle."
+ The polarization vector e for the polarized fraction 2 ol scattered photons is perpendicular to the scattering plane and is defined by (Anecl1969).. where QF2 is the propagation. directional. unit.. vector of the scattered. photon.," The polarization vector $\mathbf{e}\,'$ for the polarized fraction $P$ of scattered photons is perpendicular to the scattering plane and is defined by \citep{Angel69}, where $\hat{\mathbf{\Omega}}'$ is the propagation directional unit vector of the scattered photon."
+ For the remaining 1P fraction of photons e' is randomly distributed. in. the plane perpendicular to Q’.," For the remaining $1-P$ fraction of photons $\mathbf{e}\,'$ is randomly distributed in the plane perpendicular to $\hat{\mathbf{\Omega}}'$."
+ 1n. the case where the radiation is intially unpolarized. the degree. of linear. polarization induced hy Compton scattering is given by (seee.g.Dolan1967:Mattetal. 1996).. For polarized incident emission. the degree of linear polarization induced is given by. details can be found. in AleNamara οἱal.(2008a).," In the case where the radiation is intially unpolarized, the degree of linear polarization induced by Compton scattering is given by \citep[see e.g.][]{Dolan67, Matt96}, , For polarized incident emission, the degree of linear polarization induced is given by, Further details can be found in \cite{McNamara08a}."
+. νου clifferent reference frames are used in the computation: the plasma rest frame (PLE) which is the [rame co-moving with the emitting jet plasma: the electron rest [rame (EE) and the observer rest frame (OL)., Three different reference frames are used in the computation: the plasma rest frame (PF) which is the frame co-moving with the emitting jet plasma; the electron rest frame (EF) and the observer rest frame (OF).
+ In the EF the z-axis is defined by the direction. of the electron propagation @.. while in the PE the z-axis is along the magnetic field of the jet.," In the EF the $z$ -axis is defined by the direction of the electron propagation $\hat{\bbeta}_{\rm e}$, while in the PF the $z$ -axis is along the magnetic field of the jet."
+ In the case where the PF and the EF are not in a standard configuration (i.e. the EF does not move along the z-direction of the PIE and the corresponding axes of the two frames are not parallel) a new rotated frame in standard configuration is first defined in the PE and then transformations between frames are performed., In the case where the PF and the EF are not in a standard configuration (i.e. the EF does not move along the $z$ -direction of the PF and the corresponding axes of the two frames are not parallel) a new rotated frame in standard configuration is first defined in the PF and then transformations between frames are performed.
+ We model. the scattering of three. dilferent types of photons in our analysis: unpolarizecl blackhocky disk photons. unpolarized blackbody CMD. photons and intrinsically polarized svnchrotron photons.," We model the scattering of three different types of photons in our analysis: unpolarized blackbody disk photons, unpolarized blackbody CMB photons and intrinsically polarized synchrotron photons."
+ The disk photons are emitted in the OF from a multicolour disk model., The disk photons are emitted in the OF from a multicolour disk model.
+ We. assume that. the photons are. intially unpolarized., We assume that the photons are intially unpolarized.
+ Ehe probability density of a disk seed photon can be written as. where £(R2) is the Hux density at disk racius A? (see equation 4)) and L is the total luminosity of the disk.," The probability density of a disk seed photon can be written as, where $F(R)$ is the flux density at disk radius $R$ (see equation \ref{Flux}) ) and $L$ is the total luminosity of the disk."
+ We sample p(4?) w tabulation ancl determine the blackbody temperature at A., We sample $p(R)$ by tabulation and determine the blackbody temperature at $R$.
+ Phe photon energy is then sampled. from the Plank unction DH) as described in Pozdnvakov., The photon energy is then sampled from the Plank function $B_{\nu}(R)$ as described in \citet{P83}.
+Sobol&Sun- Some of the disk photons may impact the jet and. undergo scattering. these photons are transported hrough the jet until they escape.," Some of the disk photons may impact the jet and undergo scattering, these photons are transported through the jet until they escape."
+ CAIB photons may. also undergo scattering within the jet (see Fig., CMB photons may also undergo scattering within the jet (see Fig.
+ 2. (a))., \ref{emissiontype} (a)).
+ We sample the CMD. photons from a Xackbody. distribution with à temperature >) Ik in the OF at a redshift z.," We sample the CMB photons from a blackbody distribution with a temperature $T_{\rm bb} = 2.8 \, (1+z)$ K in the OF at a redshift $z$ ."
+ The initially unpolarized photons are emitted isotropically around the jet., The initially unpolarized photons are emitted isotropically around the jet.
+ H£ a photon enters the jet it may scatter with the jet. electrons. and ome polarized., If a photon enters the jet it may scatter with the jet electrons and become polarized.
+" Synchrotron radiation is emitted. by the jet electrons with electron Lorentz [actors ,/55ROAway.", Synchrotron radiation is emitted by the jet electrons with electron Lorentz factors $\gamma_{min} \leq \gamma \leq \gamma_{\rm max}$.
+ We emit the photons isotropically in the Ek with an initial enerey chosen from a powerlaw distribution., We emit the photons isotropically in the EF with an initial energy chosen from a powerlaw distribution.
+ Svachrotron radiation is polarized and depends on the magnetic field B of the jet. thus the initial polarization vector e is chosen in the PE so that it is perpendicular to both B and €.," Synchrotron radiation is polarized and depends on the magnetic field $\mathbf{B}$ of the jet, thus the initial polarization vector $\mathbf{e}$ is chosen in the PF so that it is perpendicular to both $\bmath{B}$ and $\bmath{\Omega}$."
+ We Lorentz transform the svnchrotron photons from the EF to the PE using. and where the Doppler factor à.=[5.1100⋅BO]+ and ‘pl ancl ο denote quantities in the plasma and. electron rest frame. respectively.," We Lorentz transform the synchrotron photons from the EF to the PF using, 	 and where the Doppler factor $\delta_{\rm e} = \{ \gamma_{\rm e}[1 - \beta_{\rm e}(\hat{\bmath{\Omega}}\cdot\hat{\bmath{\beta}}_{\rm e})]\}^{-1}$ and 'pf' and 'ef' denote quantities in the plasma and electron rest frame, respectively."
+ Once the incident. photon energy. direction and polarization have been selected. the photon is transported. through the jetand the distance the photon has travelled is determined.," Once the incident photon energy, direction and polarization have been selected, the photon is transported through the jetand the distance the photon has travelled is determined."
+ Η the photon encounters an electron beforeleaving the jet.," If the photon encounters an electron beforeleaving the jet,"
+by using a 3D radiative transfer code (o.g.. IxvIalis Baheall 1987. Nilouris et al.,"by using a 3D radiative transfer code (e.g., Kylafis Bahcall 1987, Xilouris et al."
+ 1997). we choose to mask the data in the regions near the planes instead.," 1997), we choose to mask the data in the regions near the planes instead."
+ Thirdly. by adopting our simple model. Eq. (," Thirdly, by adopting our simple model, Eq. ("
+4). we do not include the effects of a truncated. galactic disc along the line of sight.,"4), we do not include the effects of a truncated galactic disc along the line of sight."
+ However. the elfects of neglecting a line-obsieht truncation are small (see Ixregel et al.," However, the effects of neglecting a line-of-sight truncation are small (see Kregel et al."
+ 2001): while it may cause us to underestimate the cise scalelength of ESO4AL6-G44 by ~15%... the ellect is SUA for the other galaxies in our sample. and is in practice counteracted by residual extinction. by dust. if anv. at the z-heights. used in our studs.," 2001): while it may cause us to underestimate the disc scalelength of ESO446-G44 by $\sim 15$, the effect is $\le 8$ for the other galaxies in our sample, and is in practice counteracted by residual extinction by dust, if any, at the -heights used in our study."
+ In addition. this effect is negligible for the accurate determination of the cise as such. the main aim of our studv.," In addition, this effect is negligible for the accurate determination of the disc as such, the main aim of our study."
+ 3elore applving our fitting routine to the full observed luminosity density distributions. the skv-subtracted images needed to be. prepared.," Before applying our fitting routine to the full observed luminosity density distributions, the sky-subtracted images needed to be prepared."
+ First. foreground. stars. and background. galaxies were masked out.," First, foreground stars and background galaxies were masked out."
+ Secondlv. profiles parallel to the minor axis were taken at various clistances along the major axis.," Secondly, profiles parallel to the minor axis were taken at various distances along the major axis."
+ These were subsequentIy inspected for extinction ellects. leading to the masking out of the regions described below for the individual galaxies.," These were subsequently inspected for extinction effects, leading to the masking out of the regions described below for the individual galaxies."
+ We will now discuss the fits to the luminosity density of cach galaxy individually. thereby addressing a number of problems encountered. in cach case.," We will now discuss the fits to the luminosity density of each galaxy individually, thereby addressing a number of problems encountered in each case."
+ In. all fits the galactic centres were fixed at the values determined by fitting ellipses, In all fits the galactic centres were fixed at the values determined by fitting ellipses
+The calculated evolutionary tracks in the IIR diagram are plotted in Fie.,The calculated evolutionary tracks in the HR diagram are plotted in Fig.
+ L., 4.
+ The bold curves denote the phases in which the orbital period decreases., The bold curves denote the phases in which the orbital period decreases.
+ For comparison with observations. we also show the locations of four Aleol svste1is which have decreasing orbital periods aud woell-iicasurced paralucters sunmnuarized in Table 1.," For comparison with observations, we also show the locations of four Algol systems which have decreasing orbital periods and well-measured parameters summarized in Table 1."
+ As shown in the fleure. the calculated results are roughly in accordance with the observations. sugeesting that they may represeut the possible evolutionary sequences that lead to the formation of these systenis.," As shown in the figure, the calculated results are roughly in accordance with the observations, suggesting that they may represent the possible evolutionary sequences that lead to the formation of these systems."
+ Fieure 5 compares the calculated results of Aleol binary evolution with the observational data in the TigDas plane., Figure 5 compares the calculated results of Algol binary evolution with the observational data in the $T_{\rm eff} - P_{\rm orb}$ plane.
+ It seems that our model is also compatible with the observed spectral types of the four svstenis., It seems that our model is also compatible with the observed spectral types of the four systems.
+ It is au open question why some of the Algol systems show orbital shrinkage while others not., It is an open question why some of the Algol systems show orbital shrinkage while others not.
+ Although it is ogenerally believed. that there must be orbital augular momentum loss durius the evolution of these binaries. the appropriate mechanisia(s) for angular moment loss has not been verified.," Although it is generally believed that there must be orbital angular momentum loss during the evolution of these binaries, the appropriate mechanism(s) for angular momentum loss has not been verified."
+ Based on theoretical niodoel of Spruit&Taam(2001) and Taam&Spruit(2001).. we Sugeest that a CD disk may play au müiportaut role in determining the orbital change in Aleol systems.," Based on theoretical model of \citet{spru01} and \citet{taam01}, we suggest that a CB disk may play an important role in determining the orbital change in Algol systems."
+ Note that our CD disk lypotlesis docs rot exclude the effect of Guaenuetized) mass loss., Note that our CB disk hypothesis does not exclude the effect of (magnetized) mass loss.
+ On the contrary. the formation ofthe CD disk may )»e Closely related to the mass loss processes in the duuv evolution.," On the contrary, the formation of the CB disk may be closely related to the mass loss processes in the binary evolution."
+ The existence of the CD. disks nay be revealed. by their infrared radiation. as nn Gc Tau (Roddieretal. 1996)..," The existence of the CB disks may be revealed by their infrared radiation, as in GG Tau \citep{rodd96}. ."
+ The Monte-C'urlo simulation by Woodetal.(1999). has reproduced he morphology of the near IR radiation which is in accord with the observations by Rodcdieretal. (1996)., The Monte-Carlo simulation by \citet{wood99} has reproduced the morphology of the near IR radiation which is in accord with the observations by \citet{rodd96}.
+. Iu this paper we have calculated the evolutionary sequence for Aleol biuaries with typical masses and orbital periods. takine iuto account both maenetic braking aud the CD disk.," In this paper we have calculated the evolutionary sequence for Algol binaries with typical masses and orbital periods, taking into account both magnetic braking and the CB disk."
+ The detailed calculations indicate that the orbital evolution of Aleol binaries can be significantly affected by the CB disk. especially for Case A mass transfer.," The detailed calculations indicate that the orbital evolution of Algol binaries can be significantly affected by the CB disk, especially for Case A mass transfer."
+ With adequate values of d. it is possible to account for the decrease of the orbital periods iu some Aleol binimries.," With adequate values of $\delta$, it is possible to account for the decrease of the orbital periods in some Algol binaries."
+ Di addition.the existence of CB," In addition,the existence of CB"
+"For the clusters in the distant samples. we perform a simple spectral analysis extracting the spectrum of each source from a circular region with radius Rk,= 4O""centered. around the X-ray centroid.","For the clusters in the distant samples, we perform a simple spectral analysis extracting the spectrum of each source from a circular region with radius $R_{ext} = 40$ centered around the X-ray centroid."
+ The background is always obtained from empty regions of the chip in which the source ts located., The background is always obtained from empty regions of the chip in which the source is located.
+ This is possible since all the distant clusters have an extension of less than 34 aremin (corresponding to a range in physical size of 1100-1500 kpe). as opposed to the 8 aremin size of the ACIS-I/-S chips.," This is possible since all the distant clusters have an extension of less than 3 arcmin (corresponding to a range in physical size of 1100-1500 kpc), as opposed to the 8 arcmin size of the ACIS-I/-S chips."
+ The background file is scaled to the source file by the ratio of the geometrical area., The background file is scaled to the source file by the ratio of the geometrical area.
+ In principle. the background regions may partially overlap with the outer virialized regions of the clusters.," In principle, the background regions may partially overlap with the outer virialized regions of the clusters."
+ However. the cluster emission from these regions is negligible with respect to the instrumental background. and does not affect our results.," However, the cluster emission from these regions is negligible with respect to the instrumental background, and does not affect our results."
+ Our background subtraction procedure. on the other hand. has the advantage of providing the best estimate of the background for that specific," Our background subtraction procedure, on the other hand, has the advantage of providing the best estimate of the background for that specific"
+ (seereviewsFan2009:Charbonneau2010).. (c.g...Guerrero&Ixapylaà2011). (¢.¢..Nersaleetal.2007)... (seereviewMiesch2005).," \cite[see reviews][]{Fan09, Cha10}. \cite[e.g.,][]{Cal95, Fan08, Jou09, Web11}. \cite[e.g.,][]{Cli03, Vas08, Gue11}. \cite[e.g.,][]{Ker07}. \cite[see review][]{Mie05}."
+. Such models have captured the formation of magnetic structures aud cycles ii solw (Chizaructal.2010:Racineet2011) and rapidlv-rotatius sun-like stellar models (Brown=ctal.2010:Brown—2011.hereinafterBloaudDll. respoectivelv). viekliug differeutial rotation. dvuamo action. aud large-scale maguctic fields. but not buovant magnetic loops that rise toward the top of the convective laver.," Such models have captured the formation of magnetic structures and cycles in solar \cite{Ghi10, Rac11} and rapidly-rotating sun-like stellar models \cite[][hereinafter B10 and B11, respectively]{Bro10, Bro11b}, yielding differential rotation, dynamo action, and large-scale magnetic fields, but not buoyant magnetic loops that rise toward the top of the convective layer."
+ Here we report on a elobal couvective dynamo siuulation of a sun-like star rotating at thice times the mean solar angular velocity. (30... such as our Sun did when it was younger aud as do many solar analogues (Petitetal.2008).," Here we report on a global convective dynamo simulation of a sun-like star rotating at three times the mean solar angular velocity $\Omega_\odot$ ), such as our Sun did when it was younger and as do many solar analogues \cite{Pet08}."
+. This simulation (i) attains a differential rotation profile created by the interplay of couvectiou. rotation and stratification (e...Bran&Toomre2002:AlieschToonie2009).. (i) forms global-scale toroidal maguctic structures that undergo cvcles of magnetic activity and reversals of elobal polarity. aud (i) achieves buovant maguctic loops from the strongest portions of the toroidal structures which rise from the base of the convection zone.," This simulation (i) attains a differential rotation profile created by the interplay of convection, rotation and stratification \cite[e.g.,][]{Bru02, Mie09}, (ii) forms global-scale toroidal magnetic structures that undergo cycles of magnetic activity and reversals of global polarity, and (iii) achieves buoyant magnetic loops from the strongest portions of the toroidal structures which rise from the base of the convection zone."
+ This work extends the work of 010 aud Bll in which simulations of rapidlv-rotating suus with moderate levels of diffusion were able to accomplish (4) and Gi)., This work extends the work of B10 and B11 in which simulations of rapidly-rotating suns with moderate levels of diffusion were able to accomplish (i) and (ii).
+ The formation of buovaut loops is facilitated in our current work bv adopting a dvuamic Simagoriusky siberid-scale model (Cteriuauoetal.1991).. which serves to minimize the diffusion of welleresolved structures.," The formation of buoyant loops is facilitated in our current work by adopting a dynamic Smagorinsky subgrid-scale model \cite{Ger91}, which serves to minimize the diffusion of well-resolved structures."
+" We have conducted 3-D MITD simulations of turbulent convection aud dynamo action in a spherical shell spamming the bulk of the convection zone from 0.722, to (.97R.. involving a density contrast of 25. aud rotating at ο, (1210 ulIlz.ouceevery 9.3 days)."," We have conducted 3-D MHD simulations of turbulent convection and dynamo action in a spherical shell spanning the bulk of the convection zone from $R_\odot$ to $R_\odot$ involving a density contrast of 25, and rotating at $\Omega_\odot$ (1240 nHz,onceevery 9.3 days)."
+ We use the auclastic προΊσα. harmonic (ASTD) code, We use the anelastic spherical harmonic (ASH) code
+star-forming.,star-forming.
+ We calculate the median values of SER and of SSFR. and the 25 and 75 percentiles.," We calculate the median values of SFR and of SSFR, and the 25 and 75 percentiles."
+ From Fig. 1..," From Fig. \ref{SFR},"
+ we note a change with redshift. in both the field and clusters. as the average SFRs at z~0.7 are shifted to higher values compared to z0.5 at the same mass.," we note a change with redshift, in both the field and clusters, as the average SFRs at $z\sim 0.7$ are shifted to higher values compared to $z \sim 0.5$ at the same mass."
+ The main result of Fig., The main result of Fig.
+ 1. is that. at both redshifts. cluster SFRs are on average systematically lower than field SFRs at the same mass.," \ref{SFR} is that, at both redshifts, cluster SFRs are on average systematically lower than field SFRs at the same mass."
+ Clusters have in general lower median SFRs than the field., Clusters have in general lower median SFRs than the field.
+" This ts due to a population of cluster galaxies lying below the field 25 percentile. that represent of the whole cluster population. thus à ~10% excess of galaxies with ""reduced"" SFR for their mass."," This is due to a population of cluster galaxies lying below the field 25 percentile, that represent of the whole cluster population, thus a $\sim 10$ excess of galaxies with “reduced” SFR for their mass."
+ To avoid the influence of the mass distribution. we performed [000 Monte Carlo simulations extracting randomly from the field sample a subsample with— the same mass distribution as the clusters.," To avoid the influence of the mass distribution, we performed 1000 Monte Carlo simulations extracting randomly from the field sample a subsample with the same mass distribution as the clusters."
+ The SFR distributions are shown as histograms in Fig. l.., The SFR distributions are shown as histograms in Fig. \ref{SFR}.
+ At z«0.6. due to poor numberstatistics.. a Kolmogorov-Smirnov test cannot reject the null hypothesis of similar SER distributions in. clusters and field. finding a probability less than in of the cases and a probability >95% in of the simulations.," At $z<0.6$, due to poor number, a Kolmogorov-Smirnov test cannot reject the null hypothesis of similar SFR distributions in clusters and field, finding a probability less than in of the cases and a probability $>95\%$ in of the simulations."
+ At zo»0.6. the K-S test rejects the hypothesis of similar cluster and field distributions with a 7955€ probability in of the simulations.," At $z>0.6$, the K-S test rejects the hypothesis of similar cluster and field distributions with a $>$ probability in of the simulations."
+ We note that. although a correlation is evident when considering galaxies over a wide mass range (see bottom windows in Fig. D) ," We note that, although a correlation is evident when considering galaxies over a wide mass range (see bottom windows in Fig. \ref{SFR}) ),"
+the SFR-Mass relation ts flat above our mass limit., the SFR-Mass relation is flat above our mass limit.
+ A Spearman test yields a significant positive correlation only for the field at z>0.6 (99.9%)). and no correlation in all other cases.," A Spearman test yields a significant positive correlation only for the field at $z>0.6$ ), and no correlation in all other cases."
+ In Fig. 2..," In Fig. \ref{SFR2},"
+ we show the results considering also REL galaxies as star-forming., we show the results considering also REL galaxies as star-forming.
+ We recall that at least for some of them the [OIL] emission likely arises from ongoing SE., We recall that at least for some of them the [OII] emission likely arises from ongoing SF.
+ Now the difference between field and clusters is more striking. and becomes progressively more pronounced towards more massive galaxies.," Now the difference between field and clusters is more striking, and becomes progressively more pronounced towards more massive galaxies."
+" of the whole cluster population have SFRs below the field 25 percentile. therefore ~25% have ""reduced"" SFRs for their mass compared to the field."," of the whole cluster population have SFRs below the field 25 percentile, therefore $\sim25$ have “reduced” SFRs for their mass compared to the field."
+ The K-S test on mass-matched cluster and field samples rejects the null hypothesis of similarity between the two environments with a probability of (z« 0.6) and >95% (z> 0.6) in all the simulations (see histograms in Fig. 2))., The K-S test on mass-matched cluster and field samples rejects the null hypothesis of similarity between the two environments with a probability of $z<0.6$ ) and $>95\%$ $z>0.6$ ) in all the simulations (see histograms in Fig. \ref{SFR2}) ).
+ As before. no SFR-Mass correlation is detected by a Spearman test. except for the field at z>0.6 at99.," As before, no SFR-Mass correlation is detected by a Spearman test, except for the field at $z>0.6$ at."
+967. Our results highlight a change in the SFR-Mass relation with environment., Our results highlight a change in the SFR-Mass relation with environment.
+ To quantify this change. we compute the mean SFR in our mass-matched simulations. therefore removing the effects of different mass distributions.," To quantify this change, we compute the mean SFR in our mass-matched simulations, therefore removing the effects of different mass distributions."
+" Including both redshift bins. for the sample of 24;/m*BEL galaxies the mean SFR in clusters is 1,350.15 times lower than in the field."," Including both redshift bins, for the sample of $\micron$ +BEL galaxies the mean SFR in clusters is $\pm$ 0.15 times lower than in the field."
+ Including REL galaxies. it is 1.63--0.20 times lower than the In Figs.," Including REL galaxies, it is $\pm$ 0.20 times lower than the In Figs."
+ 3. and 4. we show the results for the SSER- relation.," \ref{SSFR} and \ref{SSFR2}, , we show the results for the SSFR-Mass relation."
+ Cluster and field galaxies follow a qualitatively similar decreasing trend of SSFR with mass (Spearman anticorrelation probability always > 99.9%). but cluster galaxies tend to have a lower SSPR than field galaxies of similar mass. as expected from the results discussed above.," Cluster and field galaxies follow a qualitatively similar decreasing trend of SSFR with mass (Spearman anticorrelation probability always $>99.9\%$ ), but cluster galaxies tend to have a lower SSFR than field galaxies of similar mass, as expected from the results discussed above."
+ A K-S test confirms this when red galaxies are included (probabilities always >95%)., A K-S test confirms this when red galaxies are included (probabilities always $>95\%$).
+ For 24;m-+BEL galaxies. the differences are not statistically significant. with the K-S test giving a probability =95% only in of the simulations at zo»0.6. and only in of simulations at z«0.6.," For $\micron$ +BEL galaxies, the differences are not statistically significant, with the K-S test giving a probability $>95\%$ only in of the simulations at $z>0.6$, and only in of simulations at $z<0.6$."
+ Including both redshift bins. the ratio of average SSFR between field and clusters above our mass limit is 1.20--0.14 for 24j/m-BEL galaxies. and 1.31-£0.17 for all galaxies.," Including both redshift bins, the ratio of average SSFR between field and clusters above our mass limit is $\pm$ 0.14 for $\micron$ +BEL galaxies, and $\pm$ 0.17 for all galaxies."
+ This result shows that. in all environments. the mass growth rate at a given mass decreases with time (ef.," This result shows that, in all environments, the mass growth rate at a given mass decreases with time (cf."
+ left and right panel of Figs., left and right panel of Figs.
+ 3. and 4)) and it is lower for higher mass galaxies (SSFR and mass are anticorrelated)., \ref{SSFR} and \ref{SSFR2}) ) and it is lower for higher mass galaxies (SSFR and mass are anticorrelated).
+ However. a fraction of the star-forming cluster galaxies are building-up their stellar mass at a lower rate than field galaxies: and of 24;:m+BEL and all galaxies. respectively. lie below the field 25 percentile.," However, a fraction of the star-forming cluster galaxies are building-up their stellar mass at a lower rate than field galaxies: and of $\micron$ +BEL and all galaxies, respectively, lie below the field 25 percentile."
+ Moreover. the cluster trends are steeper than the field trends (best-fit slopes differ by >Io). again suggesting that SF in more massive galaxies differs more strongly with environment than SF in lower mass galaxies.," Moreover, the cluster trends are steeper than the field trends (best-fit slopes differ by $>1\sigma$ ), again suggesting that SF in more massive galaxies differs more strongly with environment than SF in lower mass galaxies."
+ With the aim of investigating whether cluster and group galaxies. separately. differ from the field. we divide the sample into clusters with σ>400kms7! and groups with c<400kms7! (Fig. 5).," With the aim of investigating whether cluster and group galaxies, separately, differ from the field, we divide the sample into clusters with $\sigma > 400 \rm \, km \, s^{-1}$ and groups with $\sigma < 400 \rm \, km \, s^{-1}$ (Fig. \ref{CGF}) )."
+ We only consider the highest-z bin. as the lowest-z bin has too few group galaxies to study any trend.," We only consider the highest-z bin, as the lowest-z bin has too few group galaxies to study any trend."
+ Figure 5. shows that the SF in the cluster environment deviates from the field trend. while group galaxies seem to follow the SFR-Mass relation of the field.," Figure \ref{CGF} shows that the SF in the cluster environment deviates from the field trend, while group galaxies seem to follow the SFR-Mass relation of the field."
+ The K-S test perfromed on mass-matched samples yields a (without red galaxies) and a (with red galaxies) probability that clusters have a different SFR distribution from groups and field., The K-S test perfromed on mass-matched samples yields a (without red galaxies) and a (with red galaxies) probability that clusters have a different SFR distribution from groups and field.
+ In contrast. the test cannot reject the hypothesis that the groups and field have a similar distribution.," In contrast, the test cannot reject the hypothesis that the groups and field have a similar distribution."
+ Our group data are not sufficient to draw firm conclusions. but. if confirmed. our finding suggests that the group environment is not influential for the link between SF activity and mass. therefore strangulation could not be relevant. at least in groups. and only cluster-specifie processes could be important.," Our group data are not sufficient to draw firm conclusions, but, if confirmed, our finding suggests that the group environment is not influential for the link between SF activity and mass, therefore strangulation could not be relevant, at least in groups, and only cluster-specific processes could be important."
+ Having removed the groups. the mean SFR in clusters at zo»0.6 is 1.93+0.02 (without red) to 2.13+0.02 (with. red) times lower than in the field.," Having removed the groups, the mean SFR in clusters at $z>0.6$ is $\pm$ 0.02 (without red) to $\pm$ 0.02 (with red) times lower than in the field."
+ This is the first attempt to establish whether the relation between SF activity and galaxy mass depends on environment., This is the first attempt to establish whether the relation between SF activity and galaxy mass depends on environment.
+ We find that this relation in. clusters ts significantly different from the field. at all redshifts when REL galaxies are included. and at least at z>0.6 for 24;m+BEL galaxies.," We find that this relation in clusters is significantly different from the field, at all redshifts when REL galaxies are included, and at least at $z>0.6$ for $\micron$ +BEL galaxies."
+ Discriminating star-forming from AGN red galaxies will allow us to quantify with higher precision the environmental effects., Discriminating star-forming from AGN red galaxies will allow us to quantify with higher precision the environmental effects.
+ The observed differences between the SFRs in clusters and in the field can be considered a lower limit to the real gap., The observed differences between the SFRs in clusters and in the field can be considered a lower limit to the real gap.
+ In fact. we are surely considering as cluster members also star-forming galaxies that are either in projection or just recently accreted by the cluster and have not been affected yet by its influence.," In fact, we are surely considering as cluster members also star-forming galaxies that are either in projection or just recently accreted by the cluster and have not been affected yet by its influence."
+ This result cannot arise from having severely underestimated the dustcorrection to the SF© Rio., This result cannot arise from having severely underestimated the dustcorrection to the $SFR_{[OII]}$ .
+ If we were to adjust the [OIL] extinction to bring the field and clusters into agreement. many of the [ΟΠ detected sources should be detected at 24/m. and they are not.," If we were to adjust the [OII] extinction to bring the field and clusters into agreement, many of the [OII] detected sources should be detected at $\micron$ , and they are not."
+"(e.g.Blitz1993). formation{(ο,al.ο,.Bowud&Young1999:WongBlitzLerovet2008:Dieielet.al.2008).. (e...Ikeu-al.2007:Bieielet2008).. (Daddietal.2010:Cenzel2010).. collapse(e.g...[οίνος&Ostri","\citep[e.g.,][]{BLITZ93}. \citep[e.g.,][]{ROWND99,WONG02,LEROY08,BIGIEL08}, \citep[e.g.,][]{KENNICUTT07, BIGIEL08}. \citep{DADDI10,GENZEL10}. \citep[e.g.,][]{MCKEE07}."
+ker and molecular gas is an inportaut iuput aud benchinar for models attempting to reproduce todavs galaxies or ealaxy populations., and molecular gas is an important input and benchmark for models attempting to reproduce today's galaxies or galaxy populations.
+ The importance of this topic has led to several stucies of the relationship between surface deusities of IT» aud the star formation rate., The importance of this topic has led to several studies of the relationship between surface densities of $_2$ and the star formation rate.
+ Mauv of these focus on single ealaxies (e.g...al.Heveret2001:Ieunieutt2007:2010:Ralunanetal.2010) or a small sample (e.g.Wilsonetal.2009:Warren 2010).," Many of these focus on single galaxies \citep[e.g.,][]{HEYER04,KENNICUTT07,SCHUSTER07,BLANC09,VERLEY10,RAHMAN10} or a small sample \citep[e.g.,][]{WONG02,LEROY08,BIGIEL08,WILSON09,WARREN10}."
+.. Restricted by je availability of sensitive aud wide-field molecular gas -laps. studies of large sets of ealaxies (e.g...Youngetal.al.2002:Leroyet2005). mostlv used iuteeratec ueasurenieuts or a few poiutiugs por galaxy.," Restricted by the availability of sensitive and wide-field molecular gas maps, studies of large sets of galaxies \citep[e.g.,][]{YOUNG96,KENNICUTT98A,ROWND99,MURGIA02,LEROY05} mostly used integrated measurements or a few pointings per galaxy."
+ To date. no 1oniogeneous analvsis of the correlation between the star ormatiou rate and Us surface densities iu a large set of rearby galaxies at good spatial resolutiou exists.," To date, no homogeneous analysis of the correlation between the star formation rate and $_2$ surface densities in a large set of nearby galaxies at good spatial resolution exists."
+ Iu this letter we take this next logical step. comparing uolecular gas traced by CO. emissiou to recent star formation traced by ultraviolet auc infraree enission at lL kpe resolution across a large sample of 30 nearby galaxies.," In this letter we take this next logical step, comparing molecular gas — traced by CO emission — to recent star formation — traced by ultraviolet and infrared emission — at $1$ kpc resolution across a large sample of 30 nearby galaxies."
+ This sjuupleis significantly larger and ignore diverse than that of Bigicletal.(2008. DOS).., This sampleis significantly larger and more diverse than that of \citet[][hereafter B08]{BIGIEL08}.
+" From 2007-2010. the HERA CO-Line Extragalactic Survey (IERACLES.firstmapsarepre-seutedinLeroyetal;2009) collaboration used the IRAM 30-1 to coustruct maps of CO J=2>1 cnussion frou, [8 nearby galaxies."," From 2007-2010, the HERA CO-Line Extragalactic Survey \citep[HERACLES, first maps are presented
+ in][]{LEROY09} collaboration used the IRAM 30-m to construct maps of CO $J=2\rightarrow1$ emission from 48 nearby galaxies."
+ Because the targets overlap surveys by Spitzer GuostlySINGS.Ienuicutt and GALEX (uostlvtheNCS.CaldePaz 2007).. excellent uniltiwaveleneth data are available," Because the targets overlap surveys by \citep[mostly
+ SINGS,][]{KENNICUTT03} and GALEX \citep[mostly the
+ NGS,][]{GILDEPAZ07}, , excellent multiwavelength data are available"
+NNNOTADPAGG. NASA Origins of Solar Systems grant NNCGOSGILOG. and is contributed in part to NASA Astrobiology Institute grant NCC2-1056.,"NNX07AP46G, NASA Origins of Solar Systems grant NNG05GI10G, and is contributed in part to NASA Astrobiology Institute grant NCC2-1056."
+Pulsations have been detected [rom stars on a wide range of timescales. from several tens of seconds in the case of white ciwarfs to several vears in the case of red-elant stars.,"Pulsations have been detected from stars on a wide range of timescales, from several tens of seconds in the case of white dwarfs to several years in the case of red-giant stars."
+ These pulsations manifest themselves through a periodic variation of the stellar brightness., These pulsations manifest themselves through a periodic variation of the stellar brightness.
+ Pulsating stars can also be found over a wide range of parameter space (temperature. luminosity) in the Ht diagram (eg Jelfery 2008).," Pulsating stars can also be found over a wide range of parameter space (temperature, luminosity) in the HR diagram (eg Jeffery 2008)."
+ X detailed study of the photometric variability. of individual systems can give insight to the physical conditions deep inside the star (eg Wurtz 2004)., A detailed study of the photometric variability of individual systems can give insight to the physical conditions deep inside the star (eg Kurtz 2004).
+ In recent vears many photometric surveys have been undertaken. leading to a corresponding increase in the number of known stellar pulsators.," In recent years many photometric surveys have been undertaken, leading to a corresponding increase in the number of known stellar pulsators."
+ Factors such as cadence. depth. sky coverage and duration makes any individual survey more (or less) likely to discover specific tvpes of stellar pulsator.," Factors such as cadence, depth, sky coverage and duration makes any individual survey more (or less) likely to discover specific types of stellar pulsator."
+ The Rapid Temporal Survey (RATS) is a deep. high-cadence photometric survey covering nearly. 40 square degrees which took place between 2003 and 2010 (Ramsay Llakala 2005. Darclay ct al 2011).," The Rapid Temporal Survey (RATS) is a deep, high-cadence photometric survey covering nearly 40 square degrees which took place between 2003 and 2010 (Ramsay Hakala 2005, Barclay et al 2011)."
+ This strategy allows us to detect sources which vary in their intensity on a timescale of a few minutes to several hours., This strategy allows us to detect sources which vary in their intensity on a timescale of a few minutes to several hours.
+ 1n our first set of wide-field camera data taken in 2003. we identified a small number of blue stars which pulsate on a period between 4070 mins.," In our first set of wide-field camera data taken in 2003, we identified a small number of blue stars which pulsate on a period between 40–70 mins."
+ An analysis of their optical spectra indicated. they were SN Phe stars or 0 Set stars (Ramsay ct al 2006)., An analysis of their optical spectra indicated they were SX Phe stars or $\delta$ Sct stars (Ramsay et al 2006).
+ SX Phe stars are old. metal-poor stars which are likely to be halo objects (see Nemec Mateo 1990 for a review).," SX Phe stars are old, metal-poor stars which are likely to be halo objects (see Nemec Mateo 1990 for a review)."
+ The à Set stars show similar characteristics to the SX Phe stars but have solar metallicities anc more likely to be located in the thin disk (see Dreger 2000 for a review)., The $\delta$ Sct stars show similar characteristics to the SX Phe stars but have solar metallicities and more likely to be located in the thin disk (see Breger 2000 for a review).
+ Although ὁ Set-like pulsations have been detected, Although $\delta$ Sct-like pulsations have been detected
+"Integrating equations (1)). (2)) aud (3)). we have where 2, is the radial width of the WS where the shell emerges. AL,","Integrating equations \ref{eq:jetM}) ), \ref{eq:jetz}) ) and \ref{eq:jetR}) ), we have where $R_o$ is the radial width of the WS where the shell emerges. $\dot{M}_o$,"
+". D,OF and Pop are 1he initial mass. lougittcinal momentum and transverse ixneutum flows input from the WS iuto the shell."," $\dot{P}_{oz}$ and $\dot{P}_{oR}$ are the initial mass, longitudinal momentum and transverse momentum flows input from the WS into the shell."
+ If à negligible portion of the shocked-eas mass aud momentum are ost from the inner shell into the cocoon. then the total shell mass flow and momentum flow at a distauce /? are M. P. axd Pr.," If a negligible portion of the shocked-gas mass and momentum are lost from the inner shell into the cocoon, then the total shell mass flow and momentum flow at a distance $R$ are $\dot{M}$, $\dot{P}_{z}$ and $\dot{P}_{R}$."
+ In the absence of shell mass losses. the velocity field witiu he shell must be localy j»arallel to the shell surface (the bow shock). although tiere nay be gradieuts in the he velocity across the shell thickuess perpeucdicular to tle shell surface.," In the absence of shell mass losses, the velocity field within the shell must be locally parallel to the shell surface (the bow shock), although there may be gradients in the of the velocity across the shell thickness perpendicular to the shell surface."
+ Uicer this assumption. t atio eftp is the same throughout the shell thickness at any 2. and equal o the slope dz/dIt of t shell’s surface.," Under this assumption, the ratio $v_z/v_R$ is the same throughout the shell thickness at any $R$, and equal to the slope $dz/dR$ of the shell's surface."
+ Since οςορ is coustaut across the shell thickness. the ratio of the total inomentuinlows P.=fe(Ros)p2nRds and Pr=fep(R.s)p2sRds(vp.P. also gives the local slo ol the shell surface (here the integral over s denotes suuuimation across the shell thickness).," Since $v_z/v_R$ is constant across the shell thickness, the ratio of the total momentumflows $\dot P_z =\int v_z (R,s)\, \rho\ 2\pi R ds$ and $\dot P_R =\int v_R (R,s)\, \rho\ 2\pi R ds= (v_R/v_z)\dot P_z $ also gives the local slope of the shell surface (here the integral over $s$ denotes summation across the shell thickness)."
+" T shape of the shell can thus be derived by integrating The locus of the bow shock/outflow shell is therefore given by with - =Oat R=HR,", The shape of the shell can thus be derived by integrating The locus of the bow shock/outflow shell is therefore given by with $z=0$ at $R=R_o$.
+ With expression (8)) eeivine[eJ the locus of the outflow shell. we may now turn to the distribution of velocities within the shell.," With expression \ref{eq:Ijets}) ) giving the locus of the outflow shell, we may now turn to the distribution of velocities within the shell."
+ Consider first the limit in which the material added to the shell at aly poiut mixes instantaneously with all the material already. flowing alone in the shell at that yosition., Consider first the limit in which the material added to the shell at any point mixes instantaneously with all the material already flowing along in the shell at that position.
+" Hfthe uew"" and ""existing""© momentum were to mix thoroughly.ott, then there would be no velocity gradients across the thiekuess of the shell."," If the “new” and “existing” momentum were to mix thoroughly, then there would be no velocity gradients across the thickness of the shell."
+ The velocity of shell material in the shock frame would be equal to the value at auy point. given by," The velocity of shell material in the shock frame would be equal to the value at any point, given by"
+ απο 2853 was discovered cling the first observations of the Calactic Centre region performed by theGRANAT satellite in Spring 1990., $-$ GRS $-$ 2853 was discovered during the first observations of the Galactic Centre region performed by the satellite in Spring 1990.
+ The source was detected dy the low-enerev (Ll30 keV) imaging telescope ART-P in the March 21April δ observatious and was teutatively associated (\landrou 19090)) to the softEINSTEIN source LE 2850 (Watsonctal. 1981))., The source was detected by the low-energy (4–30 keV) imaging telescope ART-P in the March 24--April 8 observations and was tentatively associated \cite{Mand90}) ) to the soft source 1E $-$ 2850 \cite{Wats81}) ).
+" Further analysis of the same data (SuuvaevL990.Suuvaevetal.1991.Pavlinskyctal. 1991)) refined the source position. obtaining a=his, à=28752/5|"" (D1950. error radius 15"". confidence)."," Further analysis of the same data \cite{Suny90,Suny91,Pavl94}) ) refined the source position, obtaining $\alpha=17^{\rm h}41^{\rm m}50^{\rm s}$, $\delta=-28^{\circ}52^{\prime}54\arcsec$ (B1950, error radius $45\arcsec$, confidence)."
+ The possible assoclatioun with LE 2850 aud also with the rearby transient. GS 289 (Mitsudaetal. 19903). was ruled out., The possible association with 1E $-$ 2850 and also with the nearby transient GS $-$ $-$ 289 \cite{Mits90}) ) was ruled out.
+ The average |20 keV intensity of CRS 2853 Was 0.6d7 uCrab. correspondiug to (1.6+0.1)«10%eres3 af SJ5 kpe. aud the source spectrum could be fitted by a thermal brenisstrahluug with a temperature of ~8 seV. During the saneGRANAT observations. CRS 2853 was detected above the 3.50 detection level of the soft Camuna-ray telescope SICALA in the 10.100 keV baud (Cliurazovetal.1993.Vargasct 1997)).," The average 4--20 keV intensity of GRS $-$ 2853 was $9.6\pm0.7$ mCrab, corresponding to $(1.6\pm0.1)\times10^{36}{\rm erg~s}^{-1}$ at 8.5 kpc, and the source spectrum could be fitted by a thermal bremsstrahlung with a temperature of $\sim8$ keV. During the same observations, GRS $-$ 2853 was detected above the $3.5\sigma$ detection level of the soft Gamma-ray telescope SIGMA in the 40–100 keV band \cite{Chur93,Varg97}) )."
+ Ou the other ματς. both ART-P aud SICAIA did not detect the source ~Lo months later during the Fall 1990 observation Canrpaieu. sugeesting GRS 2853 is transient in nature.," On the other hand, both ART-P and SIGMA did not detect the source $\sim 4$ months later during the Fall 1990 observation campaign, suggesting GRS $-$ 2853 is transient in nature."
+ A 236 upper limit of 1.2 miCrab was obtained by ART-P in the |20 keV baud (Pavlinsiyetal. 1991)). thus nuplviug a drop iu the source intensity of at least a factor of 7.," A $3\sigma$ upper limit of 1.2 mCrab was obtained by ART-P in the 4–20 keV band \cite{Pavl94}) ), thus implying a drop in the source intensity of at least a factor of 7."
+ Moreover.GRANAT failed to detect GRS 2853 in all the subsequent caupaigus on the Galactic Centre (Spring 1991. Fall 1991. Spring 1992. e.c. Churazovetal.1993. 1997)).," Moreover, failed to detect GRS $-$ 2853 in all the subsequent campaigns on the Galactic Centre (Spring 1991, Fall 1991, Spring 1992, e.g. \cite{Chur93,Pavl94,Varg97}) )."
+" The ποιος was not observed iu detailed mappings of the Calactic Centre region by soft N-rav instruments likeTEIN (0.5[.5 keV. 30 upper lait of ~0.7 mCrab. SCC Watsonetal. 1981)). (330. keV. (Ls uC'rab. οκetal. 1987)). and byROSAT τοις2.5 keV. 0,1 niCrab. Predehl&Trüuirper199 1)). lus confriuüug its transient nature."," The source was not observed in detailed mappings of the Galactic Centre region by soft X-ray instruments like (0.5–4.5 keV, $3\sigma$ upper limit of $\sim0.7$ mCrab, see \cite{Wats81}) ), (3–30 keV, $<0.8$ mCrab, \cite{Skin87}) ), and by (0.8--2.5 keV, $<0.1$ mCrab, \cite{Pred94}) ), thus confirming its transient nature."
+ More recently. 10 detections of CRS 2853 were reported by RXTE-ASM in the 210 keV cnerev baud since February 19906.," More recently, no detections of GRS $-$ 2853 were reported by -ASM in the 2–10 keV energy band since February 1996."
+ Iu the next section we briefly introduce the Wide Field Caincras telescopes aud report on the observations of CRS 2853., In the next section we briefly introduce the Wide Field Cameras telescopes and report on the observations of GRS $-$ 2853.
+ The resolved spectroscopy of the burst data is preseuted in Section 3.Pa while the iapact of our results on the kuowledge of the source are discussed. iu Section 1.," Time resolved spectroscopy of the burst data is presented in Section 3, while the impact of our results on the knowledge of the source are discussed in Section 4."
+ Iu particular. we propose GRS 2853 as a transient low-mass X-ray binary harbourius a neutron star and we eive au estimate of the source distance.," In particular, we propose GRS $-$ 2853 as a transient low-mass X-ray binary harbouring a neutron star and we give an estimate of the source distance."
+ Oue of the main scieutific objectives of the Wide Field Cameras (WEC) on board theX satellite is the study of the timine/spectral behavior of both transient and persistent sources of the Galactic Bulee region. X- binaries in particular. ou time scales from seconds to vears.," One of the main scientific objectives of the Wide Field Cameras (WFC) on board the satellite is the study of the timing/spectral behavior of both transient and persistent sources of the Galactic Bulge region, X-ray binaries in particular, on time scales from seconds to years."
+ To this cud. an observation program of svstcatic monitoring of the Ser A sky region is being carriccd out," To this end, an observation program of systematic monitoring of the Sgr A sky region is being carried out"
+"enough to account for the IR luminosity (which has values between10°4 erg s! and 1036 erg s!, for the sources XTE J1859+226 (b) and GRS1915+105 (a), respectively).","enough to account for the IR luminosity (which has values between$10^{34}$ erg $s^{-1}$ and $10^{36}$ erg $s^{-1}$, for the sources XTE J1859+226 (b) and GRS1915+105 (a), respectively)."
+" It is also possible that additional emission results from electrons that were re-accelerated in optically thin regions away from the reconnection zone (behind shock waves along the jet excited by the plasmon clouds, as emphasized above)."," It is also possible that additional emission results from electrons that were re-accelerated in optically thin regions away from the reconnection zone (behind shock waves along the jet excited by the plasmon clouds, as emphasized above)."
+" This idea is supported for instance by IR observations of GRS 1915-105, which indicate a separation between the IR emission and thecentral source ~0.2 arcsec or ~2500 AU at 12.5 kpc (Sams et al."," This idea is supported for instance by IR observations of GRS 1915+105, which indicate a separation between the IR emission and thecentral source $\sim 0.2$ arcsec or $\sim
+2500$ AU at 12.5 kpc (Sams et al."
+" 1996), ie.,"," 1996), ie.,"
+ further away from the predicted launching region by the present model (Fig., further away from the predicted launching region by the present model (Fig.
+ 2)., 2).
+" Finally, because the magnetic power released by reconnection is much greater than the observed radio luminosity, only a small fraction of this energy is necessary to accelerate particles."," Finally, because the magnetic power released by reconnection is much greater than the observed radio luminosity, only a small fraction of this energy is necessary to accelerate particles."
+ Then most of the energy released would be used to heat the coronal and also the disk gas surface (by thermal conduction along the field lines)., Then most of the energy released would be used to heat the coronal and also the disk gas surface (by thermal conduction along the field lines).
+" This supplementary heating could perturb the accretion disk and result in the variability of the soft x-ray emission observed during the flare phase of GRS 1915+105 (Dahwan et al.,"," This supplementary heating could perturb the accretion disk and result in the variability of the soft x-ray emission observed during the flare phase of GRS 1915+105 (Dahwan et al.,"
+ 2000)., 2000).
+" According to the scenario presented here, Fig."," According to the scenario presented here, Fig."
+ 1 would correspond to the configuration of the system immediately before a radio flare., 1 would correspond to the configuration of the system immediately before a radio flare.
+" Thus, this could be the system configuration at the end of the ""hard' SPLS and the transition from the 'hard' to the 'soft SPLS would be due to violent magnetic reconnection events between the magnetic field lines of the inner disk region and those that are anchored into the black hole."," Thus, this could be the system configuration at the end of the 'hard' SPLS and the transition from the 'hard' to the 'soft' SPLS would be due to violent magnetic reconnection events between the magnetic field lines of the inner disk region and those that are anchored into the black hole."
+ It is important to emphasize that the physical origin of the hard x-ray emission in the SPLS is yet controversial in the literature., It is important to emphasize that the physical origin of the hard x-ray emission in the SPLS is yet controversial in the literature.
+" Most models claim for the inverse Compton scattering as the dominant radiative mechanism (e.g. Zdziarski, 2000), and photons with MeV energies suggest that the scattering occurs in a non-thermal corona (e.g. Zdziarski, 2001)."," Most models claim for the inverse Compton scattering as the dominant radiative mechanism (e.g. Zdziarski, 2000), and photons with MeV energies suggest that the scattering occurs in a non-thermal corona (e.g. Zdziarski, 2001)."
+" For Poutanen Fabian (1999), the origin of scattering electrons could be the coronal active magnetic regions that rise from the disk, which agrees with the scenario proposed here."," For Poutanen Fabian (1999), the origin of scattering electrons could be the coronal active magnetic regions that rise from the disk, which agrees with the scenario proposed here."
+" Furthermore, the SPLS tends to dominate the spectrum of the BHXRBs as the luminosity reaches values near the Eddington rate, which also agrees with the model presented here."," Furthermore, the SPLS tends to dominate the spectrum of the BHXRBs as the luminosity reaches values near the Eddington rate, which also agrees with the model presented here."
+" Despite the huge difference in scales, AGNs/quasars and microquasars have similar morphologies (Mirabel Rodrigues 1998)."," Despite the huge difference in scales, AGNs/quasars and microquasars have similar morphologies (Mirabel Rodrigues 1998)."
+" Indeed, some studies indicate that the similarity between these systems is more than morphological: they are possibly subject to the same physical processes (e.g., de Gouveia Dal Pino 2005 and references therein for a review)."," Indeed, some studies indicate that the similarity between these systems is more than morphological: they are possibly subject to the same physical processes (e.g., de Gouveia Dal Pino 2005 and references therein for a review)."
+" If this is the case, then a generalization of the above scenario to the extragalactic sources is almost straightforward."," If this is the case, then a generalization of the above scenario to the extragalactic sources is almost straightforward."
+" Indeed, relativistic jets from AGNs are also observed to produce relativistic episodic ejections that originate synchrotron law spectrum with similar spectral indices as above."," Indeed, relativistic jets from AGNs are also observed to produce relativistic episodic ejections that originate synchrotron power-law spectrum with similar spectral indices as above."
+" Considering the same assumptions presented above, we obtain similar scaling relations for AGNs."," Considering the same assumptions presented above, we obtain similar scaling relations for AGNs."
+ Figure 3 depicts a synthesis of the generalization of the magnetic reconnection scenario for relativistic sources including both microquasars and AGNs., Figure 3 depicts a synthesis of the generalization of the magnetic reconnection scenario for relativistic sources including both microquasars and AGNs.
+ The diagram shows the calculated magnetic power released in violent magnetic reconnection events as a function of the central source mass., The diagram shows the calculated magnetic power released in violent magnetic reconnection events as a function of the central source mass.
+" The symbols correspond to the observed radio luminosities of superluminal components (stars for microquasars, circles and triangles for the low luminous AGNSs, i.e., LINERs (see below) and Seyfert galaxies, respectively, and squares for luminous AGNs)."," The symbols correspond to the observed radio luminosities of superluminal components (stars for microquasars, circles and triangles for the low luminous AGNs, i.e., LINERs (see below) and Seyfert galaxies, respectively, and squares for luminous AGNs)."
+ The same parameter space described in Sect., The same parameter space described in Sect.
+" 3.5 was used for the AGNs, except that the black hole mass now spans the range 5Mo to 10!°Mo."," 3.5 was used for the AGNs, except that the black hole mass now spans the range $5 M_{\odot}$ to $10^{10} M_{\odot}$."
+" Figure 3 indicates that the magnetic power released during violent reconnection events is able to explain the emission of relativistic radio blobs from both microquasars and the so-called low-luminous AGNs (LLAGNs), which include Seyfert galaxies and LINERs (these latter have emission-line luminosities that are typically a factor of 10? orders of magnitude smaller than those of luminous AGNs; e.g., Ho et al."," Figure 3 indicates that the magnetic power released during violent reconnection events is able to explain the emission of relativistic radio blobs from both microquasars and the so-called low-luminous AGNs (LLAGNs), which include Seyfert galaxies and LINERs (these latter have emission-line luminosities that are typically a factor of $10^2$ orders of magnitude smaller than those of luminous AGNs; e.g., Ho et al."
+ 1997)., 1997).
+ This establishes a correlation between the magnetic reconnection power of stellar mass and supermassive black holes according to Eq. (, This establishes a correlation between the magnetic reconnection power of stellar mass and supermassive black holes according to Eq. (
+12) spanning over 10° orders of magnitude in mass of the sources.,12) spanning over $10^9$ orders of magnitude in mass of the sources.
+ We notice that a simple power-law dependence of Wz with the black hole mass M may be derivedfrom Eq. (, We notice that a simple power-law dependence of $\dot{W_B}$ with the black hole mass $M$ may be derivedfrom Eq. (
+"12) if we consider that the height of the reconnection zone in the corona, l, is parameterized by Rx, which in turn is Ry= 6GM/c*.","12) if we consider that the height of the reconnection zone in the corona, $l$ , is parameterized by $R_X$ , which in turn is $R_X = 3 R_{S} = 6 GM/c^2$ ."
+ This implies that according to the present, This implies that according to the present
+"by LESS. which are respectively where /7, is the detector threshokl energv and D is the spectral index if the 5-rav enission is a power law.","by HESS, which are respectively where $E_{th}$ is the detector threshold energy and $\Gamma$ is the spectral index if the $\gamma$ -ray emission is a power law."
+ By requiring that the following integral over LESS field of view ancl over (he luminosity range given in Eq.(13)) gives the number of sources LESS detects above 6 percent of the Crab flux we will obtain the normalisation factor € In Eq.(14)) pius 18 the density of SNRs and PWNe., By requiring that the following integral over HESS field of view and over the luminosity range given in \ref{lumfun}) ) gives the number of sources HESS detects above 6 percent of the Crab flux we will obtain the normalisation factor $c$ In \ref{integral}) ) $\rho_{sources}$ is the density of SNRs and PWNe.
+" The integral flux above £i [rom resolved and unresolved sources [rom a region in the Galaxy which extends from D,,5, to D,,,, in heliocentric distance. [rom 5,5, to 5,,, in latitude and [rom /,,5, to /,,,, in longitude is"," The integral flux above $E_0$ from resolved and unresolved sources from a region in the Galaxy which extends from $D_{min}$ to $D_{max}$ in heliocentric distance, from $b_{min}$ to $b_{max}$ in latitude and from $l_{min}$ to $l_{max}$ in longitude is"
+Figure 5 shows our best fits to the vertical density profiles for the MW. simulation for radial bins Ro= kp. R=1.6 kp. R=7.YD pe and R=912 kpc.,"Figure \ref{f:density_profiles} shows our best fits to the vertical density profiles for the MW simulation for radial bins $R = 2 - 4$ $\kpc$, $R = 4 - 6$ $\kpc$, $R = 7 - 9$ $\kpc$ and $R = 9 - 12$ $\kpc$."
+ As in JOS. the iunerinost radial biu is best fit bv a double exponential fuuction: however. all other bins are better fit bv the stim of two double sech? profiles. which is in agreement with the theoretical work by Spitzer(1912) and observational results for other galaxies by Yoachimi&Dalcauton(2006).," As in J08, the innermost radial bin is best fit by a double exponential function; however, all other bins are better fit by the sum of two double $^2$ profiles, which is in agreement with the theoretical work by \citet{Spitzer1942} and observational results for other galaxies by \citet{Yoachim2006}."
+". Despite the small discrepancy between secl? expouential fits. the simulation is iu good qualitative agreement with the SDSS-based analysis of the Milkv Wav (05),"," Despite the small discrepancy between $^2$ exponential fits, the simulation is in good qualitative agreement with the SDSS-based analysis of the Milky Way (J08)."
+ Iu addition. Figure 5 shows the vertical density profiles as divided iuto two populations: radial migrators 2kpe) andanin situ population (JRΠεν)<2 kpc).," In addition, Figure \ref{f:density_profiles} shows the vertical density profiles as divided into two populations: radial migrators $|R - R_{form}| > 2 \kpc$ ) and an in situ population $|R - R_{form}| \leq 2 \kpc$ )."
+ We note the second component fit for radial bius R=7 O9kpeand R=912 kpc is entirely dominated o» star particles that have uniegrated., We note the second component fit for radial bins $R = 7 - 9$ $\kpc$ and $R = 9 - 12$ $\kpc$ is entirely dominated by star particles that have migrated.
+ We expect that stars that moved from elsewhere combined with stars hat were born locally should not uaturally couspire to xoduce a single continuous profile., We expect that stars that moved from elsewhere combined with stars that were born locally should not naturally conspire to produce a single continuous profile.
+ For the solar cvlinder. R=79kpc. we find that the nodel distribution of stars as a function of |:| resembles a double sech profile. with the “transition” height of 20.275kpe. comparable to 1kpe found by JOS for the Galactic disk.," For the solar cylinder, $R = 7 - 9 \kpc$, we find that the model distribution of stars as a function of $|z|$ resembles a double $^2$ profile, with the “transition” height of $|z|\sim$$0.75
+\kpc$, comparable to $\sim$$1 \kpc$ found by J08 for the Galactic disk."
+ We have fouud the best-fit scale heights ο be 381 pe and 913 pc. in qualitative agreement with he best-fit scale heights of 270 pc aud 1200 pec for the SDSS data (JOS).," We have found the best-fit scale heights to be 381 pc and 913 pc, in qualitative agreement with the best-fit scale heights of 270 pc and 1200 pc for the SDSS data (J08)."
+ The scale height ratio suggested by the siauulationissliehtlv low ~2.1. instead of ~3 from the data.," The scale height ratio suggested by the simulation is slightly low — $\sim$$2.4$, instead of $\sim$$3$ from the data."
+ Moreover. the μπαμπά—ος thick disk to thin dis- normalization is sliehtle discrepant —0.1L. rather than 0.12 sugeested by the SDSS data.," Moreover, the simulation's thick disk to thin disk normalization is slightly discrepant — $\sim$$0.14$, rather than $\sim$$0.12$ suggested by the SDSS data."
+ Overall.— we find that the thick disk iu the MW sinulation. formed through the process of radial migration. is qualitatively very similar to the observed Calactic thick disk.," Overall, we find that the thick disk in the MW simulation, formed through the process of radial migration, is qualitatively very similar to the observed Galactic thick disk."
+ As Figure L clearly attests. this conclusiou does not apply to the LSB simulation.," As Figure \ref{f:density_comp} clearly attests, this conclusion does not apply to the LSB simulation."
+ The mecian metallicity of the Milkv. Way cüsk exhibits a clear vertical eradicut (sec Figure 9. bottom two panels. 108).," The median metallicity of the Milky Way disk exhibits a clear vertical gradient (see Figure 9, bottom two panels, I08)."
+ Notably. the spatial variation of the median ietallicitv does not follow the distribution of the stellar uunber density (105).," Notably, the spatial variation of the median metallicity does not follow the distribution of the stellar number density (I08)."
+ The middle paucl of Figure 6 shows that the MW sinmlatioun reproduces a qualitatively simular metallicity distribution: we note that a coustaut additive offset of 0.2 dex has been applied to [Fo/1I] throughout the simulation so that the mecian metallicity in the plane of the disk in the solar cxliuder matches observatious., The middle panel of Figure \ref{f:color_maps} shows that the MW simulation reproduces a qualitatively similar metallicity distribution; we note that a constant additive offset of 0.2 dex has been applied to [Fe/H] throughout the simulation so that the median metallicity in the plane of the disk in the solar cylinder matches observations.
+ As expected. at low galactic latitudes aud sanall radii. the volume is dominated by high (near solar) iuetallicities.," As expected, at low galactic latitudes and small radii, the volume is dominated by high (near solar) metallicities."
+ At hieher latitudes the volume is increasingly mctal poor., At higher latitudes the volume is increasingly metal poor.
+" As with inetallicitv. previous studies have found a eracieut in the median V;, with respect to : in the Milky Wav (sce Figures 5. 8. and 9 in Boud (2010. hereafter DIOY.. and references therein)."," As with metallicity, previous studies have found a gradient in the median $V_{\phi}$ with respect to $z$ in the Milky Way (see Figures 5, 8, and 9 in Bond (2010, hereafter \nocite{Bond2010}, and references therein)."
+" These authors couclided that V;,, is also well characterized bv a nou-CGaussian distribution (seeBinney2010.fortheiuplieddistributionfunction)...", These authors concluded that $V_{\phi}$ is also well characterized by a non-Gaussian distribution \citep[see][for the implied distribution function]{Binney2010}.
+" The bottom paucl of Figure 6 shows that the AAW simulation reproduces qualitatively simular V;, properties. incliding a strong eracienut."," The bottom panel of Figure \ref{f:color_maps} shows that the MW simulation reproduces qualitatively similar $V_{\phi}$ properties, including a strong $|z|$ gradient."
+" ""Thus ou a gross scale the MN sinimlation qualitatively matches patterus observed m mass density. metallicitv aud rotational velocity in the Milly Way disk."," Thus on a gross scale the MW simulation qualitatively matches patterns observed in mass density, metallicity and rotational velocity in the Milky Way disk."
+ We now look at the solar evliuder in greater detail., We now look at the solar cylinder in greater detail.
+ Tere we study in detail the distributions of age. stellar lnass. rotational velocity anc ietallicitv as fictions : for the solar cxliuder.," Here we study in detail the distributions of age, stellar mass, rotational velocity and metallicity as functions of $|z|$ for the solar cylinder."
+ To be consistent with theanalvsis of high ealactic latitude SDSS data by JOS aud Tos. we select model particles from an annulus with 7 kpe «OR <9 Ipc.," To be consistent with theanalysis of high galactic latitude SDSS data by J08 and I08, we select model particles from an annulus with 7 $\kpc$ $\leq R \leq$ 9 $\kpc$."
+ This radial cut spaus 2.22.8 scale leneths roni the center of modol galaxy (scale leugth 3.2. kpc). and covers the Suus location of ~3 disk scale leneths roni the Milkv Way ceuter (scale leugth ~2.6Ipc) (see uas-corrected value. Table LO. JS).," This radial cut spans 2.2–2.8 scale lengths from the center of model galaxy (scale length $3.2$ $\kpc$ ), and covers the Sun's location of $\sim$ 3 disk scale lengths from the Milky Way center (scale length $\sim$$2.6 \kpc$ ) (see bias-corrected value, Table 10, J08)."
+ The behavior of the MAW simulation iu this particular volue is illustrated in Figure 7.., The behavior of the MW simulation in this particular volume is illustrated in Figure \ref{f:ringberg}. .
+ The top panel shows 10 age distribution as a function of |:|., The top panel shows the age distribution as a function of $|z|$.
+ As expected. rear the midplane. the population is dominated by vouug stars. and the mean age monotonically iucreases witli increasing distance from the midplauc.," As expected, near the midplane, the population is dominated by young stars, and the mean age monotonically increases with increasing distance from the midplane."
+ Onlv very old stars are found at large N ~0.5kpc the mean age is already ~5 Gyr., Only very old stars are found at large $|z|$: by $|z| \sim$$0.5 \kpc$ the mean age is already $\sim$ 5 Gyr.
+ Recall that Figure 2) shows that with increasing distance frou the mudplane. the stellar population becomes dominated by older stars that formed closer to the center of the disk.," Recall that Figure \ref{f:rform_cont} shows that with increasing distance from the midplane, the stellar population becomes dominated by older stars that formed closer to the center of the disk."
+ These two figures taken together eive a coherent picture of the net dynamic effect on the system: Racial migration is able to change the extent of a stars vertical oscillation. as a well as its rotational velocity (muiddle panel. Figure 7)).," These two figures taken together give a coherent picture of the net dynamic effect on the system: Radial migration is able to change the extent of a star's vertical oscillation, as a well as its rotational velocity (middle panel, Figure \ref{f:ringberg}) )."
+ Iu the MW simulatio we find a vertical eradicut of rotational velocity ofτσ -0T kins!kpe (compared to a eradiceut of -30 ausdkpe found by I08)., In the MW simulation we find a vertical gradient of rotational velocity of -17 $\kmskpc$ (compared to a gradient of -30 $\kmskpc$ found by I08).
+ Note that stars which sieuificautlv lag the rotation of the disk in the midplane ave traditionally been regarded as nienibers of the thic- cixk., Note that stars which significantly lag the rotation of the disk in the midplane have traditionally been regarded as members of the thick disk.
+ Another observed trend that is plwsically well notivated by the uet outward aud upward movement of stars over time is the decline of metallicity with increased wieght (bottom pancl. Figure 7)).," Another observed trend that is physically well motivated by the net outward and upward movement of stars over time is the decline of metallicity with increased height (bottom panel, Figure \ref{f:ringberg}) )."
+ In the AAW simulation. he metallicity distribution chauges witli with a best- eradieunt of ~0.19 dexIpc again iu qualitative agreement withthe measured value for the Milkv. Wav of ~0.30 dexkpe+ (Ios.," In the MW simulation, the metallicity distribution changes with $|z|$ with a best-fit gradient of $\sim$$0.19$ $\dexkpc$ , again in qualitative agreement withthe measured value for the Milky Way of $\sim$$0.30$ $\dexkpc$ (I08)."
+ To wnderstand how this treud arises. recall that the top panel of Figure 7 shows that the stellar population," To understand how this trend arises, recall that the top panel of Figure \ref{f:ringberg} shows that the stellar population"
+Let the unlensed dilferential counts of the background galaxies. be denoted by AodiatNES where no is the number of galaxies per unit solid angle. ancl ο is the Hux.,"Let the unlensed differential counts of the background galaxies be denoted by $\frac{dn_{0}(S)}{dS}$, where $n_{0}$ is the number of galaxies per unit solid angle, and $S$ is the flux."
+ The lensed counts are given by where fi is the magnification., The lensed counts are given by where $\mu$ is the magnification.
+ The observed. (Iensed) counts. in the presence of incompleteness. can be written as where O05) is the completeness fraction for a given Hux S.," The observed (lensed) counts, in the presence of incompleteness, can be written as where $\phi(S)$ is the completeness fraction for a given flux $S$ ."
+ Twpicallv. 6=1 for large Uuxes. and ὁ<1 for small Huxes.," Typically, $\phi=1$ for large fluxes, and $\phi<1$ for small fluxes."
+ In the following. we will continue to use the subscript ο to denote unlensed counts. and a tilde to denote the observed counts.," In the following, we will continue to use the subscript $_{0}$ to denote unlensed counts, and a tilde to denote the observed counts."
+ In practice. it is convenient to consider the total counts down to a detection limit 5 In our case (see Fig. 5))," In practice, it is convenient to consider the total counts down to a detection limit $S$ In our case (see Fig. \ref{fig-slopes}) )"
+ and for many applications. the intrinsic unlensed counts are well described by a power Law of the form where 23 is the Lux slope parameter. which is related to the magnitude slope parameter a=d(logno)/dim (Equation 2)) by 3=2.56.," and for many applications, the intrinsic unlensed counts are well described by a power law of the form where $\beta$ is the flux slope parameter, which is related to the magnitude slope parameter $\alpha \equiv d(\log n_{0})/dm$ (Equation \ref{eq:alpha}) ) by $\beta=2.5\alpha$."
+ In this case. Equation CX1)) takes the simple form In this case. the observed Iensed counts are given hy As a result. the mean integrated.> lensecl counts CN observed in the field are given by where my=.Asoτα is the observed. surface. number clensity of unlensed galaxies.," In this case, Equation \ref{eq:dnds_lens}) ) takes the simple form In this case, the observed lensed counts are given by As a result, the mean integrated lensed counts $\langle \widetilde{N}
+\rangle$ observed in the field are given by where $\tilde{n}_{0} = \int_{S}^{\infty} dS' \phi \frac{dn_{0}}{dS}$ is the observed surface number density of unlensed galaxies."
+ ‘The probability of finding the 2 galaxies at position 4; is piXD(8;)πω1," The probability of finding the $i^{\rm th}$ galaxies at position ${\mathbf \theta}_{i}$ is $p_{i} \propto \tilde{n}({\mathbf
+\theta}_{i}) = \tilde{n}_{0} [\mu({\mathbf
+\theta}_{i})]^{\beta-1}$."
+ Thus. the likelihood of finding N galaxies in our incomplete survey at positions @; is given by where (ΑΔN35) stands for the Poisson distribution with mean EN}.," Thus, the likelihood of finding $N$ galaxies in our incomplete survey at positions ${\mathbf \theta}_{i}$ is given by where $P(N;\langle N \rangle)$ stands for the Poisson distribution with mean $\langle N \rangle$."
+ The log likelihood function is thus where constant terms have been dropped., The log likelihood function is thus where constant terms have been dropped.
+ ‘These two expressions are identical to the corresponding expressions in SAL. except that incomplete counts CV? and Do replace the complete counts £N? and. my. respectively.," These two expressions are identical to the corresponding expressions in SKE, except that incomplete counts $\langle \widetilde{N}
+\rangle$ and $\tilde{n}_{0}$ replace the complete counts $\langle N
+\rangle$ and $n_{0}$, respectively."
+ Therefore. the cllect of incompleteness is very simple to take into account. (," Therefore, the effect of incompleteness is very simple to take into account. ("
+We will drop the tilde in the rest of the text).,We will drop the tilde in the rest of the text).
+ In fact. there is no reason to restrict the sample to galaxies brighter than the completeness limit.," In fact, there is no reason to restrict the sample to galaxies brighter than the completeness limit."
+ On the contrary. it is desirable to include faint galaxies to improve the signal-to-nolse ratio.," On the contrary, it is desirable to include faint galaxies to improve the signal-to-noise ratio."
+ Note that this simplification results from the factorisation of gà in Equation (AG)) ancl thus only holds if the intrinsic unlensed. counts follow a power Law., Note that this simplification results from the factorisation of $\mu$ in Equation \ref{eq:dnds_obs})) and thus only holds if the intrinsic unlensed counts follow a power law.
+ Note also that the likelihood. £ is independent of the completeness function O09). , Note also that the likelihood ${\mathcal L}$ is independent of the completeness function $\phi(S)$ 
+foregrounds. beam wncertaiutics. calibration or noise estimation crrors) to the final cosmological parameters.,"foregrounds, beam uncertainties, calibration or noise estimation errors) to the final cosmological parameters."
+ This is not straightforward in the hivbrid scheme used by the WMAP code., This is not straightforward in the hybrid scheme used by the WMAP code.
+ Second. this new approach corresponds to the exact low-f pixel space Likelihood part of the WMAP code. nof the approximate high-( MASTER part.," Second, this new approach corresponds to the exact $\ell$ pixel space likelihood part of the WMAP code, not the approximate $\ell$ MASTER part."
+" Still. our method can handle arbitrary high (Cs. Third. once the one-time pre-processing step las becu completed. the computational expense of our estimator is determined by the cost of (4,44 spline evaluations. while a pixel space approach requires a matrix inversion. aud therefore scales as "," Still, our method can handle arbitrary high $\ell$ 's. Third, once the one-time pre-processing step has been completed, the computational expense of our estimator is determined by the cost of $\ell_{\textrm{max}}$ spline evaluations, while a pixel space approach requires a matrix inversion, and therefore scales as $\mathcal{O}(N_{\textrm{pix}}^3)$."
+For the cases considered iu this paper. the CPU time OUNS).required for the CBR WMADP estimator up to (=200 was ~0.2 nülliseconds. while it was e20 seconds for the pixel space approach up to (=32. for amap with 2500 pixels.," For the cases considered in this paper, the CPU time required for the GBR WMAP estimator up to $\ell=200$ was $\sim0.2$ milliseconds, while it was $\sim20$ seconds for the pixel space approach up to $\ell=32$, for a map with $~2500$ pixels."
+ Iu order to validate our estimator. we applied it to a low-resolution simulated data set. and conrpared it to slices through the exact joiut Likelihood as computed by brute-force evaluation in pixel space.," In order to validate our estimator, we applied it to a low-resolution simulated data set, and compared it to slices through the exact joint likelihood as computed by brute-force evaluation in pixel space."
+ The agreement between the two approaches was excellent., The agreement between the two approaches was excellent.
+ We then applied the same estimator to the S-vear WALAP temperature data. and estimated both a new powcr spectrum and new cosmological parameters within a standard six-paramcter ACDM model.," We then applied the same estimator to the 5-year WMAP temperature data, and estimated both a new power spectrum and new cosmological parameters within a standard six-parameter $\Lambda$ CDM model."
+ The results from these calewlatious are interesting., The results from these calculations are interesting.
+ Fist. our power spectrum is statistically τον similar to the official WALAP spectimm. with no visible biases seen and relative fluctuations within the level predicted by noise aud sky cut.," First, our power spectrum is statistically very similar to the official WMAP spectrum, with no visible biases seen and relative fluctuations within the level predicted by noise and sky cut."
+" Nevertheless. we do find siguificaut differences in terius of cosmological parameters. aud most notably in the spectral iudex of scalar perturbations. εἰς, "," Nevertheless, we do find significant differences in terms of cosmological parameters, and most notably in the spectral index of scalar perturbations, $n_{\textrm{s}}$ ."
+"Specifically, we fud ως=0.973001L which is oulv 10σ away from uuitv aud 0.66 higher than the official WALAP result. ος=0.965+£0.011,"," Specifically, we find $n_{\textrm{s}} = 0.973\pm
+0.014$, which is only $1.9\sigma$ away from unity and $0.6\sigma$ higher than the official WMAP result, $n_{\textrm{s}} = 0.965\pm
+0.014$."
+" This result resembles very much the outcome of a re- we did with the 3-vear WALAP temperature data (Eviksenctal.2007a).. for which we found abias of 0.le inv, compared to the official WAIAP results."," This result resembles very much the outcome of a re-analysis we did with the 3-year WMAP temperature data \citep{eriksen:2007a}, for which we found abias of $0.4\sigma$ in $n_{\textrm{s}}$ compared to the official WMAP results."
+ This bias was due to the sub-optimal \MIASTER-based likelihood approximation (Ilivonetal.2002:Verde2003) used by the WMAP team betweerf =12 and 30. whereas we used an exact estimator nm he same range.," This bias was due to the sub-optimal MASTER-based likelihood approximation \citep{hivon:2002,verde:2003}
+ used by the WMAP team between $\ell=12$ and 30, whereas we used an exact estimator in the same range."
+ This study later prompted the WALAP o change their codes to use an exact Likehhood evaluator up to (—30., This study later prompted the WMAP to change their codes to use an exact likelihood evaluator up to $\ell=30$.
+ Iu the same study. we aso tried to increase the (- range for our exac estimator to (—50. but found zuall differences.," In the same study, we also tried to increase the $\ell$ -range for our exact estimator to $\ell=50$, but found small differences."
+" We therefore coicluded. perhaps somewha prematurely, that an exact csstimator up to (=30 was sufficient for obtaining accurate results."," We therefore concluded, perhaps somewhat prematurely, that an exact estimator up to $\ell=30$ was sufficient for obtaining accurate results."
+ On the coutrary. in this paper we find still fiid significant chanecs when increasing the exact estimatex up to {—200.," On the contrary, in this paper we find still find significant changes when increasing the exact estimator up to $\ell=200$."
+ Iu retrospect. this shoukL perhaps not come as a conplete surprise. when τοιdiziug that the iupact on a particular cosmolosical p:waleter typically depeuds logarithmically ou (.," In retrospect, this should perhaps not come as a complete surprise, when realizing that the impact on a particular cosmological parameter typically depends logarithmically on $\ell$."
+" For instance. Tamineche&Lewiρα(2008) considered a simple ]»ower spectruni model wit[um a single free amplitude. €,=qCE, and fouud that. for a eiven likelihood estimator to be “statistically unbiased. the systematic errors in that sale estimator nist fall off faster than ~1/f."," For instance, \citet{hamimeche:2008} considered a simple power spectrum model with a single free amplitude, $C_{\ell} = q\,C_{\ell}^{\textrm{fid}}$, and found that, for a given likelihood estimator to be “statistically unbiased”, the systematic errors in that same estimator must fall off faster than $\sim 1/\ell$."
+ A sinular consideration holds for ως., A similar consideration holds for $n_{\textrm{s}}$.
+ Iutuitivelv. i. is as much affected by (=2 ta 10 as it is between (=20 aud 100.," Intuitively, $n_{\textrm{s}}$ is as much affected by $\ell=2$ to 10 as it is between $\ell=20$ and 100."
+" Iu the previous 3-vear WALAP re-analysis paper. we increased the range of the accurate Likelihood estimator from (=12 to 30. corresponding to a factor of 2.5 in f. aud removed a bias of ~0.0 iu n,"," In the previous 3-year WMAP re-analysis paper, we increased the range of the accurate likelihood estimator from $\ell=12$ to 30, corresponding to a factor of 2.5 in $\ell$, and removed a bias of $\sim0.4\sigma$ in $n_{\textrm{s}}$."
+ In this paper. we increase the rauge from f=30 to 200. corresponding to a factor of 6.7 in ἐν and fiud i1 additional bias of 0.60.," In this paper, we increase the range from $\ell=30$ to 200, corresponding to a factor of 6.7 in $\ell$, and find an additional bias of $0.6\sigma$."
+ However. iucreasing £ from 3 ο 50 correspouds oulv to a factor of 1.7 in ἐν aud this appears to be too small to produce a statistically significant result.," However, increasing $\ell$ from 30 to 50 corresponds only to a factor of 1.7 in $\ell$, and this appears to be too small to produce a statistically significant result."
+ The nan conclusionsfrom this work are two-fold., The main conclusionsfrom this work are two-fold.
+ First. it secs that au accirate likelihood description is required to higher (s than previously believed. aud at least up to (= 200. in order to obtain unbiased results.," First, it seems that an accurate likelihood description is required to higher $\ell$ 's than previously believed, and at least up to $\ell=200$ , in order to obtain unbiased results."
+Within the framework of the SSC model the power emitted by IC is related (to (he power of the seed photons.,Within the framework of the SSC model the power emitted by IC is related to the power of the seed photons.
+ Photons of different energies that are emitted by the same particles should in principle produce similar pulse profiles., Photons of different energies that are emitted by the same particles should in principle produce similar pulse profiles.
+ In our model one expects. therefore. (hat ihe pulse profiles in X-ray and in gamma-rays are similar because (the secondary. plasma emits svuchrotron radiation in X-rays and IC scatters UV photons into the VUE band.," In our model one expects, therefore, that the pulse profiles in X-ray and in gamma-rays are similar because the secondary plasma emits synchrotron radiation in X-rays and IC scatters UV photons into the VHE band."
+ Ànd indeed. (he ratio of the amplitudes of the two pulses in the pulse profile of the Crab pulsar changes consistently in (he X-rays / soft eanuimacray band aid in the high energy eamima-ray band.," And indeed, the ratio of the amplitudes of the two pulses in the pulse profile of the Crab pulsar changes consistently in the X-rays / soft gamma-ray band and in the high energy gamma-ray band."
+ In. X-rays the main pulse dominates over (he inter pulse., In X-rays the main pulse dominates over the inter pulse.
+ The ratio changes towards higher energies and reverses in the soft gamma-ray band at about MMeV.; Similarly. the main pulse dominates at MMeV.(see.€.g..alsolorthepulseprofilesatlowerenergiesAbdo&etal.2010) while at GGeV. the inter pulse clearly dominates over the main pulse Alin&VERITASCollaboration(2011).," The ratio changes towards higher energies and reverses in the soft gamma-ray band at about MeV. Similarly, the main pulse dominates at MeV\citep[see, \eg also for the pulse profiles at lower energies][]{2010ApJS..187..460A} while at GeV the inter pulse clearly dominates over the main pulse \cite{VERITASPSRDetection}."
+. In addition. as we have argued that though IC losses may be energetically dominant (or similar to curvature emission). in the INN regime thev do not lead to a equilibrium distribution ol Lorentz factors.," In addition, as we have argued that though IC losses may be energetically dominant (or similar to curvature emission), in the KN regime they do not lead to a equilibrium distribution of Lorentz factors."
+ ence we expect highly non-stationary magnetospheric plasma lows., Hence we expect highly non-stationary magnetospheric plasma flows.
+ This will lead to highly non-stationary radiative properties., This will lead to highly non-stationary radiative properties.
+ Since within the SSC model the soft and hard photon fields are related. we might expect some 5-rav - X-ray correlation.," Since within the SSC model the soft and hard photon fields are related, we might expect some $\gamma$ -ray - $X$ -ray correlation."
+ Though ibis the soft UV. photons (hat are scattered (ο the GeV energies. and. formally. one expects UV-QGeV correlation. since X-ravs and UV form a continuous spectral distribution. one also expects .X-rav - GeV correlation as well.," Though it is the soft UV photons that are scattered to the GeV energies, and, formally, one expects UV-GeV correlation, since $X$ -rays and UV form a continuous spectral distribution, one also expects $X$ -ray - GeV correlation as well."
+ Thus we expect short (me-scale statistical correlation between X-ray and > rays photons., Thus we expect short time-scale statistical correlation between $X$ -ray and $\gamma$ rays photons.
+ llere we discuss the dependence of the 5-ray. Iuminosity on the spin-down power in the raciation-reaction limit., Here we discuss the dependence of the $\gamma$ -ray luminosity on the spin-down power in the radiation-reaction limit.
+ Generalizing Eq. (2.1)).," Generalizing Eq. \ref{01}) ),"
+ the total luminosity radiated by a primary beam of Goldreich-Julian density ie;=DO/(2xec) in the radiation reaction limited regime L.-— .- eciDnez, the total luminosity radiated by a primary beam of Goldreich-Julian density $n_{GJ}= {B \Omega /( 2 \pi e c)} $ in the radiation reaction limited regime L_c = e c B _G
+ the total luminosity radiated by a primary beam of Goldreich-Julian density ie;=DO/(2xec) in the radiation reaction limited regime L.-— .- eciDnezn, the total luminosity radiated by a primary beam of Goldreich-Julian density $n_{GJ}= {B \Omega /( 2 \pi e c)} $ in the radiation reaction limited regime L_c = e c B _G
+ the total luminosity radiated by a primary beam of Goldreich-Julian density ie;=DO/(2xec) in the radiation reaction limited regime L.-— .- eciDnezna, the total luminosity radiated by a primary beam of Goldreich-Julian density $n_{GJ}= {B \Omega /( 2 \pi e c)} $ in the radiation reaction limited regime L_c = e c B _G
+ ~100—104vr. ~20—200AU , $\sim 100-10^4 \yr$ $\sim 20-200 \AU$ 
+with steeper light profiles (larger helt concentration) and simaller sizes (see also Paper ID.,with steeper light profiles (larger light concentration) and smaller sizes (see also Paper II).
+ Iu Sevtert 2 hosts the light is less centrally concentrated aud the bulge component less prominent., In Seyfert 2 hosts the light is less centrally concentrated and the bulge component less prominent.
+ Iu addition. Sevtert 2 hosts tend to wave nore disturbed morplologics aud overall larger sizes.," In addition, Seyfert 2 hosts tend to have more disturbed morphologies and overall larger sizes."
+ These differences are sugeestive of an evolutionary connection between the two Seyfert types. that iiehlt be related o recent interactions/moergors.," These differences are suggestive of an evolutionary connection between the two Seyfert types, that might be related to recent interactions/mergers."
+ This sugeestion is consistent with our results in Paper TL indicating larger dust couteut aud disk star formation in Sevtert 2 galaxies.," This suggestion is consistent with our results in Paper II, indicating larger dust content and disk star formation in Seyfert 2 galaxies."
+ Iu. this respect. it is iutriguiug that the Wari Sevfert 2 xoperties are intermediate between those of Wari Severt ls and Cold ealaxies.," In this respect, it is intriguing that the Warm Seyfert 2 properties are intermediate between those of Warm Seyfert 1s and Cold galaxies."
+ The later reside iu ate-tvpe hosts of strongly interacting systems. have shallower ight profiles aud. as their »ptica and IR colours indicate. are probably dominated by strong disk star formaion (Paper ID).," The latter reside in late-type hosts of strongly interacting systems, have shallower light profiles and, as their optical and IR colours indicate, are probably dominated by strong disk star formation (Paper II)."
+ In Paper V we wil further explore the link between the interaction characteristics ancl he optical properties of our samples and will present evidence for a possible evolutionary scenario., In Paper V we will further explore the link between the interaction characteristics and the optical properties of our samples and will present evidence for a possible evolutionary scenario.
+ Before this. iu Paper IV. we will explore the colour distributious iu all our objects aud the inplications for their stellar conteut.," Before this, in Paper IV, we will explore the colour distributions in all our objects and the implications for their stellar content."
+ Tam grateful to mv thesis advisors George Mile and Walter Jaffe for providing me with stimulation and support throughout the completion of this project., I am grateful to my thesis advisors George Miley and Walter Jaffe for providing me with stimulation and support throughout the completion of this project.
+ This research has mace use of the NASA/TIPAC Extragalactic Database (NED) which is operated by he Jet Propulsion Laboratory. California Iustitute of Technology. wider contract with the National Aeronautics and Space Administration.," This research has made use of the NASA/IPAC Extragalactic Database (NED) which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration."
+ Part of this work was completed while the author held a National Research Council - NASA GSEC Research Associateship., Part of this work was completed while the author held a National Research Council - NASA GSFC Research Associateship.
+2001).. therefore. they are less reliable (han (ποσο based only upon spectroscopic information.,", therefore, they are less reliable than those based only upon spectroscopic information."
+ Groups of galaxies are very interesting objects in the universe. since {μον are the most common svstems of galaxies and (he evolution of galaxies inside them plavs an important role al early stages of cluster galaxies evolution.," Groups of galaxies are very interesting objects in the universe, since they are the most common systems of galaxies and the evolution of galaxies inside them plays an important role at early stages of cluster galaxies evolution."
+ How important is (he variation of galaxy luminosities with (he main physical properties of these svstems?, How important is the variation of galaxy luminosities with the main physical properties of these systems?
+ Iow does the environment allect the different galaxy populations?., How does the environment affect the different galaxy populations?.
+ Large statistical samples of galaxy groups are needed in order to these questions be addressed., Large statistical samples of galaxy groups are needed in order to these questions be addressed.
+ The largest galaxy redshift survey at the present is the Fourth Data Release of the Sloan Digital Skv Survev (hereafter DRA: Adelman-AeCarthyetal. 2005))., The largest galaxy redshift survey at the present is the Fourth Data Release of the Sloan Digital Sky Survey (hereafter DR4; \citealt{dr4}) ).
+ This catalog covers a very wide area on the sky. has hieh quality information in five broadhands and comprises two times (he number of galaxies in the current Final Release of the 21FGRS.," This catalog covers a very wide area on the sky, has high quality information in five broadbands and comprises two times the number of galaxies in the current Final Release of the 2dFGRS."
+ The huge size of the Sloan galaxy. catalog will provide us of the largest galaxy. group catalog at (he present with very important photometric and spectroscopic information., The huge size of the Sloan galaxy catalog will provide us of the largest galaxy group catalog at the present with very important photometric and spectroscopic information.
+ The main purpose of (this work is to compute the LFs of galaxies in groups under different. conditions in order to understand the behavior of galaxies in intermediate density environments and consequently provide important clues on galaxy. formation and evolution., The main purpose of this work is to compute the LFs of galaxies in groups under different conditions in order to understand the behavior of galaxies in intermediate density environments and consequently provide important clues on galaxy formation and evolution.
+ The lavout of this paper is as follows., The layout of this paper is as follows.
+ In section 2 we describe the galaxy sample aud the eroup identification., In section 2 we describe the galaxy sample and the group identification.
+ The full analvsis of LEs as a function of the band. mass ranges. color and the brightest galaxy is in section 3.," The full analysis of LFs as a function of the band, mass ranges, color and the brightest galaxy is in section 3."
+ Finally in section 4 we summarize and discuss the results., Finally in section 4 we summarize and discuss the results.
+ The Sloan Digital Sky Survey (5D55) has validated. aud made publicly available its Fourth Data Release (DIA: Adelman-MeCarthyetal. 2005)) which is a photometric and spectroscopic survey constructed with a dedicated 2.5m telescope al Apache Point Observatory in New Mexico., The Sloan Digital Sky Survey (SDSS) has validated and made publicly available its Fourth Data Release (DR4; \citealt{dr4}) ) which is a photometric and spectroscopic survey constructed with a dedicated $2.5 \ {\rm m}$ telescope at Apache Point Observatory in New Mexico.
+ The SDSS DIU consists of 6670deg? of five-band. ugri z. imaging data and 673.280 (4783 deg?) spectra of galaxies. quasars and stars;," The SDSS DR4 consists of $6670~ {\rm deg}^2$ of five-band, $u \ g \ r \ i \ z$ , imaging data and 673,280 $ 4783~ {\rm deg}^2$ ) spectra of galaxies, quasars and stars."
+ The DR4 Main Galaxy. Sample (MGS: Straussetal. 2002)) includes roughly 400.000 galaxies with redshift measurements up to ze0.3 and an upper apparent magnitude limit of 11.71 in the r baa.," The DR4 Main Galaxy Sample (MGS; \citealt{mgs}) ) includes roughly 400,000 galaxies with redshift measurements up to $z\sim0.3$ and an upper apparent magnitude limit of 17.77 in the $r$ band."
+ Galaxy groups used in (his work have been identilied in the MG of DR4 using the same procedure as in Alerchan&Zauclivarez(2005)., Galaxy groups used in this work have been identified in the MGS of DR4 using the same procedure as in \citet{mz05}.
+. The method consists in using a lrienc-ol- algorithm similar to that developed by Huclira&Geller(1982)., The method consists in using a friend-of-friend algorithm similar to that developed by \citet{hg82}.
+. The algorithun links galaxies(7. j) which satisIv that D;;<DyB(z) and Vj;€VyR(z) where Dj; is the projected," The algorithm links galaxies$i,j$ ) which satisfy that $D_{ij}\le D_0 \ R(z)$ and $V_{ij}\le V_0 \ R(z)$ where $D_{ij}$ is the projected"
+To explore the stability. of surface ice in more detail. we used a numerical heat transport model (Guilbert-Lepoutre et al.,"To explore the stability of surface ice in more detail, we used a numerical (three-dimensional) heat transport model (Guilbert-Lepoutre et al."
+ 2011) aud considered a rotating body having finite thermal conductivity. a realistic rotation period (6 hr) ancl three obliquities vec. 45° and 90°.," 2011) and considered a rotating body having finite thermal conductivity, a realistic rotation period (6 hr) and three obliquities $\psi =$ $\degr$, $\degr$ and $\degr$."
+ The surface thermal inertia was taken to be / = 15 MENS units (Mueller and Dlommaert 2004). emissivity € = 0.9 and we assumed that the spin vector al perihelion occupied a plane that also contained the Sun.," The surface thermal inertia was taken to be $I$ = 15 MKS units (Mueller and Blommaert 2004), emissivity $\epsilon$ = 0.9 and we assumed that the spin vector at perihelion occupied a plane that also contained the Sun."
+ Ice Dond albedos from 0.2 <Ax 0.8 were considered., Ice Bond albedos from 0.2 $\le A \le$ 0.8 were considered.
+ We computed the temperature on the surface as a Iunction of position and time., We computed the temperature on the surface as a function of position and time.
+ Sample solutions lor the peak annual temperature as a Dunction of latitude are shown in Figure (4))., Sample solutions for the peak annual temperature as a function of latitude are shown in Figure \ref{temp_vs_lat}) ).
+ The model results. plotted there refer to a spherical body with 90° obliquity moving in (he orbit of Cybele and with [our values of the ice albedo. as marked.," The model results plotted there refer to a spherical body with $\degr$ obliquity moving in the orbit of Cybele and with four values of the ice albedo, as marked."
+ For this large obliquity. the equator shows the lowest temperatures.," For this large obliquity, the equator shows the lowest temperatures."
+ For albedos 70.6. the equatorial reeions never rise above 131 Ix. the critical temperature for retaining surlace ice over (he vr collision interval.," For albedos $>$ 0.6, the equatorial regions never rise above 131 K, the critical temperature for retaining surface ice over the $^3$ yr collision interval."
+ In Figure (5)) we show the cumulative fraction of the surface whose peak temperature remains less (han a given value. 7.," In Figure \ref{cuplot_06}) ) we show the cumulative fraction of the surface whose peak temperature remains less than a given value, $T$."
+ For simplicity. we show only the results lor albedo = 0.6 and tabulate results for other albedos in Table (4)).," For simplicity, we show only the results for albedo = 0.6 and tabulate results for other albedos in Table \ref{conduction}) )."
+ A vertical dashed line in Figure (5)) shows the critical temperature 7 = 131 Ix (cL., A vertical dashed line in Figure \ref{cuplot_06}) ) shows the critical temperature $T$ = 131 K (c.f.
+ Figure (4))). below which surface ice can survive against sublimation for longer than the oE vr collision time.," Figure \ref{temp_vs_lat}) )), below which surface ice can survive against sublimation for longer than the $^3$ yr collision time."
+ For the particular case shown. αἱ ο = O°. almost of the surface remains al obliquely.Z7« 131 K all around the orbit. a result. of permanently cold polar regions iluminated ," For the particular case shown, at $\psi$ = $\degr$, almost of the surface remains at $T <$ 131 K all around the orbit, a result of permanently cold polar regions illuminated obliquely."
+Table (4)) shows even EL 4-02. these polar caps survive over a substantial action ol the surface (f~ na0.28).," Table \ref{conduction}) ) shows that, even for $A$ = 0.2, these polar caps survive over a substantial fraction of the surface $f \sim$ 0.28)."
+ ibtdAt high obliquity. οἳ = 907. the coldregion shrinks to ~20% of the surface and is instead as an equatorial band.," At high obliquity, $\psi$ = $\degr$, the cold region shrinks to $\sim$ of the surface and is distributed instead as an equatorial band."
+ The high obliquity models show the highest of all temperatures (at the poles) but retain cold nearly equatorial bands. as shown.," The high obliquity models show the highest of all temperatures (at the poles) but retain cold nearly equatorial bands, as shown."
+ For A> 0.6. these equatorial bands remain below 131 IX throughout the orbit aud so are potential regions for the preservation of surface ice on. 10* vr timescales.," For $A >$ 0.6, these equatorial bands remain below 131 K throughout the orbit and so are potential regions for the preservation of surface ice on $^3$ yr timescales."
+ The smallest. temperature variation over (he surface is found for intermediate obliquities (i— 45° i, The smallest temperature variation over the surface is found for intermediate obliquities $\psi =$ $\degr$ ).
+ For A€ 0.6 and this obliquity. no part of the surface is cold enough to retain ice for Q3 vr (Figure 5)).," For $A \le$ 0.6 and this obliquity, no part of the surface is cold enough to retain ice for $^3$ yr (Figure \ref{cuplot_06}) )."
+ Taken together. the numerical model results show that ice can persist for ~10 timescales on the surfaces of asteroids in Cvbele-like and Themis-like orbits. over a wide range of ice albedos ancl asteroid obliquities.," Taken together, the numerical model results show that ice can persist for $\sim$ $^3$ yr timescales on the surfaces of asteroids in Cybele-like and Themis-like orbits, over a wide range of ice albedos and asteroid obliquities."
+ 911all obliquiües strongly favor the long-term survival of ice in cold. polar regions. but ice can also survive near the equators of asteroids having obliquities ~90°.," Small obliquities strongly favor the long-term survival of ice in cold, polar regions, but ice can also survive near the equators of asteroids having obliquities $\sim$ $\degr$."
+ For albedos A 20.6. a millimeter-thick ice laver can survive for 10 vr on or more of the surface regardless of the oblieuitv.," For albedos $A >$ 0.6, a millimeter-thick ice layer can survive for $\sim$ $^3$ yr on or more of the surface regardless of the obliquity."
+ In this sense. the conduction models are compatible with the scenario proposed above. in which buried ice is excavated," In this sense, the conduction models are compatible with the scenario proposed above, in which buried ice is excavated"
+Planets are believed to form in the gas ancl cust discs that surround newly [ormecd stars.,Planets are believed to form in the gas and dust discs that surround newly formed stars.
+ In principle. planets can form purely from gravity. via direct gravitational instability in the initial circumstellar disc (Boss1997).," In principle, planets can form purely from gravity, via direct gravitational instability in the initial circumstellar disc \citep{Boss97}."
+. ILowever. this mechanism of planet formation requires higher initial densities (han commonly presumed.," However, this mechanism of planet formation requires higher initial densities than commonly presumed."
+ For a standard MMSN. (minimum, For a standard MMSN (minimum
+Iu eq. (12)). (153). (185)),"In eq. \ref{tcent_thin}) ), \ref{tcent_thick}) ), \ref{tcent_mir}) )"
+" and (19)). it is convenieut to express the optical depth 7, in ternis of the optical extinction ο”. Jos, where the (1ratio in parentheses can be computed for cach components with the parameters listed in Table B2 for the Osseukopf Wenning (1991) opacity."," and \ref{gamma_exp}) ), it is convenient to express the optical depth $\tau_a$ in terms of the optical extinction $A_V$, _a, where the ratio in parentheses can be computed for each components with the parameters listed in Table \ref{op_param} for the Ossenkopf Henning (1994) opacity."
+ We see that a cloud is optically thin to V-NIR radiation for Ay«3. to FIR radiation for dτοῦ. and to the cosinie background radiation for sly:G«104.," We see that a cloud is optically thin to V-NIR radiation for $A_V\ll 3$, to FIR radiation for $A_V\ll 700$, and to the cosmic background radiation for $A_V\ll 6\times 10^4$."
+" Thus. with Black (1991) standard mterstellar field. the central tempcrature in a ""diffuse"" cloud core of typical ceutre-to-edee extiuctiou Ayx=1 2 is determined by optically thin absorption of V-NIR aud AIR radiation (the effects of au external UV field are not considered here) whereas for 10Z4,z00 is cletermines bv optically thick absorption of V-NIR radiation aud optically thin absorption of MIR. aud FIR radiation."," Thus, with Black (1994) standard interstellar field, the central temperature in a “diffuse” cloud core of typical centre-to-edge extinction $A_V\simeq 1$ –2 is determined by optically thin absorption of V-NIR and MIR radiation (the effects of an external UV field are not considered here) whereas for $10\simless A_V \simless 400$ is determined by optically thick absorption of V-NIR radiation and optically thin absorption of MIR and FIR radiation."
+" TusertingC» the numerical values Ooeiveu in Appeucix D (and keeping terms up to second order in τι). the formmlac above give Χοκςeee. gpyulisJine BA,” TTA,EM 6.2.(QOO3LAAoo... where 7; is in K. Combining eq. (21)) (21))."," Inserting the numerical values given in Appendix B (and keeping terms up to second order in $\tau_a$ ), the formulae above give ), ) 43 -77, ) 6.2-0.0031, where $T_d$ is in $K$ Combining eq. \ref{vnirthin}) \ref{mirthin}) ),"
+ one can obtain simple analyticalexpressions for the dus telpcrature (im Jv) at the clouds centre. for Ay~1 2. and for LOSAy-thin)”=LOO.," one can obtain simple analyticalexpressions for the dust temperature (in $K$ ) at the cloud's centre, , for $A_V\simeq 1$ –2, and for $10 \simless A_V\simless 400$."
+ Noticο thatincreasing the iutensity of the radiation fick by a facor g results in an increase of the dust temperature by a factor y1/56, Notice that increasing the intensity of the radiation field by a factor $g$ results in an increase of the dust temperature by a factor $g^{1/5.6}$.
+ We have checked these approximate formulae against the uunerical comptations aud the results are shown in Figure2.. Tere predict, We have checked these approximate formulae against the numerical computations and the results are shown in Figure \ref{analytic_temp}.
+ionswe compare in the first place the πιο]σα] model for the temperature a the center of a DounorEbert sphere with tha derived from eq. (25)), Here we compare in the first place the numerical model predictions for the temperature at the center of a Bonnor–Ebert sphere with that derived from eq. \ref{tdlow}) )
+ and 26))., and \ref{tdhigh}) ).
+ We see that tje analytic formula is m agreement with the more accurate muuerical results within ~1 Is. CET isladeyy e£ iuterest to compare the expectations from the analytic formula wih that calculated as a function of radius in a given Donnor-Ebert sphere., We see that the analytic formula is in agreement with the more accurate numerical results within $\sim 1$ K. It is also of interest to compare the expectations from the analytic formula with that calculated as a function of radius in a given Bonnor-Ebert sphere.
+" Here. we merely use the ""iutuitive approximation that the appropriate optical depth to substitute iu eq. (26))"," Here, we merely use the “intuitive approximation” that the appropriate optical depth to substitute in eq. \ref{tdhigh}) )"
+ is the mami opticalepth along a ray to the exterior of the core., is the minimum optical depth along a ray to the exterior of the core.
+ We see in Fie., We see in Fig.
+ 3 that this approximation is also quite reasonable sugecsting that. in a nunber of situations. oue can obtain an estimate of the dust temperature as a function of position σας eq. (26)).," \ref{tanalytic} that this approximation is also quite reasonable suggesting that, in a number of situations, one can obtain an estimate of the dust temperature as a function of position using eq. \ref{tdhigh}) )."
+ Obviously however. this ust be Tuxused with caution.," Obviously however, this must be used with caution."
+" Nevertheless. it should as a rule be ""ure accurate than the isothermal assumption."," Nevertheless, it should as a rule be more accurate than the isothermal assumption."
+ Iu this section. we prescut our results obtained numerically inteerating eq. (1))," In this section, we present our results obtained numerically integrating eq. \ref{therm_eq}) )"
+ aud (2))., and \ref{Jrdef}) ).
+ We first. cousider inhomogeneous spherically sviuiuetric density distributious., We first consider inhomogeneous spherically symmetric density distributions.
+ Two obvious cases of interest are the singular isothermal sphere (SIS) solution aud the Bounor- sphere (BES)., Two obvious cases of interest are the singular isothermal sphere (SIS) solution and the Bonnor-Ebert sphere (BES).
+ In Fie. L. ," In Fig.\ref{sisbeb}, ,"
+weshow the results we have obtained for these density distributions., weshow the results we have obtained for these density distributions.
+ Both models are obtained from the equation of vdrostatic equilibrium, Both models are obtained from the equation of hydrostatic equilibrium
+"The closest optical/UV counterpart for X4 is WOCS ID 4003, located at a distance of 17557 (1.30) from the X- position.","The closest optical/UV counterpart for X4 is WOCS ID 4003, located at a distance of 57 $\sigma$ ) from the X-ray position."
+" Radial-velocity studies of WOCS 4003 show it is not a short-period (P <3000 days) binary, meaning it could be a single star or a long period binary."," Radial-velocity studies of WOCS 4003 show it is not a short-period (P $< 3000$ days) binary, meaning it could be a single star or a long period binary."
+" WOCS 4003 has a low probability of being a cluster member from its radial-velocity, and proper-motion studies indicate probability that it is a member (Holeetal.(2009);; I. Platais, private communication)."," WOCS 4003 has a low probability of being a cluster member from its radial-velocity, and proper-motion studies indicate probability that it is a member \citet{Hole09}; I. Platais, private communication)."
+" This object falls near the intersection of all three populations in Figure 6 and in the area dominated by active binaries in Figure 7,, if WOCS 4003 is the true optical counterpart and assuming it is at the distance of NGC 6819."," This object falls near the intersection of all three populations in Figure \ref{xcmd} and in the area dominated by active binaries in Figure \ref{xopt}, if WOCS 4003 is the true optical counterpart and assuming it is at the distance of NGC 6819."
+" Since cluster membership is highly unlikely, the luminosity given is almost certainly incorrect."," Since cluster membership is highly unlikely, the luminosity given is almost certainly incorrect."
+" With only this information, the source of the X-ray emission for X4 is still unknown."," With only this information, the source of the X-ray emission for X4 is still unknown."
+" Without a previous systematic study of the X-ray sources in many open clusters, it is difficult to predict the types and number of sources we expect to find."," Without a previous systematic study of the X-ray sources in many open clusters, it is difficult to predict the types and number of sources we expect to find."
+" If we extrapolate the relation between the number of X-ray sources and dynamical frequency in globular clusters to the regime of open clusters, we would expect almost no sources (Pooley&Hut2006,seebelow).."," If we extrapolate the relation between the number of X-ray sources and dynamical frequency in globular clusters to the regime of open clusters, we would expect almost no sources \citep[][see below]{Pooley06}."
+ Studies of two other individual open clusters have shown this not to be the case (vandenBergetal.2004;Gon-doin2005) but little is known about the formation of such populations in open cluster environments.," Studies of two other individual open clusters have shown this not to be the case \citep{vandenBerg04,Gondoin05} but little is known about the formation of such populations in open cluster environments."
+" Here we make some reasonable assumptions about the formation of some of our sources, but definitive conclusions await accurate N-body modeling of NGC 6819 and systematic analyses of a sample of open cluster X-ray populations."," Here we make some reasonable assumptions about the formation of some of our sources, but definitive conclusions await accurate $N$ -body modeling of NGC 6819 and systematic analyses of a sample of open cluster X-ray populations."
+" The presence of a qLMXB in an open cluster may be unexpected, and could be assumed to be dynamical given the known relationship between dynamics and qLMXB populations in globular clusters (Pooley&Hut 2006)."," The presence of a qLMXB in an open cluster may be unexpected, and could be assumed to be dynamical given the known relationship between dynamics and qLMXB populations in globular clusters \citep{Pooley06}."
+". Although there are currently studies underway to determine the spatial density of qLMXBs in the field (Grindlayetal.2005),, without such a measurement we are unable to determine how many primordial qLMXBs would be expected in an open cluster using field values."," Although there are currently studies underway to determine the spatial density of qLMXBs in the field \citep{Grindlay05}, without such a measurement we are unable to determine how many primordial qLMXBs would be expected in an open cluster using field values."
+" Pooley&Hut(2006) investigate X-ray sources in globular clusters, fitting a two-component model to nz of X-ray sources per unit mass) vs. (encounter frequency(number per unit mass) that separates the ycontribution from primordial and dynamical populations."," \citet{Pooley06} investigate X-ray sources in globular clusters, fitting a two-component model to $n_{x}$ (number of X-ray sources per unit mass) vs. $\gamma$ (encounter frequency per unit mass) that separates the contribution from primordial and dynamical populations."
+" The model is applied to population I, II, and III objects separately as well as to all X-ray objects combined."," The model is applied to population I, II, and III objects separately as well as to all X-ray objects combined."
+" For population I objects (dominated by qLMXBs), the best fit to this relationship results in 0.41023 primordial sources per 106 Mo."," For population I objects (dominated by qLMXBs), the best fit to this relationship results in $0.4^{+0.53}_{-0.55}$ primordial sources per $10^{6}$ $_{\odot}$."
+" Although this is a poorly constrained value, we apply it here to roughly determine the probability of finding à primordial population I object in NGC 6819."," Although this is a poorly constrained value, we apply it here to roughly determine the probability of finding a primordial population I object in NGC 6819."
+ Kaliraietal.(2001b) integrate the mass function of NGC 6819 to find a cluster mass of ~2600 Mo., \citet{Kalirai01b} integrate the mass function of NGC 6819 to find a cluster mass of $\sim$ 2600 $_{\odot}$.
+" Using this value with the upper limit for primordial X-ray sources expected per unit mass, we expect ~0.002 primordial population I objects in NGC 6819."," Using this value with the upper limit for primordial X-ray sources expected per unit mass, we expect $\sim$ 0.002 primordial population I objects in NGC 6819."
+ The low likelihood of finding a primordial qLMXB in NGC 6819 given this relationship would seem to indicate a dynamical history for X1., The low likelihood of finding a primordial qLMXB in NGC 6819 given this relationship would seem to indicate a dynamical history for X1.
+" We apply this comparison with caution, noting that the validity of extrapolating globular cluster characteristics down to the regime of open clusters is dependent, in part, on whether the age difference between globular and open clusters affects the observable X-ray population."," We apply this comparison with caution, noting that the validity of extrapolating globular cluster characteristics down to the regime of open clusters is dependent, in part, on whether the age difference between globular and open clusters affects the observable X-ray population."
+" Ivanovaetal.(2008) explore the number of qLMXBs in globular clusters over time, finding that the number of observable qLMXBs from 4 to 11 Gyr remains relatively constant for the majority of globular clusters, assuming that all qLMXBs that appear before 4 Gyr are primordial in nature."," \citet{Ivanova08} explore the number of qLMXBs in globular clusters over time, finding that the number of observable qLMXBs from 4 to 11 Gyr remains relatively constant for the majority of globular clusters, assuming that all qLMXBs that appear before 4 Gyr are primordial in nature."
+" This interesting result would seem to indicate that age does not play a major role in the number of observable qLMXBs, but the initial assumption that all sources prior to 4 Gyr are primordial limits the ability to confidentlyextrapolate their population characteristics back to the age of NGC 6819."," This interesting result would seem to indicate that age does not play a major role in the number of observable qLMXBs, but the initial assumption that all sources prior to 4 Gyr are primordial limits the ability to confidentlyextrapolate their population characteristics back to the age of NGC 6819."
+" Regardless of the formation mechanism, the presence of a qLMXB is also unexpected given the results of Hurleyetal.(2005) that found with their N-body model only two neutron stars remaining in M67 at 4 Gyr, both of which are single stars."," Regardless of the formation mechanism, the presence of a qLMXB is also unexpected given the results of \citet{Hurley05} that found with their $N$ -body model only two neutron stars remaining in M67 at 4 Gyr, both of which are single stars."
+" Given the comparison tools currently available that all point to the unlikelihood of creating a primordial qLMXB in open cluster environment, we tentatively interpret X1 to be a dynamically formed candidate qLMXB."," Given the comparison tools currently available that all point to the unlikelihood of creating a primordial qLMXB in open cluster environment, we tentatively interpret X1 to be a dynamically formed candidate qLMXB."
+ Predicting the number of expected CVs in open clusters is similarly troublesome., Predicting the number of expected CVs in open clusters is similarly troublesome.
+" Although the CV populations of globular clusters and the field have been well studied, it is unknown whether either population is strictly applicable to rich open clusters at 2 Gyr."," Although the CV populations of globular clusters and the field have been well studied, it is unknown whether either population is strictly applicable to rich open clusters at 2 Gyr."
+ We extrapolate the relationship for primordial CVs in globular clusters down to the masses of open clusters., We extrapolate the relationship for primordial CVs in globular clusters down to the masses of open clusters.
+ Pooley&Hut(2006) find a best fit of 0.27*139 primordial population II objects (dominated by CVs) per 10° Mo., \citet{Pooley06} find a best fit of $0.27^{+1.40}_{-1.99}$ primordial population II objects (dominated by CVs) per $10^{6}$ $_{\odot}$.
+" At the upper limit, this results in ~0.004 primordial CVs in NGC 6819."," At the upper limit, this results in $\sim$ 0.004 primordial CVs in NGC 6819."
+ A more suitable comparison tool may be the field CV density., A more suitable comparison tool may be the field CV density.
+" The ChamPLane survey finds a local CV space density of 0.9*05x107? pe~? at confidence, which is consistent with models of the galactic CV population (Rogeletal."," The ChamPLane survey finds a local CV space density of $0.9^{+1.5}_{-0.5}\times10^{-5}$ $^{-3}$ at confidence, which is consistent with models of the galactic CV population \citep{Rogel08}."
+ This value is also consistent within two standard deviations2008).. to the local CV space density of Patterson(1998) at all scale heights., This value is also consistent within two standard deviations to the local CV space density of \citet{Patterson98} at all scale heights.
+" Given the local mass density of the galactic plane of 0.1 Me pc-? (Kuijken&Gilmore 1989),, this results in a CV mass density of 9.0155x1075 CVs Mo’."," Given the local mass density of the galactic plane of 0.1 $_{\odot}$ $^{-3}$ \citep{Kuijken89}, this results in a CV mass density of $9.0^{+15}_{-5}\times10^{-5}$ CVs $_{\odot}^{-1}$."
+ For the mass of NGC 6819 this results in a prediction of primordial CVs in the cluster., For the mass of NGC 6819 this results in a prediction of $0.23^{+0.39}_{-0.13}$ primordial CVs in the cluster.
+" The upper limit of this 0.230prediction,14 0.62 CVs in NGC 6819, is not inconsistent with a detection of one CV in the cluster."," The upper limit of this prediction, 0.62 CVs in NGC 6819, is not inconsistent with a detection of one CV in the cluster."
+" Hurleyetal.(2005) did create one primordial CV in M67 at a cluster mass of only 1400 Mo, which may point to primordial CVs being more prevalent in open clusters than either of these comparisons predict."," \citet{Hurley05} did create one primordial CV in M67 at a cluster mass of only 1400 $_{\odot}$, which may point to primordial CVs being more prevalent in open clusters than either of these comparisons predict."
+" The possible formation scenario for the optical counterpart of X6, outlined in refsec:activebin,, may point to a dynamical history, but requires more observations and modeling to definitively determine the formation mechanism."," The possible formation scenario for the optical counterpart of X6, outlined in \\ref{sec:activebin}, may point to a dynamical history, but requires more observations and modeling to definitively determine the formation mechanism."
+" In addition, the possible sub-subgiant counterpart to X9 cannot be explained by single-star evolutionary theory."," In addition, the possible sub-subgiant counterpart to X9 cannot be explained by single-star evolutionary theory."
+" Possible explanations require under- or over-luminous binary components that may be formed through mass transfer, stellar mergers, and/or dynamical stellar exchanges (Mathieuetal. 2003).."," Possible explanations require under- or over-luminous binary components that may be formed through mass transfer, stellar mergers, and/or dynamical stellar exchanges \citep{Mathieu03}. ."
+" Thus, while not yet definitive, several X-ray sources in NGC 6819 hint at dynamical formation processes."," Thus, while not yet definitive, several X-ray sources in NGC 6819 hint at dynamical formation processes."
+ The presence of dynamically formed, The presence of dynamically formed
+dispersion <80%). with the star formation rate consider as the ultimate scaling [actor for both Ih. and non-thermal radio huninosities (see Thompson et al.,"dispersion $\la
+80\%$ ), with the star formation rate consider as the ultimate scaling factor for both IR and non-thermal radio luminosities (see Thompson et al."
+ 2006 for the latest theoretical background)., 2006 for the latest theoretical background).
+ To obtain a working value for (7) we follow Yun et al., To obtain a working value for $\mu(\nu)$ we follow Yun et al.
+" 2001 and use their value for the far-IR/radio correlation at r—1.4 GGIIz: (pL,)/Lyy,=2.66 x10x10.9. where q=2.35 (as defined by Helou et al."," 2001 and use their value for the far-IR/radio correlation at $\nu $ GHz: $\rm \left(\nu
+L_{\nu}\right)/L_{FIR}$ $\rm 2.66\times 10^{-4}\times 10^{-q}$, where q=2.35 (as defined by Helou et al."
+ 1935)., 1985).
+ For (Ljs(8. jm)/Lqa45(42. jm)) —1.3 (Yun et al.," For $\langle \rm
+L_{IR}$ $\mu $ $\rm L_{FIR}$ $\mu $ $\rangle$ =1.3 (Yun et al."
+ 2001). the coefficient in Equation 2 becomes jr(p)—1.54. and is adopted for this study (wil an inherent uncertainty of a factor of ~2 due to a varving Lygi/Lii im galaxies).," 2001), the coefficient in Equation 2 becomes $\rm \mu
+(\nu)$ =1.54, and is adopted for this study (with an inherent uncertainty of a factor of $\sim $ 2 due to a varying $\rm
+L_{FIR}/L_{IR}$ in IR-luminous galaxies)."
+ On the other hand (he non-thermal radio power depends on (he SN rate [ax as (Condon 1992)., On the other hand the non-thermal radio power depends on the SN rate $\rm f_{SN}$ as (Condon 1992).
+ Combining Equations 2 and 3 and assuming full concomitance of LR and non-thermal radio emission over the star forming regions vields the SN rate surface density Following Suckhov et al. (, Combining Equations 2 and 3 and assuming full concomitance of IR and non-thermal radio emission over the star forming regions yields the SN rate surface density Following Suckhov et al. (
+1993) the CR energy densities in ULIRGs scale with respect to that of the Galaxy as where Xa c13.85  tkkpe 7 (Melee Williams 1997). and V qiiis the diffusion velocily al which CRs escape from quiescent disks like the Milky Way. while Ving is the velocily of a SE-induced wind at which CRs are advected out of the SF regions in sturbursts.,"1993) the CR energy densities in ULIRGs scale with respect to that of the Galaxy as where $\rm \dot\Sigma_{SN,Gal}$ $\times
+$ $^{-5}$ $^{-1}$ $^{-2}$ (McKee Williams 1997), and $\rm
+V_{diff}$ is the diffusion velocity at which CRs escape from quiescent disks like the Milky Way while $\rm V_{wind}$ is the velocity of a SF-induced wind at which CRs are advected out of the SF regions in starbursts."
+ Typically VayerlOkkmss + while the maximum velocity of SF-cdriven galactic winds is Vdndansv x10* | (Suchkov et a. L993: Veilleux et al.," Typically $\rm V_{diff}$$\sim $ $^{-1}$ while the maximum velocity of SF-driven galactic winds is $\rm
+V_{wind,max}$$\sim $ $\times $ $^3$ $^{-1}$ (Suchkov et a. 1993; Veilleux et al."
+ 2005)., 2005).
+ Such high wind velocities are deduced [or 582 (Seaquist οἱ al., Such starburst-induced high wind velocities are deduced for 82 (Seaquist et al.
+ 1985) though usually μασ !., 1985) though usually $\rm V_{wind}$$\sim $ $^{-1}$.
+ From Equations 4 and 5. after setting Viindmas=3000 !. and νου 1.54. a=0.8 (for p—1.4 GGIIz). Equation | vields CR energy density in compact starbursts (CSB) of," From Equations 4 and 5, after setting $\rm
+V_{wind,max}$ $^{-1}$, and $\mu (\nu)$$\sim $ 1.54, $\alpha
+$ =0.8 (for $\nu$ GHz), Equation 1 yields CR energy density in compact starbursts (CSB) of"
+The four velocity of the source is given by: where Vis the 2-direction velocity relative to the locally uou-rotating inertial frame.,The four velocity of the source is given by: where $V_{z}$ is the $z$ -direction velocity relative to the locally non-rotating inertial frame.
+" Iu the locally now-rotating inertial frame. the velocity components are: Tf the N-rav flare motion is dominated by motion in the r-cdirection (sav. outward or inward). theu the oulv uou-zero velocity component in the locally uon-rotating inertial frame is 2,=VL."," In the locally non-rotating inertial frame, the velocity components are: If the X-ray flare motion is dominated by motion in the $r$ -direction (say, outward or inward), then the only non-zero velocity component in the locally non-rotating inertial frame is $\beta_{r}=V_{r}$."
+" It is also possible that the flare co-rotates with the disk while other direction motion can be neglected G.e. 3,=390).", It is also possible that the flare co-rotates with the disk while other direction motion can be neglected (i.e. $\beta_{r}=\beta_{\theta}=0$ ).
+ The motion constants in (ALL)) ave then equivalent to the ones given by Ruszkowski(2000)., The motion constants in \ref{eq:system9}) ) are then equivalent to the ones given by \citet{rus00}.
+". Note that any choice of the value of V. or V, ust mect the requirement of V<1 in order that the source follows a time-like world-liue.", Note that any choice of the value of $V_z$ or $V_{r}$ must meet the requirement of $V^2<1$ in order that the source follows a time-like world-line.
+" For the special case that the flare is on axis aud has only ¢-/+-direction motion (ic. 0.>0 and 0). the two motion coustauts reduce to: We asstune the X-ray huuinositv iu the source comoving frame to be isotropic aud iu the form of a power lav Leo=Lyk"". where E. is the photon euergy in the source comoving frame."," For the special case that the flare is on axis and has only $r$ $z$ -direction motion (i.e. $\theta \rightarrow 0$ and $\beta_{\varphi}=\beta_{\theta}=0$ ), the two motion constants reduce to: We assume the X-ray luminosity in the source comoving frame to be isotropic and in the form of a power law $L_{E_e}=L_0 E^{-\alpha}_{\rm e}$, where $E_{\rm e}$ is the photon energy in the source comoving frame."
+ The simulation shows that the anisotropic effect is very weak for multiple scattering which dominates the hard N-ray cussion. although the soft N-rav. emission is significantly anisotropic in the case of single scattering approximation (Janis. Czeruv Zvveki20003.," The Monte-Carlo simulation shows that the anisotropic effect is very weak for multiple scattering which dominates the hard X-ray emission, although the soft X-ray emission is significantly anisotropic in the case of single scattering approximation (Janiuk, Czerny Żyycki2000)."
+" Tn the frame rotating with the disk. chiding the ‘k-correctiou. the illuminating fiux on the disk clement (r>r|dr. ye>| dy) with an incident ang5 of On, (cosÜu,— p-:n/p-uqa. where n is the equator normal) is given by (e.g. Yu Lu 2000: Ruszkowski 2000): where fiq is the fraction of total photons which corresponds to the eid. clemeuts defined by dqdr. qa is the redshift factor for a photon propagating from the source to the disk erid elements eiven by and 5 is the Loreutz factor of the relative motion of the disk clement aud the locally non-rotating inertial fraane."," In the frame rotating with the disk, including the `k-correction', the illuminating flux on the disk element $r\rightarrow r+dr$, $\varphi
+\rightarrow \varphi+d\varphi$ ) with an incident angle of $\theta_{\rm in}$ $\cos\theta_{\rm in}=-{\bf p}\cdot{\bf n}/{\bf p}\cdot{\bf u_{\rm d}}$, where ${\bf n}$ is the equator normal) is given by (e.g. Yu Lu 2000; Ruszkowski 2000): where $f_{\rm sd}$ is the fraction of total photons which corresponds to the grid elements defined by $d\varphi dr$, $g_{\rm sd}$ is the redshift factor for a photon propagating from the source to the disk grid elements given by and $\gamma$ is the Lorentz factor of the relative motion of the disk element and the locally non-rotating inertial frame."
+ For &2 ry. Where ry. is the radius of the last stable orbit. the 7 factor is given by:," For $r\geq r_{\rm ms}$ , where $r_{\rm ms}$ is the radius of the last stable orbit, the $\gamma$ factor is given by:"
+More than fifty plaucts have now boeon icentifiec outside the Solar svstein (που for a sunanuurv oftιο known planetary svstenis).,More than fifty planets have now been identified outside the Solar system (see for a summary of the known planetary systems).
+ While the majority of thex5 rave been foind via the measurement of a stellar radia velocity changes due to the presence of au orbiting planet. nore recently the dimunine of stellar ligl due to the transi of a giaut planet has been observed. (Charbounecan et al.," While the majority of these have been found via the measurement of a stellar radial velocity changes due to the presence of an orbiting planet, more recently the dimming of stellar light due to the transit of a giant planet has been observed, (Charbonneau et al."
+ 2000: Παν e al., 2000; Henry et al.
+ 2000: Queloz et a., 2000; Queloz et al.
+ 200: Castellauno e al., 2000; Castellano et al.
+ 2000: Dovle et al., 2000; Doyle et al.
+ 2000: Caucli 2100)ο, 2000; Gaudi 2000).
+ ηicalv. lich curves are seen to dip by 124 as the pauet transits the star. correspoxdiug to planetary rac lo 0.10(WLLLRy.," Typically, light curves are seen to dip by $1-2\%$ as the planet transits the star, corresponding to planetary radii of $0.10-0.14\ R_s$."
+" Cüven current observational limits. hese teclunicWes Call oulv be 1:sed to identify planets within a few tens of 2111ΜΟΝ,"," Given current observational limits, these techniques can only be used to identify planets within a few tens of parsecs."
+ Gravitational uicroleusius has the poteutia to discover panets owy a nmmchl ereater distance., Gravitational microlensing has the potential to discover planets over a much greater distance.
+ Perturbiug he maguificatiou cistributiou. aud lence the form of the nicroleusiie liehlit cmrve. planets can be ideutiBed orbiting colupact oyjects iu the Calactic Talo (Mao Daczyvüsski 1901: Coud Loc» 1992: Bolatto Falco 199I: Deunett Rhic 1996: Wiuuewnss 1997).," Perturbing the magnification distribution, and hence the form of the microlensing light curve, planets can be identified orbiting compact objects in the Galactic Halo (Mao Paczyńsski 1991; Gould Loeb 1992; Bolatto Falco 1994; Bennett Rhie 1996; Wambsganss 1997)."
+ More receulv. Craft Cardi (2000) aid Lewis Ibata (2000) turned ier attention to usi18o eravitational microleusimg to ideutifv ylanets orbiting sas in the Calactic Dueo. that is detecting plancts orbiting the ina uücroleusius event. rather tiui orbiting the ens.," More recently, Graff Gaudi (2000) and Lewis Ibata (2000) turned their attention to using gravitational microlensing to identify planets orbiting stars in the Galactic Bulge, that is detecting planets orbiting the in a microlensing event, rather than orbiting the lens."
+ Asstmine that he star plus planet system is swept bv a fold carstic formed bv a bimary leus. it was shown that the feeble ight that is refiected from the plaret (upte)LO+L.) can be imaguified to observable levels. resulting iu a OAL~ O.03inag fluctuation.," Assuming that the star plus planet system is swept by a fold caustic formed by a binary lens, it was shown that the feeble light that is reflected from the planet ${\rm (up~ to
+\sim10^{-4}L_*)}$ can be magnified to observable levels, resulting in a $\delta M\sim0.03$ mag fluctuation."
+" The scattering responsible or reflecting the stellar light from, t16 planet also acts to yolarize it and Lewis Ibata (2000) demoustrated that his signature too is boosted to observable levels. probing he physical conditious iu the planetary atinosphere."," The scattering responsible for reflecting the stellar light from the planet also acts to polarize it and Lewis Ibata (2000) demonstrated that this signature too is boosted to observable levels, probing the physical conditions in the planetary atmosphere."
+" Following these initial studies. Ashtou Lewis (2001) considered t1e influence of the plauetary oase ou the form of the wicrolensing lielt curve. finding tiat when the planet appears cresceut-like. the naenification can be substatially ereater than the sinple. circularly saunietric uxxdel that was euiploved in hese earier studies,"," Following these initial studies, Ashton Lewis (2001) considered the influence of the planetary phase on the form of the microlensing light curve, finding that when the planet appears crescent-like, the magnification can be substantially greater than the simple, circularly symmetric model that was employed in these earlier studies."
+" This oper also focuses ou the identification of planets orbitiis stars that are the sources durus eravitatioial wicrocusing evens,", This paper also focuses on the identification of planets orbiting stars that are the sources during gravitational microlensing events.
+ Instead o: lookiug at the refleced ight. however. this paper examines Section 7? compares the results of this study with previous iuvestigations of the influence of spots on he stellar surface. as well as investigating the frequency of planetary transits diving nicroleusime events.," Instead of looking at the reflected light, however, this paper examines Section \ref{discussion} compares the results of this study with previous investigations of the influence of spots on the stellar surface, as well as investigating the frequency of planetary transits during microlensing events."
+ Section preseuts the couchisious of this study., Section \ref{conclusions} presents the conclusions of this study.
+ A wealth of information las resulted from the rapid erowth iu the field of Calactic and Local Group eravitational microlensing over the last decade., A wealth of information has resulted from the rapid growth in the field of Galactic and Local Group gravitational microlensing over the last decade.
+ The reader is davected to the review of Paczvüsski (1996) which covers lis backeround material iu detail., The reader is directed to the review of Paczyńsski (1996) which covers this background material in detail.
+ Most stucies of. Galactic ido nmücroleusing lave ocused: iponu unicrolensing by isolated compact objects. which display simple. bell-shapec light curves.," Most studies of Galactic halo microlensing have focused upon microlensing by isolated compact objects, which display simple, bell-shaped light curves."
+ A fraction of events. however. displav several rapid. asvnuiuetric eatures which are characteristic of nücrolensiug by a jnary svstem (Alcock et al.," A fraction of events, however, display several rapid, asymmetric features which are characteristic of microlensing by a binary system (Alcock et al."
+ 2000).," 2000),"
+the halo size and changes of environment with distance from the galaxy.,the halo size and changes of environment with distance from the galaxy.
+" Finally in addition to environment being defined by galaxy positions within the volume, we also measure environment as inferred from the background dark matter distribution."," Finally, in addition to environment being defined by galaxy positions within the volume, we also measure environment as inferred from the background dark matter distribution."
+ To obtain the neighbourhood dark matter environment in the Millennium Simulation the full volume is broken into a three dimensional grid with side-length 2h-'Mpc.," To obtain the neighbourhood dark matter environment in the Millennium Simulation the full volume is broken into a three dimensional grid with side-length $2\,h^{-1} {\rm Mpc}$."
+" At the centre of each grid element a three dimensional Gaussian density is calculated using the local dark matter particles, smoothed over three different scales: 2.5, 5, and 10h!Mpc."," At the centre of each grid element a three dimensional Gaussian density is calculated using the local dark matter particles, smoothed over three different scales: $2.5$, $5$, and $10\,h^{-1} {\rm Mpc}$."
+" This Gaussian smoothed density is similar to the kernel smoothed density of Equation 7,, but with a dark matter particle mass term in the sum."," This Gaussian smoothed density is similar to the kernel smoothed density of Equation \ref{eq:sphdens}, but with a dark matter particle mass term in the sum."
+ In Section 4 we apply a number of fixed aperture methods to the mock galaxy catalogue and measure local density around each galaxy., In Section \ref{sec:res} we apply a number of fixed aperture methods to the mock galaxy catalogue and measure local density around each galaxy.
+" This allows us to quantify the properties that aperture measured densities best probe, and compare with the previously described nearest neighbour estimators."," This allows us to quantify the properties that aperture measured densities best probe, and compare with the previously described nearest neighbour estimators."
+" To investigate the different properties of each galaxy environment measure, in this section we consider how they correlate with (1) the host dark matter halo mass, (2) the underlying dark matter environment, and (3) the colour of the galaxies."," To investigate the different properties of each galaxy environment measure, in this section we consider how they correlate with (1) the host dark matter halo mass, (2) the underlying dark matter environment, and (3) the colour of the galaxies."
+ 'To facilitate this we have converted the output of each to a ‘percentage rank’ for each galaxy., To facilitate this we have converted the output of each to a `percentage rank' for each galaxy.
+" This is computed by listing the galaxies in order of increasing density, then assigning them a percentage based on where they appear in that list, with zero percent being the least dense and one hundred percent the most dense."," This is computed by listing the galaxies in order of increasing density, then assigning them a percentage based on where they appear in that list, with zero percent being the least dense and one hundred percent the most dense."
+" Therefore, a galaxy with a percentage rank of ninety five has five percent of the galaxies in the sample denser than it and ninety five percent less dense than it."," Therefore, a galaxy with a percentage rank of ninety five has five percent of the galaxies in the sample denser than it and ninety five percent less dense than it."
+" This normalisation provides a fairer comparison between environment estimators and probes their relative rather than absolute distributions across the environment spectrum, which would otherwise be definition dependent."," This normalisation provides a fairer comparison between environment estimators and probes their relative rather than absolute distributions across the environment spectrum, which would otherwise be definition dependent."
+" By design, the most fundamental property for a galaxy within our model is its dark matter halo mass."," By design, the most fundamental property for a galaxy within our model is its dark matter halo mass."
+ Halo mass determines both the spatial distribution of the galaxy population and the individual galaxy properties., Halo mass determines both the spatial distribution of the galaxy population and the individual galaxy properties.
+" Therefore, each environment measure should reveal some underlying correlation."," Therefore, each environment measure should reveal some underlying correlation."
+" Typically halo masses of 10121Μο correspond to the field, ~10!?54-!M to groups and ~101?A7!Mo to clusters."," Typically halo masses of $\sim 10^{12}\,h^{-1} M_{\odot}$ correspond to the field, $\sim 10^{13.5}\,h^{-1} M_{\odot}$ to groups and $\sim 10^{15}\,h^{-1} M_{\odot}$ to clusters."
+ Figure 3 shows contours of the abundance of galaxies, Figure \ref{fig:NHalo} shows contours of the abundance of galaxies
+and a flow that appears to clissipate rather close to the core.,and a flow that appears to dissipate rather close to the core.
+ Jets associated with QSOs mav (vpically be followed further than those associated with BL Lac objects. so perhaps one facet of the class distinction relates to the amplitude of precessional perturbation {ο which the jet is subjected.," Jets associated with QSOs may typically be followed further than those associated with BL Lac objects, so perhaps one facet of the class distinction relates to the amplitude of precessional perturbation to which the jet is subjected."
+ Evidently a small helical perturbation (transverse velocity of order of longitudinal velocity) does lead to disruption of the flow. 2001).., Evidently a small helical perturbation (transverse velocity of order of longitudinal velocity) does lead to disruption of the flow \citep{hhr}.
+ Our inunecliale goal is (o generate a well-saaipled erid of simulations for a range of Lorentz [actor ancl precession cone semi-angle. to establish both the largest amplitude of precession lor which flows of a given Lorentz [actor avoid macroscopic break-up. and the physical origin of that break up.," Our immediate goal is to generate a well-sampled grid of simulations for a range of Lorentz factor and precession cone semi-angle, to establish both the largest amplitude of precession for which flows of a given Lorentz factor avoid macroscopic break-up, and the physical origin of that break up."
+ The extent to which propagating shocked fIuid. thought to explain moving VLBI components. will follow the helical channel. or follow approximately ballistie trajectories. is also a subject for future study.," The extent to which propagating shocked fluid, thought to explain moving VLBI components, will follow the helical channel, or follow approximately ballistic trajectories, is also a subject for future study."
+ This work was supported in part by NSF grants AST 9617032 and PILY 9979985. by NRAC allocation MCA938025 al NCSA. and by the Ohio Supercomputer Center.," This work was supported in part by NSF grants AST 9617032 and PHY 9979985, by NRAC allocation MCA93S025 at NCSA, and by the Ohio Supercomputer Center."
+ We thank the releree [ον many useful comments that helped to improve the presentation of this work., We thank the referee for many useful comments that helped to improve the presentation of this work.
+distance of Ro=SA kpc). respectively.,"distance of $R=8.4$ kpc), respectively."
+ The dense spots present in the first column of the panels are due to recent SN explosions (sec. for example. the top-left panel).," The dense spots present in the first column of the panels are due to recent SN explosions (see, for example, the top-left panel)."
+ Instead. the more structured distribution visible in the middle column is due to ejecta trapped and condensed. within clouds forming at [larger z from the lifted disk eas which sullers thermal instability and. cools down to Z=107? Kk (the minimum temperature allowed in our simulations).," Instead, the more structured distribution visible in the middle column is due to ejecta trapped and condensed within clouds forming at larger $z$ from the lifted disk gas which suffers thermal instability and cools down to $T=10^4$ K (the minimum temperature allowed in our simulations)."
+ These clouds are characterized by large column densities. and are visible in lig.," These clouds are characterized by large column densities, and are visible in Fig."
+ 3 as red spots., \ref{fig:colden} as red spots.
+ The cold. dense gas must eventually [all back onto the disk.," The cold, dense gas must eventually fall back onto the disk."
+ Figure 4. shows the column density of the clouds with negative vertical velocity (the cloud. gas is selected as gas with 7x5Of IX).," Figure \ref{fig:vneg} shows the column density of the clouds with negative vertical velocity (the cloud gas is selected as gas with $T\leq
+5\times 10^4$ K)."
+ This figure clearly demonstrates how the gas liftec by the fountains falls back remaining mostly within the ring of the disk in which the active area is locatec., This figure clearly demonstrates how the gas lifted by the fountains falls back remaining mostly within the ring of the disk in which the active area is located.
+ This result confirms what was found in Paper Land indicates, This result confirms what was found in Paper I and indicates
+We also iuspected the Canadian Galactic Plane Survey (COPS. Tavloretal. 2003)) at 1.1 GIIz and the Cree Bank [55 GIIz northern sky survey (CGB6. Coudon et al.,"We also inspected the Canadian Galactic Plane Survey (CGPS, \cite{tay03}) ) at 1.4 GHz and the Green Bank 4.85 GHz northern sky survey (GB6, Condon et al."
+ 1991. 1993. 1991).," 1991, 1993, 1994)."
+ Some traces of the extended feature quoted im previous sections are also visible around Cveuus N-3 in both surveys., Some traces of the extended feature quoted in previous sections are also visible around Cygnus X-3 in both surveys.
+ Therefore. this adds even iore," Therefore, this adds even more confidence to its reality."
+ of , In the left panel of Fig.
+"display the COPS πμασο,", \ref{surveys} we display the CGPS image.
+ These traces of extended enussion are better seen here than in the CDB6 single dish with ⇁⋅lower augular↜⋅⋅↴ resolution., These traces of extended emission are better seen here than in the GB6 single dish survey with lower angular resolution.
+ ⋅↸∖∐∏↴∖∷∖↴↕∪∐∐∐∖↸⊳∐⋜∐∐↴∖↴⊔↴∖↴∙↕↕⋜⋯↖↽↸⊳⋜↕↴∖↴↸∖∙⊓⋅⋯∖↕⊔⋜↧↑↸⊳↕∐∐≼↴⋁↴⋈∖⋜∐⊔ The presence of this extended emission around Ceuus N-3 iu both surveys encouraged us to attempt a crude estimate of its spectral iudex., The presence of this extended emission around Cygnus X-3 in both surveys encouraged us to attempt a crude estimate of its spectral index.
+ Our VLA (6 cin) aud Ryle Telescope (2 cui) nuages are likely missive some flux density from very short baselines aud we preferred not to use them for such purpose., Our VLA (6 cm) and Ryle Telescope (2 cm) images are likely missing some flux density from very short baselines and we preferred not to use them for such purpose.
+ Since we were ouly interested in exteuded cussion. the ceutral core of Cveuus N-9 was initially subtracted from both the CGPS and GB6 images.," Since we were only interested in extended emission, the central core of Cygnus X-3 was initially subtracted from both the CGPS and GB6 images."
+ The CGPS image was later appropriately couvolved with an elliptical Gaussian so that its one arc-lminute angular resolution matched the lower one of the (BG (about 3are-iiuute)., The CGPS image was later appropriately convolved with an elliptical Gaussian so that its one arc-minute angular resolution matched the lower one of the GB6 (about 3 arc-minute).
+ In the process. we also transformed its brightuess seale from brightucss temperature to Jv |.," In the process, we also transformed its brightness scale from brightness temperature to Jy $^{-1}$."
+ The COPS image was also reerided from galactic to equatorial coordinates., The CGPS image was also regrided from galactic to equatorial coordinates.
+" Ouce the COPS and CBG images were correctly matched. we combined them in order to create a imap of spectral index a (5,x Vv"")."," Once the CGPS and GB6 images were correctly matched, we combined them in order to create a map of spectral index $\alpha$ $S_{\nu} \propto \nu^{\alpha}$ )."
+ The result is reported iu the right panel of Fig. 5.., The result is reported in the right panel of Fig. \ref{surveys}.
+ Here. the spectral index of the extended radio feature under discussion hints at a negative value az0.5 and sugecstingOO a non-thermal ciission mechanism.," Here, the spectral index of the extended radio feature under discussion hints at a negative value $\alpha \simeq -0.5$ and suggesting a non-thermal emission mechanism."
+ This would be iu natural agreement with relativistic electrons accelerated in the teruination shock a distorted . jet., This would be in natural agreement with relativistic electrons accelerated in the termination shock of a distorted jet.
+In. contrast. other bright exteuded∙ radio we∉⋡↭∐∱↕⋠∐∖∐↫↾⋠⋝↕↾≭↨⋅⋠∖⋮↕↕↕↾⋮∖↽∙↕∐↾∐⋠∖↕⋠��↕⊳↾↻⋮↴⋯∖↕↭↕⊳∏⋮↴∙≭∏ sources in the field have à values predominantly positive as expected from ordinary star formune regious with thermal survev," In contrast, other bright extended radio sources in the field have $\alpha$ values predominantly positive as expected from ordinary star forming regions with thermal emission mechanisms."
+ ∪↴⋝↴∖↴↸∖↥⋅↖⇁⋜↧↑↕∪∐↴∖↴⋜↧↑≼∐⇡∶↸∖↥⋅↸∖∐↖↖⇁⋜↧↖⇁↸∖↕↸∖∐∶↴⋁↑∐∖↴⋜⋯∖↴∖↴↑↕∐∐↸∖↸∖≼∐∖≼↧↑∪ confidently constrain the spectral iudex value.," In any case, true matching beam observations at different wavelengths are still needed to confidently constrain the spectral index value."
+ The possible Cyeuus. N-5 lobe or extended radio feature of other kind las an integrated fiux density of about 133 iJy at the 6 emi waveleneth aud extends over a solid anele of about 6.8&10. sr., The possible Cygnus X-3 lobe or extended radio feature of other kind has an integrated flux density of about 133 mJy at the 6 cm wavelength and extends over a solid angle of about $6.8\times 10^{-7}$ sr.
+ These values have been measured using the full iuulti-coufiguration radio map in Fig. 2.., These values have been measured using the full multi-configuration radio map in Fig. \ref{cd+d+b}.
+ Fora=0.5. this implies a radio liinosity of 6.8\]y? ore | in the 01-100 GIIz at 9 kpc.," For $\alpha=-0.5$, this implies a radio luminosity of $6.8 \times 10^{32}$ erg $^{-1}$ in the 0.1-100 GHz at 9 kpc."
+" Asstuning standard equipartition conditions. the minimum cucrey content is estimated as 3.2.10"" ere with the equipartition maguetic field being5 2-«10.? Ci"," Assuming standard equipartition conditions, the minimum energy content is estimated as $3.2 \times 10^{47}$ erg with the equipartition magnetic field being $2\times 10^{-5}$ G."
+Fic. 1.,. .—
+" Evolution of he kinetic temperature of the eas line), the spin temperature of neutral hydrogen dine). aud the CAIB temperature linc) in the simulation.Fic. "," Evolution of the kinetic temperature of the gas ), the spin temperature of neutral hydrogen ), and the CMB temperature ) in the simulation. .—"
+2. The upper panel slows the transiuitted radio flux density over a relatively narrow velocity range (FOO kin 1j assundue III 21«1à absorption bv the neutral ICAL, The upper panel shows the transmitted radio flux density over a relatively narrow velocity range (700 km $^{-1}$ ) assuming HI 21cm absorption by the neutral IGM.
+ Three different redshifts are displaved: +=12 (red). :=10 (violet). and :=8 (blue).," Three different redshifts are displayed: $z=12$ (red), $z=10$ (violet), and $z=8$ (blue)."
+ The abscissa for the upper panel is velocity. while that or the middle and lower panels is the corresponding comoving plysical scale.," The abscissa for the upper panel is velocity, while that for the middle and lower panels is the corresponding comoving physical scale."
+ The middle panel shows the kinetic dines) and the spin lines) temperature of the the ucutral ICAL over the range of distances that contribute to the velocity range nmcdicate ou the top panel., The middle panel shows the kinetic ) and the spin ) temperature of the the neutral IGM over the range of distances that contribute to the velocity range indicated on the top panel.
+ The bottom panel shows the neutral hydrogen density.Fic. , The bottom panel shows the neutral hydrogen density. .—
+3. The joint distribution of the spin temperature. Ta. and neutral hydrogen density. xyi(l | 0) (in units of the mean hivdrogeu deusitv] of gas aloug random lines of sight at two redshifts: (6) ;=12 (upper panel) iid (0) 2—8 (lower panel).," The joint distribution of the spin temperature, $_{\rm S}$ , and neutral hydrogen density, $_{\rm HI}$ (1 + $\delta$ ) (in units of the mean hydrogen density) of gas along random lines of sight at two redshifts: ) $z=12$ (upper panel) and ) $z=8$ (lower panel)."
+ The solid curves inclicate constant III 2lem optical depth. τ. as labeled.Fic. ," The solid curves indicate constant HI 21cm optical depth, $\tau$, as labeled. .—"
+4. The radio spectrun of the powerful radio ealaxy Cyvguus A at ;—0.057 (Ps)=1.1ς1079 ore 1TIIz !: Baars et al., The radio spectrum of the powerful radio galaxy Cygnus A at $z = 0.057$ $P_{151} = 1.1 \times 10^{36}$ erg $^{-1}$ $^{-1}$; Baars et al.
+ 1977)., 1977).
+ The dash line is a first order polvnonual fit to the (log) data. corresponding o a power-law of index —1.05+0.03.," The dash line is a first order polynomial fit to the (log) data, corresponding to a power-law of index $-1.05\pm 0.03$."
+ The solid line is a second order polynomial fit. Fic. δα, The solid line is a second order polynomial fit. . .—
+. The simulated spectrum from 100 MITz o 200 MIIZzZ of a source with σου = 20 iuJv at >=10 usine the Cyeuus A spectral model and assuniue TT 210i absorption by the ICAL, The simulated spectrum from 100 MHz to 200 MHz of a source with $_{120}$ = 20 mJy at $z = 10$ using the Cygnus A spectral model and assuming HI 21cm absorption by the IGM.
+ Thermal roise las been added using the specifications of the SISA and asstuning LO davs integration with 1 kIIz wide spectral channels. ο., Thermal noise has been added using the specifications of the SKA and assuming 10 days integration with 1 kHz wide spectral channels. .—
+ The same as 5a. but showing au expanded view of the spectral region around the frequency corresponding to the redshift ΤΗ 2tei line at the source redshift (129 MIIz).," The same as 5a, but showing an expanded view of the spectral region around the frequency corresponding to the redshift HI 21cm line at the source redshift (129 MHz)."
+ The solid line is the Cyenus A model spectrum without nolse or absorption.Fic. , The solid line is the Cygnus A model spectrum without noise or absorption. .—
+θα. The simulated spectrum from 100 MITz to 200 AMTz of a source with S494 2 35 11Jv at 2=8 using the Cryenus A spectral model aud assuuiusg III 2lei absorption by the IGAL, The simulated spectrum from 100 MHz to 200 MHz of a source with $_{120}$ = 35 mJy at $z = 8$ using the Cygnus A spectral model and assuming HI 21cm absorption by the IGM.
+ Thermal noise has been added using the specifications of the SILA aud assuniue 10 davs integration with 1 kITz wide spectral channels. 6b, Thermal noise has been added using the specifications of the SKA and assuming 10 days integration with 1 kHz wide spectral channels. .—
+. The same as Ga. but showing au expanded view of the spectral region around the frequency correspouding to the redshift IIT 21« line at the source redshift (157.8 ΑΠ).," The same as 6a, but showing an expanded view of the spectral region around the frequency corresponding to the redshift HI 21cm line at the source redshift (157.8 MHz)."
+ The solid line is the Cyenns A model spectrum without noise or absorption.Fic. , The solid line is the Cygnus A model spectrum without noise or absorption. .—
+7 The xiis noise of the spectrmm shown imn Figure la (2=8 source) measured in bands of 1 MITZz width with 1 kIIz channels.Fic. , The rms noise of the spectrum shown in Figure 4a $z = 8$ source) measured in bands of 1 MHz width with 1 kHz channels. .—
+8. The cunulative distribution of channel values for a LO MITz wide band using 1 kITz chamnels for the spectrum iu Fieure la (2=8 source)., The cumulative distribution of channel values for a 10 MHz wide band using 1 kHz channels for the spectrum in Figure 4a $z = 8$ source).
+ The solid line shows the noise distribution at frequencies just below 157.8 Az. je.," The solid line shows the noise distribution at frequencies just below 157.8 MHz, ie."
+" just before the on-set of III 21cm absorption by the neutral IGAL while the dotted line shows the distribution just above 157.5 ΑΠΕ, ie."," just before the on-set of HI 21cm absorption by the neutral IGM, while the dotted line shows the distribution just above 157.8 MHz, ie."
+ just after the on-set of ITI οτι absorption by the IGALFic. 9., just after the on-set of HI 21cm absorption by the IGM. .—
+ The predicted παο: of radio sources with PiiνL3«103 ere 1 + over the entire sky per unit redshift vs. redshift usine the two modelsdescribed. in section lL., The predicted number of radio sources with $P_{151} \ge 4.3\times 10^{34}$ erg $^{-1}$ $^{-1}$ over the entire sky per unit redshift vs. redshift using the two modelsdescribed in section 4.
+ The normalization is set at 2=[by Jarvis et al. (, The normalization is set at $z = 4$by Jarvis et al. (
+2001) extrapolated to lower DIuniuositv.,2001) extrapolated to lower luminosity.
+niaximises the fibre multiplex since there is no need for eross-beanmeswitching.,maximises the fibre multiplex since there is no need for cross-beam-switching.
+ An alternative. assessment of the data quality is presented. in Fig. 12.., An alternative assessment of the data quality is presented in Fig. \ref{stacked spectra}.
+ Here we create stacked spectra from the 219 WigeleZ observations from 2009., Here we create stacked spectra from the 219 WiggleZ observations from 2009.
+ ln each case the spectrum is presented with and without the PCA sky subtraction., In each case the spectrum is presented with and without the PCA sky subtraction.
+ Insullicient N|S observations are available to generate the comparison spectrum at this timo., Insufficient N+S observations are available to generate the comparison spectrum at this time.
+ The suppression of svstematic residuals in the spectra. is Clearly visible in the PCA processed. data., The suppression of systematic residuals in the spectra is clearly visible in the PCA processed data.
+ No continuum subtraction was applied to the fibres in Fig., No continuum subtraction was applied to the fibres in Fig.
+ 12. and. so the weak source continuum from the high equivalent. width emission line targets of the WigeleZ survey is visible in the averaged spectrum., \ref{stacked spectra} and so the weak source continuum from the high equivalent width emission line targets of the WiggleZ survey is visible in the averaged spectrum.
+ The most apparent benefit of the PCA analysis to the user. for. faint-object spectroscopy. d8 the reduction in systematic cosmetic defects.," The most apparent benefit of the PCA analysis to the user, for faint-object spectroscopy, is the reduction in systematic cosmetic defects."
+ In. order to. provide a quantitative assessment of the improvement a detailed simulation has been undertaken., In order to provide a quantitative assessment of the improvement a detailed simulation has been undertaken.
+ AXO0moega survey data was selected from six clear dark nights from the September 2008 observing campaign for the WigeleZ project., AAOmega survey data was selected from six clear dark nights from the September 2008 observing campaign for the WiggleZ project.
+ This provided science. frames., This provided $\times$ sec science frames.
+ As discussed. previously. WiegeleZ ssec.targets are primarily high. equivalent —width faint (rCXD)—22) emission-line sources hence cach of the ~354 science fibres larechy represents blank sky observation.," As discussed previously, WiggleZ targets are primarily high equivalent width faint $r$ $\sim$ 22) emission-line sources hence each of the $\sim$ 354 science fibres largely represents blank sky observation."
+ Additionally. each cate frame contains 25 fibre allocated to oeselected blank sky. positions which are free from the low evel Poisson noise contribution from the faint continuum still present in fibres allocated to WigeleZ science targets.," Additionally, each date frame contains 25 fibre allocated to preselected blank sky positions which are free from the low level Poisson noise contribution from the faint continuum still present in fibres allocated to WiggleZ science targets."
+ 1n what follows the blank sky fibre provide the idealised est data set while the science fibres observations are used ο provide a larger sample size., In what follows the blank sky fibre provide the idealised test data set while the science fibres observations are used to provide a larger sample size.
+ Each CCD frame was processed using the software environment ancl the standard. cata reduction recipes., Each CCD frame was processed using the software environment and the standard data reduction recipes.
+ For each individual reduced. data frames ssec integration) a svnthetic spectrum was generated. from a template carly type galaxy., For each individual reduced data frames sec integration) a synthetic spectrum was generated from a template early type galaxy.
+ Phe galaxy template used was the intermediate age carb type spectrum from the GDDS survey (Abrahametal.2004)., The galaxy template used was the intermediate age early type spectrum from the GDDS survey \citep{gdds}.
+. Phe galaxies where distributed: in redshift between O.7«z« 1.3. whieh is the imit imposed by the rest frame spectral data matchec o the observed. spectral. range., The galaxies where distributed in redshift between $<$ $<$ 1.3 which is the limit imposed by the rest frame spectral data matched to the observed spectral range.
+ Ελπίς redshift range see he Calcium IL]Ix. absorption lines and the Balmer break ransit the strongest OIL sky-line regions in the re AAOmecga spectrum., This redshift range see the Calcium H+K absorption lines and the Balmer break transit the strongest OH sky-line regions in the red AAOmega spectrum.
+ Spectra where generated for a series of source intensities scaled to match count rates equivalen o apparent /(Vega)-20-24. with the appropriate targe 'oisson noise included.," Spectra where generated for a series of source intensities scaled to match count rates equivalent to apparent $I$ (Vega)=20-24, with the appropriate target Poisson noise included."
+ The PCA sky subtraction was then applied to each reduced. seience frame and image stacks where created a each simulated. magnitude level with and without the PCA correction., The PCA sky subtraction was then applied to each reduced science frame and image stacks where created at each simulated magnitude level with and without the PCA correction.
+ A redshift analysis was undertaken. with the fully automated mode of the cross-correlation code typically used with cata from the δα/XXOmoega system., A redshift analysis was undertaken with the fully automated mode of the cross-correlation code typically used with data from the 2dF/AAOmega system.
+ The fully automated: analysis was chosen in preference to the more typical semi-interactive methods more common to faint object programs undertaken with AAOmega citet WakeQ6.1toss06)) to avoid. subjective results.," The fully automated analysis was chosen in preference to the more typical semi-interactive methods more common to faint object programs undertaken with AAOmega \\citet{Wake06,Ross06}) ) to avoid subjective results."
+ For this idealised test case. only a single cross correlation template was used. the same spectrum used in the generation of the mock data.," For this idealised test case, only a single cross correlation template was used, the same spectrum used in the generation of the mock data."
+ Phe prior for the redshift distribution was match o that of the mock test data input range., The prior for the redshift distribution was match to that of the mock test data input range.
+ The redshift analysis results are shown in Fig. 13.., The redshift analysis results are shown in Fig. \ref{mockdata}.
+ The upper panel indicates the measured. redshift returned: by as a function of input mocel redshift for the £(Vega)=2 1 simulated data set for the full bhour (85 frame)., The upper panel indicates the measured redshift returned by as a function of input model redshift for the $I$ (Vega)=21 simulated data set for the full hour (85 frame).
+ With he PCA processed data we recover the correct redshift or of simulated: sources., With the PCA processed data we recover the correct redshift for of simulated sources.
+ The. misidentified: spectra are preferentially at. the highest. redshifts. where the /- xuxd normalisation results in the lowest signal-to-noise ratio spectra.," The misidentified spectra are preferentially at the highest redshifts, where the $I$ -band normalisation results in the lowest signal-to-noise ratio spectra."
+ In the absence of the PCA sky correction. the success rate is reduced to with a number of distinct zilure modes clearly present.," In the absence of the PCA sky correction, the success rate is reduced to with a number of distinct failure modes clearly present."
+ A comparison of the model spectra. derived. from Xank sky libre spectra rather than those containing WigeleZ science targets demonstrates that the residual noise contribution from the faint WigeleZ targets present in the science spectra used for this analvsis does contribute a reduced. ellicicney for the test and will therefore contribute to the slope measured for the sky error reduction rate seen in Fig. 9..," A comparison of the model spectra derived from blank sky fibre spectra rather than those containing WiggleZ science targets demonstrates that the residual noise contribution from the faint WiggleZ targets present in the science spectra used for this analysis does contribute a reduced efficiency for the test and will therefore contribute to the slope measured for the sky error reduction rate seen in Fig. \ref{PCA
+stack}."
+ The lower panel of Fig., The lower panel of Fig.
+ 13. shows the declining recovery rate for fainter model spectra., \ref{mockdata} shows the declining recovery rate for fainter model spectra.
+ At fainter input magnitudes. the fully automated code. which was not specifically tuned for such a regime. does remarkably well even at 4223.," At fainter input magnitudes, the fully automated code, which was not specifically tuned for such a regime, does remarkably well even at $I$ =23."
+" With the planed upgrade of the red arm of AXOmoega to a new generationhigh-resisHivilig CCD (boosting quantum ellicienev in the range SOOnm-lym by a factor of four) it will be possible for AXOmega to continue to conduct ellicent wide field surveys for early tvpe galaxies (Cannon.ter.Eisenstiene£a£,2006). out to at redshift z~ 1.", With the planed upgrade of the red arm of AAOmega to a new generation CCD (boosting quantum efficiency in the range $\mu$ m by a factor of four) it will be possible for AAOmega to continue to conduct efficent wide field surveys for early type galaxies \citep{2SLAQ} out to at redshift $\sim$ 1.
+ We have demonstrated. that the use of a Principle Components Analysis sky subtraction with dedicated: sky fibres observations in a manner similar to that. prescribed bv Wurtz&Mink(2000). ancl Wild&Lewett(2005) (and: subsequently implemented. for the GAALA survey by Parkinsonοἱαἱ.(in prep.))). provides excellent. quality astronomical spectra over long duration exposures.," We have demonstrated that the use of a Principle Components Analysis sky subtraction with dedicated sky fibres observations in a manner similar to that prescribed by \citet{Kurtz00} and \citet{Wild05} (and subsequently implemented for the GAMA survey by \citet{Parkinson}) ), provides excellent quality astronomical spectra over long duration exposures."
+ Lor moderate length exposures. hhours on source) the spectral quality is comparable to that obtained. using the NUS observational technique. and likely for significantly longer exposures.," For moderate length exposures hours on source) the spectral quality is comparable to that obtained using the N+S observational technique, and likely for significantly longer exposures."
+ We therefore conclude that for fibre based multi-object spectroscopy. in instances were high multiplex ancl observing efficiency. are critical. the PCA approach is superior to nod-andeshullle based techniques for moderate signal-to-noise extragalactic survey spectroscopy.," We therefore conclude that for fibre based multi-object spectroscopy, in instances were high multiplex and observing efficiency are critical, the PCA approach is superior to nod-and-shuffle based techniques for moderate signal-to-noise extragalactic survey spectroscopy."
+ We wish to. thank Scott Croom. Will Saunders. and ναι). Glazebrook for invaluable cliscussion regarding the concepts discussed throughout this work.," We wish to thank Scott Croom, Will Saunders and Karl Glazebrook for invaluable discussion regarding the concepts discussed throughout this work."
+ Εις investigation would not have been possible without the hare work and dedication of the stall at the Anelo-Australia Observatory.," This investigation would not have been possible without the hard work and dedication of the staff at the Anglo-Australia Observatory,"
+Model D with thin thermal emission (apec) is adopte as the best model for IC 5063 from the time average spectra.,Model B with thin thermal emission ) is adopted as the best model for IC 5063 from the time averaged spectra.
+ The thermal enissiou is required because it euificantlv reduces (7 in all 3 models (A. D. iux C) and its addition produces acceptable fits.," The thermal emission is required because it significantly reduces $\chi^2$ in all 3 models (A, B, and C) and its addition produces acceptable fits."
+ Models A. D. aud C eive 4?/dof = Ish.s/177. 175/175 an 175/173. respectively.," Models A, B, and C give $\chi^2/{\rm dof}$ = 185.8/177, 175/175 and 175/173, respectively."
+ Thus. the improvement of the fit by adding a reflection component (Model DB) is fou to be significant with an F-test probability of 0.007. whereas the introduction of an additional absorber for the transmitted component (Model C) is not significant iu introducing 47.," Thus, the improvement of the fit by adding a reflection component (Model B) is found to be significant with an F-test probability of 0.007, whereas the introduction of an additional absorber for the transmitted component (Model C) is not significant in introducing $\chi^2$."
+ The results of simultaucous analysis of the aud BAT spectra with Model B | are sunuuarized in Table 3.. and the best-fit modelis plotted in Fieures 2. and 3..," The results of simultaneous analysis of the and BAT spectra with Model B + are summarized in Table \ref{para}, and the best-fit modelis plotted in Figures \ref{data} and \ref{model}."
+ The best fit photou iudex of Dz1(6 from IC 5063 is vpical of radio galaxies., The best fit photon index of $\Gamma \approx 1.7$ from IC 5063 is typical of radio galaxies.
+ The scattered fraction faa404 is larger than that of 3€ 103.," The scattered fraction $f_{\rm
+scat} \sim 0.9\%$ is larger than that of 3C 403."
+ I is consistent with he result of ~1% (Vignalietal.1997)., It is consistent with the result of $\sim$ \citep{Vig97}.
+.. The marrow dmon-E line detected at 6.1 keV is consisteut with being emitted from the torus., The narrow iron-K line detected at 6.4 keV is consistent with being emitted from the torus.
+ Its equivalent width with respect to the reflected contiuuuu (1.5+0.2 τον) ds within a ranec of theoretical expectation. thus supporting the validity of our spectral model.," Its equivalent width with respect to the reflected continuum $1.5\pm0.2$ keV) is within a range of theoretical expectation, thus supporting the validity of our spectral model."
+ Similar o the case of 3€ 1023. the reflection. component is less absorbed iqu~--3107 73) compared with the ransnütted component (NYS23«107 723). and disfavoring the likely edge-ou geometry.," Similar to the case of 3C 403, the reflection component is less absorbed $N_{\rm H}^{\rm refl} \approx 3\times 10^{22}$ ) compared with the transmitted component $\approx 3\times 10^{23}$ ), and disfavoring the likely edge-on geometry."
+ The reflection strength of R~--0.6 relative to the averaged DAT Hux is close to the value obtained from 3€ 103.," The reflection strength of $R \approx
+0.6$ relative to the averaged BAT flux is close to the value obtained from 3C 403."
+ These results sugeest that the torus ecometry of tle two ACNs could be similay to cach other., These results suggest that the torus geometry of the two AGNs could be similar to each other.
+ The obtained eniperature aud Iwunuinositv of the componcut are in the typical range of hot ISM detected frou elliptical ealaxies (Matsushitaetal.2000). and hence supports lis interpretation of its origin. although we cannot rule out the possibility that a part of this cluission may come rom pliotoioinized plasma irradiated by the ACN. which should be regarded as an accompanying component of he scattered emission.," The obtained temperature and luminosity of the component are in the typical range of hot ISM detected from elliptical galaxies \citep{Mat00} and hence supports this interpretation of its origin, although we cannot rule out the possibility that a part of this emission may come from photoioinized plasma irradiated by the AGN, which should be regarded as an accompanying component of the scattered emission."
+ Nevertheless. since the amplitude of the scattered fraction is αλα] determined from the spectrum in d2 keV band. where no strong emission mes are present. its iufensitv is not stronely affected by the nucertaimty in the nature of the line cutting component.," Nevertheless, since the amplitude of the scattered fraction is mainly determined from the spectrum in 1–2 keV band, where no strong emission lines are present, its intensity is not strongly affected by the uncertainty in the nature of the line emitting component."
+ Since IC 5065 exhibits significant time variability during the observation. we check time variability of the spectra bv dividing the time region into two states with differcut fiux levels. epoch I (for 57.6 ks frou the beeimuing) aud epoch IT (after that).," Since IC 5063 exhibits significant time variability during the observation, we check time variability of the spectra by dividing the time region into two states with different flux levels, epoch I (for 57.6 ks from the beginning) and epoch II (after that)."
+ We adopt the Modelως D. | model for cach spectrum., We adopt the Model B + model for each spectrum.
+ Asstmine, Assuming
+Ny=10708? towards thenucleus. (Warwicketal.1989).,"$N_{\rm H}=10^{23.1}\,{\rm cm}^{-2}$ towards thenucleus \cite{WarwickKoyamaInoue1989}."
+. All this suggests the presence of an obscuring torus in Mrk 318., All this suggests the presence of an obscuring torus in Mrk 348.
+ Attempts to detect the obscuring material through radio spectroscopy have failed so far (e.g. II 1. Callimoreοal. 19993).," Attempts to detect the obscuring material through radio spectroscopy have failed so far (e.g. H I, \citeNP{GallimoreBaumO'Dea1999}) )."
+ What makes this galaxw stand out among Sevtert ealaxies is its bright and variable radio nucleus., What makes this galaxy stand out among Seyfert galaxies is its bright and variable radio nucleus.
+" Neff&deBruvu(1983) oui a compact radio core on. VLBI scales with a flux deusitv of several hundred milli--Jausky and a flat to inverted spectrum,", \citeN{NeffdeBruyn1983} found a compact radio core on VLBI scales with a flux density of several hundred milli-Jansky and a flat to inverted spectrum.
+ Ulvestadetal.(1999) preseuted more receut. VEDI observatious for two epoclis. showing a two-coniponent structure expanding with sub-relativistic speeds.," \citeN{UlvestadWrobelRoy1999} presented more recent VLBI observations for two epochs, showing a two-component structure expanding with sub-relativistic speeds."
+ They also note a fare of the radio continui enission at 15 Clg with the fux risus from 120 11Jv to 570 10Jx. between 1997.10 and 1998.75., They also noted a flare of the radio continuum emission at 15 GHz with the flux rising from 120 mJy to 570 mJy between 1997.10 and 1998.75.
+ In the following we will present aud discuss results of IK-band radio spectroscopy of this galaxy., In the following we will present and discuss results of K-band radio spectroscopy of this galaxy.
+ Data were taken in March and April 2000. uxiug the Effelsbere 100 i telescope of the AIPIER equipped with a dual channel K-baud IIEMT receiver.," Data were taken in March and April 2000, using the Effelsberg 100 m telescope of the MPIfR equipped with a dual channel K-band HEMT receiver."
+" The svsteii temperature was of order TORS on a T4 temperature scale: the beam size was 10""."," The system temperature was of order K on a $T_{\rm
+A}^{*}$ temperature scale; the beam size was $''$."
+ The data were recorded using an autocorrelator with & « 256 channels aud bandwidths of MNIIz each., The data were recorded using an autocorrelator with 8 $\times$ 256 channels and bandwidths of MHz each.
+ The eight backends were configured im two groups of four. sampling the two orthogonal linear polarizations.," The eight backends were configured in two groups of four, sampling the two orthogonal linear polarizations."
+ Frequency shifts between the four backeuds representius a given linear polarization were adjusted m such a sway that a total velocity rauge of | could be covered., Frequency shifts between the four backends representing a given linear polarization were adjusted in such a way that a total velocity range of $^{-1}$ could be covered.
+" The measureiieuts were carried out iu a dual beam switching mode (switching frequency Wz) with a beam throw of 121"" iu azimuth.", The measurements were carried out in a dual beam switching mode (switching frequency Hz) with a beam throw of $''$ in azimuth.
+ Oulv linear baselines were subtracted., Only linear baselines were subtracted.
+ Calibration was obtained by measurements of W3OII JJy according to Mauersberger.Wilson.&Henkel 1988))., Calibration was obtained by measurements of W3OH Jy according to \citeNP{MauersbergerWilsonHenkel1988}) ).
+" Poiting could be checked on Mrk 3Ls itself and was found to be stable to within 7""", Pointing could be checked on Mrk 348 itself and was found to be stable to within $''$.
+ The galaxy was initially observed to look for auinonuia CNIT;) and exclopropeuxlideue (C302) absorption against the bright radio uucleus., The galaxy was initially observed to look for ammonia $_3$ ) and cyclopropenylidene $_3$ $_2$ ) absorption against the bright radio nucleus.
+ No absorption features were found., No absorption features were found.
+ The transitions. rest frequeucies. aud upper limits of these observations are isted in Table 1..," The transitions, rest frequencies, and upper limits of these observations are listed in Table \ref{abslines}."
+ Since the II2O. (6465 5595) trausition is very close we also observed the redshifted 22.253508 CIIz line aud detected an enmuüssiou. feature ou March. 17., Since the $_2$ O $6_{16}-5_{23}$ ) transition is very close we also observed the redshifted 22.23508 GHz line and detected an emission feature on March 17.
+ We repeated he observations on five subsequent days aud detected the ine cach tine., We repeated the observations on five subsequent days and detected the line each time.
+ The six spectra are shown in Fig. 1.., The six spectra are shown in Fig. \ref{h2ospec}.
+ We also produced a combined spectruu of all days (Fie. 2..," We also produced a combined spectrum of all days (Fig. \ref{specall},"
+ op) and fitted the IT2O line with a Gaussian profile., top) and fitted the $_2$ O line with a Gaussian profile.
+ A one conrponeut fit vields a ceutral velocity of e=1611.641.2au [for the Hue. which is redshifted by 133 kin 1 roni the systemic velocity.," A one component fit yields a central velocity of $v=4641.6\pm1.2$km $^{-1}$ for the line, which is redshifted by 133 km $^{-1}$ from the systemic velocity."
+" The peak flux is 5,=31 12Jv with a full width at half masxinuuun of Ae=13043 kin ?", The peak flux is $S_{\nu}=34$ mJy with a full width at half maximum of $\Delta v=130\pm3$ km $^{-1}$.
+ The inteerated fluxis L7LE0.1 Iv lans.+. yielding an apparent isotropic luminosity of 120 Εν," The integrated fluxis $\pm0.1$ Jy km$^{-1}$, yielding an apparent isotropic luminosity of 420 $L_\odot$."
+" A two componcut fit vields a narrow liue at e=1677.7d(8 kins |! with Av=S8.64£3.2 kins + and $,=27 andy plus a broa ine with e=1609.5€1.9 lan ο, Ac=10L.8+1.9 au land $,=26 uy."," A two component fit yields a narrow line at $v=4677.7\pm0.8$ km $^{-1}$ with $\Delta
+v=58.6\pm3.2$ km $^{-1}$ and $S_{\nu}=27$ mJy plus a broad line with $v=4609.5\pm1.9$ km $^{-1}$, $\Delta v=101.8\pm1.9$ km $^{-1}$, and $S_{\nu}=26$ mJy."
+ The latter fit iudicates that the lue is asviuinetric aud has a xonounced blue wing., The latter fit indicates that the line is asymmetric and has a pronounced blue wing.
+ If we conrpare our two best spectra ον 17 March and 5 April we tentatively find some variability iu the shape of lis blue wine., If we compare our two best spectra from 17 March and 5 April we tentatively find some variability in the shape of this blue wing.
+ We searched for additional high velocity catures aud did rot detect anything down to a limit of 6 uJv (lo. Eins 1j in the rauge 3250 to 6200 kms.  ane down to a lait of 10-15 iiJy in the rauge 1170 to 10150 aus ?.," We searched for additional high velocity features and did not detect anything down to a limit of 6 mJy $\sigma$, 4 km $^{-1}$ ) in the range 3250 to 6200 km $^{-1}$ and down to a limit of 10-15 mJy in the range $-1470$ to 10480 km $^{-1}$."
+ The continmun fux density we find at 22 GIIz is L8dO. Jy corresponding to 1:1022 Watt |.," The continuum flux density we find at 22 GHz is $0.8\pm0.1$ Jy, corresponding to $4\cdot10^{23}$ Watt $^{-1}$."
+" We also reduced some archival data of carlicr observatious of Myk Als taken between October 1997 aud February 1998 aud combined them iuto one data set (Fig. ὃν,"," We also reduced some archival data of earlier observations of Mrk 348 taken between October 1997 and February 1998 and combined them into one data set (Fig. \ref{specall},"
+ bottom)., bottom).
+ Iu this spectzun the broad line is inareinallv detected with a dux density approximately three times lower than in the current observations., In this spectrum the broad line is marginally detected with a flux density approximately three times lower than in the current observations.
+ The sinaller and variable baseline iu these earlier observations made a reliable ideutiSicatioun of this line impossible withoutprior? information., The smaller and variable baseline in these earlier observations made a reliable identification of this line impossible without information.
+ We have clearly discovered a new meganmiaser in Mrk 318., We have clearly discovered a new megamaser in Mrk 348.
+ It» lhnuuinositv is comparable to the enisson frou NGC 3079 which coutaius the second most luninous ΠΟ naser after TNS 2226-181 (I&oekemiocretal.1995)., Its luminosity is comparable to the emission from NGC 3079 which contains the second most luminous $_2$ O maser after TXS 2226-184 \cite{KoekemoerHenkelGreenhill1995}.
+. The iue width is among the broadest found for a meganaser. sinulay to the masers in TNS 2:26-181 aud NGC 1052.," The line width is among the broadest found for a megamaser, similar to the masers in TXS 2226-184 and NGC 1052."
+ Assunidug the enüssiou is associated with material close ο the ceuter this is the first spectroscopic evidence for nolecular gas possibly obscuring the nucleus., Assuming the emission is associated with material close to the center this is the first spectroscopic evidence for molecular gas possibly obscuring the nucleus.
+ Despite the bright radio continua aud the molecular uaser hue we have found no absorption lines from either NIT3 or C3 Uo., Despite the bright radio continuum and the molecular maser line we have found no absorption lines from either $_3$ or $_3$ $_2$ .
+ This is in line with the earlier nou-detection of II absorption (Callimoreetal. 1999).., This is in line with the earlier non-detection of H absorption \cite{GallimoreBaumO'Dea1999}. .
+ Πο is an organic rine molecule which is widespread in the Galaxy and is associated with diffuse eas in the ISAL (Matthews 1985).., $_3$ $_2$ is an organic ring molecule which is widespread in the Galaxy and is associated with diffuse gas in the ISM \cite{MatthewsIrvine1985}. .
+ Tt was also found in absorption against, It was also found in absorption against
+and μας=19.1ip.,"and $I_{*\rm max} =
+19.1\,I_0$."
+" Neutron star structure models predict 7, close to or slightly above {ῃ so (hat we will favor the lower bounds and expect masses in the vicinity of 1.9M..."," Neutron star structure models predict $I_*$ close to or slightly above $I_0$ so that we will favor the lower bounds and expect masses in the vicinity of $1.9\,M_\odot$."
+ In a last step. we use the figure-ol-merit function JF. Eq. (8)).," In a last step, we use the figure-of-merit function $\mathcal{F}$, Eq. \ref{eq:MeritFunction}) ),"
+ to fit the data and also the variation of mass with spin rate according to Eq. (9))., to fit the data and also the variation of mass with spin rate according to Eq. \ref{eq:RelationSpinMasse}) ).
+ We span a vast range in the CA.FL )-plane to compute the merit function.," We span a vast range in the $(M,I_*)$ -plane to compute the merit function."
+ Note however the subtil change in unknowns compared to the previous analvtical study., Note however the subtil change in unknowns compared to the previous analytical study.
+" Now we use a constant mass for thefamal, AL. imstead of a constant mass for the rotating star. M,=M+££Q2/2(2,"," Now we use a constant mass for the, $M$, instead of a constant mass for the rotating star, $M_* = M + I_*\,\Omega_*^2/2\,c^2$."
+ Actually. the discrepancy. between both approaches is small ancl can be neglected at the end of the study.," Actually, the discrepancy between both approaches is small and can be neglected at the end of the study."
+ The results presenting the isocontours of the merit fanction are summarized in Fig. 2.., The results presenting the isocontours of the merit function are summarized in Fig. \ref{fig:MeritFunction1}.
+" The resulting region for minimization of F is marked in red ancl shapes very similarly {ο the previous diagram (M. Z,). Fie. 1.."," The resulting region for minimization of $\mathcal{F}$ is marked in red and shapes very similarly to the previous diagram $(M_*,I_*)$ , Fig. \ref{fig:ContrainteMI}."
+" The most probable mass and moment of inertia are M2:2:02.2 andl,220.5—1.5(10.km)?MU..."," The most probable mass and moment of inertia are $M \approx 2.0-2.2\,M_\odot$ and $I_* \approx 0.5-1.5 \,
+(10\;\textrm{ km})^2 \, M_\odot$."
+ The best fit according to these values is shown in Fie., The best fit according to these values is shown in Fig.
+ 3. where (he spin rate is plotted on (hie x-axis aud the (win kIIz-QDOs difference normalized to the spin rate is plotted on the y-axis., \ref{fig:Rotateur1} where the spin rate is plotted on the x-axis and the twin kHz-QPOs difference normalized to the spin rate is plotted on the y-axis.
+ First. we retrieve the segregalion between slow and fast rotators at (he correct frequency as expected.," First, we retrieve the segregation between slow and fast rotators at the correct frequency as expected."
+ Next. for [ast spinning stars. the theoretical curve agrees verv well wilh observations.," Next, for fast spinning stars, the theoretical curve agrees very well with observations."
+ Nevertheless. for slow rotation rates. (he spread around unity is significant and cannot be explained in a straightforward way by our predictions.," Nevertheless, for slow rotation rates, the spread around unity is significant and cannot be explained in a straightforward way by our predictions."
+ Clearly. some refinement of the mocel is still needed and under investigations. including other aspects of the plasma flow around an accreting magnetized neutron star.," Clearly, some refinement of the model is still needed and under investigations, including other aspects of the plasma flow around an accreting magnetized neutron star."
+ In this paper. we investigated huther theconsequences of forced oscillations induced," In this paper, we investigated further theconsequences of forced oscillations induced"
+little to increase the variable quasar population.,little to increase the variable quasar population.
+ There is a definite break in (he cumulative distribution of variable quasars al GCL-93., There is a definite break in the cumulative distribution of variable quasars at $ = 93$.
+ This break is seen at roughly the same GCL value in all time lag bins of reasonable leneth and was used to define our variable quasar population., This break is seen at roughly the same GCL value in all time lag bins of reasonable length and was used to define our variable quasar population.
+ The first-order structure functions relSFplot)) of the variable aud non-variable quasar populations show GCL to be an effective method [or separating quasars for further study., The first-order structure functions \\ref{SFplot}) ) of the variable and non-variable quasar populations show GCL to be an effective method for separating quasars for further study.
+ In a time-limited survey. the longest rest-lrame (ime lags correspond to (he most nearby objects. resulting in our observed anti-correlation between redshift and variabilitv.," In a time-limited survey, the longest rest-frame time lags correspond to the most nearby objects, resulting in our observed anti-correlation between redshift and variability."
+ We note that this as an inherent bias and do not make anv substantive claims as to the supposedly evolutionary correlation between redshift and variability seen by others., We note that this as an inherent bias and do not make any substantive claims as to the supposedly evolutionary correlation between redshift and variability seen by others.
+ Our smaller number of quasars do not allow the fine binning in redshift-magnitude-time lag space as used. by VandenBerketal.(2004) to decouple these competing effects., Our smaller number of quasars do not allow the fine binning in redshift-magnitude-time lag space as used by \citet{Vanden04} to decouple these competing effects.
+ With a range of over 4 orders of magnitude. the distribution of 2500 Ihuninosity density for (he entire quasar sample. as well as for the variable ancl non-variable sub-samiples. deviates significantly from Gaussian. complicating their statistical analvses.," With a range of over 4 orders of magnitude, the distribution of 2500 luminosity density for the entire quasar sample, as well as for the variable and non-variable sub-samples, deviates significantly from Gaussian, complicating their statistical analyses."
+ Nevertheless. (he Iuminosityv density median values of (he variable and non-variable quasars are consistent with the anti-correlation between variability. and luminosity seen in earlier work: (his result is supported by the INS test run between the two samples.," Nevertheless, the luminosity density median values of the variable and non-variable quasars are consistent with the anti-correlation between variability and luminosity seen in earlier work; this result is supported by the KS test run between the two samples."
+ The Lact. that oplical/UVN. variabilitv-Iuninosityv. anti-correlation is mirrored in both the IE. and. X-ray reeimes supports the standard unification moclel., The fact that optical/UV variability-luminosity anti-correlation is mirrored in both the IR and X-ray regimes supports the standard unification model.
+ We suggest this similar behavior arises [from the fact that the IR. ancl X-ray. emission are al least partially coupled to the UV/optical accretion disk emission. as described by the eenerally accepted model [or active galactic nuclei (see.e...Antonucei1993).," We suggest this similar behavior arises from the fact that the IR and X-ray emission are at least partially coupled to the UV/optical accretion disk emission, as described by the generally accepted model for active galactic nuclei \citep[see, e.g.,][]{Antonucci93}."
+.. IR radiation presumably originates in (he outer. cooler portions of (he accretion disk and in the torus from re-processed disk emission.," IR radiation presumably originates in the outer, cooler portions of the accretion disk and in the torus from re-processed disk emission."
+ X-ray emisson is due to UV photons fom the inner accretion disk being up-scatterecl by relativistic electrons in the corona., X-ray emisson is due to UV photons from the inner accretion disk being up-scattered by relativistic electrons in the corona.
+ Any correlations between intrinsic variations and luminosity originating in (he inner accretion disk would be expected to also be seen in both the IR. ancl X-ray. regimes., Any correlations between intrinsic variations and luminosity originating in the inner accretion disk would be expected to also be seen in both the IR and X-ray regimes.
+ A larger sample of X-ray detections and a robust strength of variability parameter. capable of reliably. characterizing an individual euasars photometric activity. are required to further investigate this multi-wavelength trend.," A larger sample of X-ray detections and a robust strength of variability parameter, capable of reliably characterizing an individual quasars' photometric activity, are required to further investigate this multi-wavelength trend."
+ The radio data. however. show a qualitatively different behavior: a positive correlation between optical variability and radio flux.," The radio data, however, show a qualitatively different behavior: a positive correlation between optical variability and radio flux."
+ Blazars teid to be primarily both radio-Ioud and hiehlv variable (see.e.g..Antonueci1993:Ulrichetal.1997.andreferencestherein).," Blazars tend to be primarily both radio-loud and highly variable \citep[see, e.g.,][and references therein]{Antonucci93,Ulrich97}."
+. Within the standard unification model. blazars ancl other core-dominant radio sources are seen when our line of sight aligns with relativistic jets perpendicular to the accretion disk.," Within the standard unification model, blazars and other core-dominant radio sources are seen when our line of sight aligns with relativistic jets perpendicular to the accretion disk."
+ The physics and continuum emission mechanisms for racio-loud AGN originate with the relativistic jets, The physics and continuum emission mechanisms for radio-loud AGN originate with the relativistic jets
+"differ by less than a factor of 3 and the cross-section is the effective size of the absorbing system as seen by the photon, we estimate that the extrapolation of the bound-free cross-section of the isolated hydrogen atom beyond Lyman edge introduces an uncertainty no larger than an order of magnitude on κι/(ν) calculated by Eq. 8..","differ by less than a factor of $3$ and the cross-section is the effective size of the absorbing system as seen by the photon, we estimate that the extrapolation of the bound-free cross-section of the isolated hydrogen atom beyond Lyman edge introduces an uncertainty no larger than an order of magnitude on $\kappa_{bf}(\nu)$ calculated by Eq. \ref{333}."
+" In helium-rich atmospheres, where the density can be as high as 2—3g/cm?, equation (6)) should be corrected when p& 0.1g/cm?, by a factor e""""-»* (Martynov1992), where wppe« is a thermal potential that arises from the correlations in the dense fluid."," In helium-rich atmospheres, where the density can be as high as $\rm 2-3 \, g/cm^3$, equation \ref{11}) ) should be corrected when $\rho \rm \wig>0.1\,g/cm^3$ , by a factor $e^{\rm w_{\rm H-pert}}$ \citep{MS}, where $\rm w_{\rm H-pert}$ is a thermal potential that arises from the correlations in the dense fluid."
+" Solving the Ornstein-Zernike equation in the Percus-Yevick approximation (Martynov1992) for fluid He, we have verified that in the helium-rich atmospheres of cool white dwarfs, e""*8-»«*<100."," Solving the Ornstein-Zernike equation in the Percus-Yevick approximation \citep{MS} for fluid $\rm He$, we have verified that in the helium-rich atmospheres of cool white dwarfs, $e^{\rm w_{\rm H-pert}}<100$."
+" Therefore, due to correlations in the most extreme case, Eq."," Therefore, due to correlations in the most extreme case, Eq."
+ 8 leads to an underestimate of Αν; of no more than two orders of magnitude., \ref{333} leads to an underestimate of $\kappa_{bf}$ of no more than two orders of magnitude.
+ We will seebelow that this does not affect our conclusions., We will seebelow that this does not affect our conclusions.
+"When ideally applied on the observed correlation functions from the model of ((29)), linear theory convolved with a velocity dispersion, the difference estimator D completely undoes’ the convolution with the velocity dispersion and returns a zero response if using the true 8 and συ.","When ideally applied on the observed correlation functions from the model of \ref{eq:modelPW}) ), linear theory convolved with a velocity dispersion, the difference estimator $D$ completely undoes the convolution with the velocity dispersion and returns a zero response if using the true $\beta$ and $\sigma_v$."
+ This is expected since D was explicitly based on the model in ((29))., This is expected since $D$ was explicitly based on the model in \ref{eq:modelPW}) ).
+ Much more interesting and realistic is when we consider the correlation functions built from the Sato and Matsubara N-body simulations., Much more interesting and realistic is when we consider the correlation functions built from the Sato and Matsubara $N$ -body simulations.
+ It is important to test if and at which scales the action of D is able to nullify or markedly reduce the overall effect of non linearities and convolution with the velocity dispersion., It is important to test if and at which scales the action of $D$ is able to nullify or markedly reduce the overall effect of non linearities and convolution with the velocity dispersion.
+ An example of the effectiveness of the prescription D from ((41)) and (42)) is provided in reffig:D.., An example of the effectiveness of the prescription $D$ from \ref{eq:defD}) ) and \ref{eq:estD}) ) is provided in \\ref{fig:D}.
+" It can be noticed that the two terms of the right-hand side of the transformed multipoles [20 and (rhs)D», are very ((41)),similar to each other."," It can be noticed that the two terms of the right-hand side (rhs) of \ref{eq:defD}) ), the transformed multipoles $D_0$ and $D_2$, are very similar to each other."
+" Also, the bulk of the baryonic feature is similar for both cases, apart from a small non linear contribution."," Also, the bulk of the baryonic feature is similar for both cases, apart from a small non linear contribution."
+ This means that our proposed processing on the monopole and quadrupole is capable of bringing the final products to a close agreement., This means that our proposed processing on the monopole and quadrupole is capable of bringing the final products to a close agreement.
+ The plot is just an example to display the general idea of the method., The plot is just an example to display the general idea of the method.
+" In the next section, we perform a full x? analysis on the simulation data formed from the total difference D to find if the best-fitting B is close to the fiducial value."," In the next section, we perform a full $\chi^2$ analysis on the simulation data formed from the total difference $D$ to find if the best-fitting $\beta$ is close to the fiducial value."
+" As already mentioned, we consider only a Gaussian velocity dispersion."," As already mentioned, we consider only a Gaussian velocity dispersion."
+" But if needed, more effective velocity dispersions might be easily implemented within the same framework in future analyses."," But if needed, more effective velocity dispersions might be easily implemented within the same framework in future analyses."
+ Two nice properties of the statistics D are that it does not depend on the mean density of the galaxies and that it is defined in terms of convolutions., Two nice properties of the statistics $D$ are that it does not depend on the mean density of the galaxies and that it is defined in terms of convolutions.
+" The first attribute can be easily proved by looking at tion((41)), where the presence of the Bessel function assures that the zero-mode contribution vanishes."," The first attribute can be easily proved by looking at \ref{eq:defD}) ), where the presence of the Bessel function assures that the zero-mode contribution vanishes."
+" The second feature potentially provides a certain degree of stability, compared with methods that make use of deconvolutions."," The second feature potentially provides a certain degree of stability, compared with methods that make use of deconvolutions."
+" To widen the range of applicability, the method will have to be further tested with N-body simulations that include also galaxy bias."," To widen the range of applicability, the method will have to be further tested with $N$ -body simulations that include also galaxy bias."
+" However, due to its simplicity and practicality, we believe that, with in case necessary phenomenological modifications to account for some of the effects of non linearities and galaxy bias, our method could be a useful tool to study RSDs."," However, due to its simplicity and practicality, we believe that, with in case necessary phenomenological modifications to account for some of the effects of non linearities and galaxy bias, our method could be a useful tool to study RSDs."
+" In a practical situation where data from a galaxy survey are available, the quantity D in ((42)) can be used directly in configuration space."," In a practical situation where data from a galaxy survey are available, the quantity $D$ in \ref{eq:estD}) ) can be used directly in configuration space."
+" Here the correlation functions can be measured, for example, by the Landy&Szalay(1993) estimator."," Here the correlation functions can be measured, for example, by the \cite{1993ApJ...412...64L} estimator."
+ We recommend the application of the following X? function in order to select from the data the best 6 and συ., We recommend the application of the following $\chi^2$ function in order to select from the data the best $\beta$ and $\sigma_v$.
+" The covariance matrix CovD;; relative to the distances s; and s; for D could be computed from for example a set of N-body simulations, like we are doing in this work, according to ((Dirm(s;. fid,)) ) fid,)) ), where irun ((D?™""(s;,is the realizationfid,) D((s;,index, that reaches the(56) total Nun=30 in our case."," The covariance matrix $\mathrm{Cov}D_{ij}$ relative to the distances $s_i$ and $s_j$ for $D$ could be computed from for example a set of $N$ -body simulations, like we are doing in this work, according to ( _v) ) ( _v) ), where $\mathrm{irun}$ is the realization index, that reaches the total $N_{\mathrm{run}}=30$ in our case."
+" The overbar indicates the mean over the realizations and the parameters § and συ are considered fixed to fiducial values, that need to be relatively close to the best-fit values."," The overbar indicates the mean over the realizations and the parameters $\beta$ and $\sigma_v$ are considered fixed to fiducial values, that need to be relatively close to the best-fit values."
+ A few extra iterations can update these values if necessary., A few extra iterations can update these values if necessary.
+" It is then natural to evaluate 2), Where we have explicitly indicated that D is a function (57)of distance separation, and cy."," It is then natural to evaluate _v), where we have explicitly indicated that $D$ is a function of distance separation, $\beta$ and $\sigma_v$."
+ We apply this statistic to the simulations results., We apply this statistic to the simulations results.
+ 'The monopole and quadrupole functions furnished to us by Sato and Matsubara are evaluated at discrete points separated by Ar=2-!Mpc.," The monopole and quadrupole functions furnished to us by Sato and Matsubara are evaluated at discrete points separated by $\Delta r=2\,h^{-1}\mathrm{Mpc}$."
+ It turns out that the derived covariance matrix (ονDj; is highly correlated and close to singular and in order to be able to stably invert it we have used singular value decomposition , It turns out that the derived covariance matrix $\mathrm{Cov}D_{ij}$ is highly correlated and close to singular and in order to be able to stably invert it we have used singular value decomposition (SVD).
+"The smaller eigenvalues, responsible for the instability,(SVD). are more affected by noise than the larger ones and it is necessary to operate a cut upon them, i.e. replace their inverses by zeros in the SVD representation of the inverse covariance matrix."," The smaller eigenvalues, responsible for the instability, are more affected by noise than the larger ones and it is necessary to operate a cut upon them, i.e. replace their inverses by zeros in the SVD representation of the inverse covariance matrix."
+" The situation is more conservative for the first eigenvalues but the statistical errors are larger, therefore a balance has to be struck."," The situation is more conservative for the first eigenvalues but the statistical errors are larger, therefore a balance has to be struck."
+1n currently [favoured cosmological models. dark matter halos grow via the merging together of smaller systems. leading to an ever-growing hierarchy. of halos.,"In currently favoured cosmological models, dark matter halos grow via the merging together of smaller systems, leading to an ever-growing hierarchy of halos."
+ Recent numerical simulations have demonstrated that the remnants of pre-existing dark matter halos which merged to become part of a larger system. (the Που). can survive for significant periods of time within the larger svstem (Mooreetal.1999:[xlvpinetal. 1998).," Recent numerical simulations have demonstrated that the remnants of pre-existing dark matter halos which merged to become part of a larger system (the “host”) can survive for significant periods of time within the larger system \cite{moore99,klypin98}."
+.. These subhalos orbit around in the potential of the host gradually: losing mass via tidal forces and spiralling in to ever smaller radii due to dvnamical friction., These subhalos orbit around in the potential of the host gradually losing mass via tidal forces and spiralling in to ever smaller radii due to dynamical friction.
+ Phese substructures (or at. least some subset of them) are presumably the abodes of satellite galaxies. such as those found in the Local Group. ancl of 11e majority of cluster galaxies.," These substructures (or at least some subset of them) are presumably the abodes of satellite galaxies, such as those found in the Local Group, and of the majority of cluster galaxies."
+ This substructure has attracted a great deal of interest since its discovery., This substructure has attracted a great deal of interest since its discovery.
+ Observational tests for its presence. though not vet conclusive. are in good. agreement with the theoretical expectations (Aletcall&Alacau2001:ChibaWP:Dalal&Ixochanek|. 2002).," Observational tests for its presence, though not yet conclusive, are in good agreement with the theoretical expectations \cite{metmad01,chiba01,dal01}."
+. There has been much work conducted in which the distribution ancl properties of substructures. their ellects on galaxy. disks. and so on were examined (Chignaetal.1998:Tormen.Diaferio&Stacel 2004).," There has been much work conducted in which the distribution and properties of substructures, their effects on galaxy disks and so on were examined \cite{ghigna98,tormen98,vdb99,zhang02,benson04,gao04,diemand04}."
+.. While the orbital parameters of substructures have been measured in the past this measurement has often been at the end. point of the substructure evolution (i.c. at the present cay) when significant dynamical evolution in the orbital parameters is expected. (e.g. Ghignaal. 1998))., While the orbital parameters of substructures have been measured in the past this measurement has often been at the end point of the substructure evolution (i.e. at the present day) when significant dynamical evolution in the orbital parameters is expected (e.g. \pcite{ghigna98}) ).
+ Exceptions to this are the works of (1997).. Vitvitskaetal.(2002). and. Whochfarert," Exceptions to this are the works of \scite{tormen97}, \scite{vitvit02} and \scite{khochfar04}."
+ (2004).. Tormen(1997) ancl Whochlar&Burkert(2004) both identified progenitors of halos in their N-bodvy simulations and. measured the orbital parameters of them. while Vitvitskaetal.(2002). searched. for. pairs of halos about to merge and measured the orbital parameters. of these.," \scite{tormen97} and \scite{khochfar04} both identified progenitors of halos in their N-body simulations and measured the orbital parameters of them, while \scite{vitvit02}
+ searched for pairs of halos about to merge and measured the orbital parameters of these."
+ These works have typically made use of rather small samples of orbits and. (perhaps consequently) have been unable to fully characterise the two dimensional distribution of orbital parameters needed to construct realistic orbits., These works have typically made use of rather small samples of orbits and (perhaps consequently) have been unable to fully characterise the two dimensional distribution of orbital parameters needed to construct realistic orbits.
+ The distribution of initial orbital paranicters of substructure halos at the time of mereing into the host system is a particularly interesting property as it represents the initial conditions which determine the later evolution of the substructure within the host., The distribution of initial orbital parameters of substructure halos at the time of merging into the host system is a particularly interesting property as it represents the initial conditions which determine the later evolution of the substructure within the host.
+ Phe cllectivencss of many processes invoked. to explain the morphological transformation of galaxies in clusters (e.g. rani pressure stripping. tidal mass loss. galaxy harassment. ete.)," The effectiveness of many processes invoked to explain the morphological transformation of galaxies in clusters (e.g. ram pressure stripping, tidal mass loss, galaxy harassment, etc.)"
+ depend crucially on the nature of the galaxy orbit. (see. for example. Moore.Lake&Katz1998:AXbacdi.BowerNavarro. 2000)).," depend crucially on the nature of the galaxy orbit (see, for example, \pcite{moore98,abadi00}) )."
+ The distribution of orbits will also determine the rate of ealaxv mergers and therefore the degree. of heating and rate of morphological transformation experienced by galaxy disks., The distribution of orbits will also determine the rate of galaxy mergers and therefore the degree of heating and rate of morphological transformation experienced by galaxy disks.
+ Taking a more practical point of view. recent analytic models of satellite halo orbits (Bensonetal.2002: have been able to follow the orbital evolution of satellites quite accurately. but these models are," Taking a more practical point of view, recent semi-analytic models of satellite halo orbits \cite{benson02,taybab04} have been able to follow the orbital evolution of satellites quite accurately, but these models are"
+data for AGN obtained over many vears.,data for AGN obtained over many years.
+ We randomly select. points along the lighteurve separated by 5 vears. (he time interval sampled by our HDF images.," We randomly select points along the lightcurve separated by 5 years, the time interval sampled by our HDF images."
+ Because long term variability surveys for low-Iuminositv AGN have not vet been published. we have conducted this test with two sets of QSO lishteurve data: a sample of PG quasars from Giveon (1999) and a sample of SA57 quasars [rom Trevese (1994).," Because long term variability surveys for low-luminosity AGN have not yet been published, we have conducted this test with two sets of QSO lightcurve data; a sample of PG quasars from Giveon (1999) and a sample of SA57 quasars from Trevese (1994)."
+ Giveon monitored 42 quasars over T vears with a tvpical sampling interval of 40 days while Trevese monitored 64 over 15 vears about once a vear., Giveon monitored 42 quasars over 7 years with a typical sampling interval of 40 days while Trevese monitored 64 over 15 years about once a year.
+ Sampling these lieltcurves every 5 vears and assuming photometric errors (vpical of our ΠΗΡΕ images. we can estimate the probability that these sources would be detected in our survey.," Sampling these lightcurves every 5 years and assuming photometric errors typical of our HDF images, we can estimate the probability that these sources would be detected in our survey."
+ We find that “75 would be detected as varving by al least 30 with only 2 observations separated bv 5 vears., We find that $\sim$ would be detected as varying by at least $\sigma$ with only 2 observations separated by 5 years.
+ Therefore. (he results of our analysis are likely to wnelerrepresent (he true number of variable nuclei by ~2025%.," Therefore, the results of our analysis are likely to underrepresent the true number of variable nuclei by $\sim$."
+.. The second incompleteness issue results rom the use of a fixed aperture to detect nuclear variability., The second incompleteness issue results from the use of a fixed aperture to detect nuclear variability.
+ We have chosen a small. fixed. aperture to minimize dilution of the AGN light from the underlving host galaxy (o be more sensitive to AGN varving within bright hosts.," We have chosen a small, fixed aperture to minimize dilution of the AGN light from the underlying host galaxy to be more sensitive to AGN varying within bright hosts."
+" We set (his aperture to the size of an unresolved point source in the HST images (0.15"" in diameter).", We set this aperture to the size of an unresolved point source in the HST images $\arcsec$ in diameter).
+ The fixed aperture will include (he same amount of lisht from an unresolved source regardless of its redshift., The fixed aperture will include the same amount of light from an unresolved source regardless of its redshift.
+ ILowever. the aperture will contain a larger fraction of the underlving. resolved. galaxy [or higher redshift objects than for low-z sources.," However, the aperture will contain a larger fraction of the underlying, resolved galaxy for higher redshift objects than for low-z sources."
+ Because of this. the dilution of the nuclear light bv the galaxy increases as a [function of redshift ancl. consequently. as a function of nuclear magnitude.," Because of this, the dilution of the nuclear light by the galaxy increases as a function of redshift and, consequently, as a function of nuclear magnitude."
+ If we assume a disk-dominated. r.—1 kpc ealaxv containing an AGN that is of the host galaxy. light and varving by Am=0.3 nagnitudes. the measured Am within our fixed aperture would be 0.15 magnitudes at z=0.2 but only 0.06 magnitudes at z-1.," If we assume a disk-dominated, $_e$ =1 kpc galaxy containing an AGN that is of the host galaxy light and varying by $\Delta$ m=0.3 magnitudes, the measured $\Delta$ m within our fixed aperture would be 0.15 magnitudes at z=0.2 but only 0.06 magnitudes at z=1."
+ The observed total nuclear magnitude would drop by ~3 nagnitudes over this interval., The observed total nuclear magnitude would drop by $\sim$ 3 magnitudes over this interval.
+ The increasing dilution with redshilt is compounded by the increasing photometric errors with magnitude., The increasing dilution with redshift is compounded by the increasing photometric errors with magnitude.
+" For an AGN that has varied by Am=0.30.2 nags. consistent with tvpieal structure function values for AGN (Trevese 1994). and reasonable estimates for the galaxy. magnitude. morphology. physical size and true AGN Taction. we estimate that incompleteness due to significant AGN dilution begins to effect our survey at magnitudes between V,,,,2:26.5 lo Vay 27.5."," For an AGN that has varied by $\Delta$ m=0.3–0.2 mags, consistent with typical structure function values for AGN (Trevese 1994), and reasonable estimates for the galaxy magnitude, morphology, physical size and true AGN fraction, we estimate that incompleteness due to significant AGN dilution begins to effect our survey at magnitudes between $_{nuc}$$\simeq$ 26.5 to $_{nuc}$$\simeq$ 27.5."
+" As can be seen in Figure 4. all of the variables lie αἱ magnitudes greater than V,,.—27.3."," As can be seen in Figure 4, all of the variables lie at magnitudes greater than $_{nuc}$ =27.3."
+ While (his may partially be due to a real astvophvsical effect. (see discussion in Section 4). (his limiting magnitude is within the range of our incompleteness estimate above.," While this may partially be due to a real astrophysical effect (see discussion in Section 4), this limiting magnitude is within the range of our incompleteness estimate above."
+" Therefore. allhough we survey galaxy. nuclei to V,,,.2229. our limit for detecting AGN at a unilormlv varving level wilhin galaxy nuclei is probably closer to ο1.5. wilh decreasing completeness fainter than V,,,7226.5."," Therefore, although we survey galaxy nuclei to $_{nuc}$$\simeq$ 29, our limit for detecting at a uniformly varying level within galaxy nuclei is probably closer to $_{nuc}$$\simeq$ 27.5, with decreasing completeness fainter than $_{nuc}$$\simeq$26.5."
+"Primitive meteorites contain daughter products of the decay of short-lived radioisotopes (SLRIs) such as “AL Ca. Mn. and ""Fe. distributed in different minerals in a way that inclicates (he parent isotopes were slill alive at the lime of their incorporation into the refractory inclusions and chondrules that record the earliest history of the solar svstem.","Primitive meteorites contain daughter products of the decay of short-lived radioisotopes (SLRIs) such as $^{26}$ Al, $^{41}$ Ca, $^{53}$ Mn, and $^{60}$ Fe, distributed in different minerals in a way that indicates the parent isotopes were still alive at the time of their incorporation into the refractory inclusions and chondrules that record the earliest history of the solar system."
+ The presence of Το is particularly significant. as its production requires stellar nucleosvnthesis (Tachibana Luss 2003: Tachibana et al.," The presence of $^{60}$ Fe is particularly significant, as its production requires stellar nucleosynthesis (Tachibana Huss 2003; Tachibana et al."
+ 2006)., 2006).
+ Given halllives on the order of ~LO° vr. the evidence for these raclioisolopes suggests that the same stellar source that synthesized them may well," Given half-lives on the order of $\sim 10^6$ yr, the evidence for these radioisotopes suggests that the same stellar source that synthesized them may well"
+ Iureceut vears. the impact of photoevaporation ou the evolution of protoplanctary disks has received much attention (22?7)..,"In recent years, the impact of photoevaporation on the evolution of protoplanetary disks has received much attention \citep{johnstone98,bally98,adams04,gorti09}."
+ Recent observations confini that there Is sufficient FUY aud X-rav flux from voung. solu-tvpe sars to drive photoevaporative mass loss iu the surroundiug uebulae (?)..," Recent observations confirm that there is sufficient FUV and X-ray flux from young, solar-type stars to drive photoevaporative mass loss in the surrounding nebulae \citep{ingleby11}."
+ Although photoevaporation las been investigated iun detail im the coutext of planet formation. it has wet to be applied to circunplauetary disks aud the formaion of regular satellites.," Although photoevaporation has been investigated in detail in the context of planet formation, it has yet to be applied to circumplanetary disks and the formation of regular satellites."
+ Receut models have looked. at he combined effects of viscous evolution aud ploocvaporation ou the evolution of protoplanetary webulae (22).," Recent models have looked at the combined effects of viscous evolution and photoevaporation on the evolution of protoplanetary nebulae \citep{clarke07,mitchell10}."
+ It is tempting to apply a direct analogy between protoplauctary and protosatellite nebulae. but it wotld be incorrect ο view a protosatellite nebula :is simply a scded down version of a protoplaretary nebula.," It is tempting to apply a direct analogy between protoplanetary and protosatellite nebulae, but it would be incorrect to view a protosatellite nebula as simply a scaled down version of a protoplanetary nebula."
+" The slow iufall of material from fjo solar nebula ouο cireumplanetary disks. as proposed bv 777. causes them to behave very iferentle frou, protoplanetary analogues."," The slow infall of material from the solar nebula onto circumplanetary disks, as proposed by \citet{canup02,canup09,ward10}, causes them to behave very differently from protoplanetary analogues."
+ Iu paricular. the combination of mass iufall and photoevaporative nass loss can allow for quasi steady-state solutions.," In particular, the combination of mass infall and photoevaporative mass loss can allow for quasi steady-state solutions."
+ The giant planets have similar ratios of the total lass contained witlin the satellite svstems to the of massthe host planets (ΑΕμα~10./! Myancc)-, The giant planets have similar ratios of the total mass contained within the satellite systems to the mass of the host planets ${\rm M_{satellites} \sim 10^{-4}\ M_{planet}}$ ).
+ Iu seueral. the mass of a satellite increases with increasing distance from the plauct. reaches a," In general, the mass of a satellite increases with increasing distance from the planet, reaches a"
+1996: Wheeler. Meier. Wilson 2002: Vasil Drumamell 2008. 2009).,"1996; Wheeler, Meier, Wilson 2002; Vasil Brummell 2008, 2009)."
+ The orbital motion of (he protostars in the cluster eravitational potential ancl the orbital motion of gas in individual circumstellar disks may each contribute to the amplification of the magnetic fields (Ixappylà el al., The orbital motion of the protostars in the cluster gravitational potential and the orbital motion of gas in individual circumstellar disks may each contribute to the amplification of the magnetic fields (Käppylä et al.
+ 2008)., 2008).
+ When equipartition is reached. B?zz8zp(r)c. where pir) is the local gas density and c; is (he effective sound speed including turbulent motions.," When equipartition is reached, $B^2 \approx 8 \pi \rho (r) c_s^2$, where $\rho (r)$ is the local gas density and $c_s$ is the effective sound speed including turbulent motions."
+" For the envelope having (the density prolile described in Appendix 1. the equipartition magnetic οποιον densityv. ad radius rods given bv ej,=B?(r)/8xοι) where ο, is the local effective sound speed Gncluding thermal and turbulent motion) and (he subscript e relers to the envelope."," For the envelope having the power-law density profile described in Appendix 1, the equipartition magnetic energy density at radius $r$ is given by $\epsilon _{Be} = B^2 (r) / 8 \pi = \rho (r) c^2 _s$ where $c_s$ is the local effective sound speed (including thermal and turbulent motion) and the subscript $e$ refers to the envelope."
+" Under (his assumption. (he magnetic energv in a shell al radius ¢ and thickness dr is which can be integrated from the inner edge of the envelope to a radius r to give Selling r=rq, gives By,=M,c2."," Under this assumption, the magnetic energy in a shell at radius $r$ and thickness $dr$ is which can be integrated from the inner edge of the envelope to a radius $r$ to give Setting $r = r_{out}$ gives $E_{Be} = M_e c_s^2$."
+" For M, = 60 M. and ος = 2 km ! this results in Eg,z5x10” eres for the equipartition magnetic enerev that could be stored in the envelope.", For $M_e$ = 60 $_{\odot}$ and $c_s$ = 2 km $^{-1}$ this results in $E_{Be} \approx 5 \times 10^{45}$ ergs for the equipartition magnetic energy that could be stored in the envelope.
+ Ifimaenetic stress reaches equipartition with the effective sound speed within circumstellar disks. a similar calculation vields Ep;=Mc; where the subscript d. refers to the disk quantities.," If magnetic stress reaches equipartition with the effective sound speed within circumstellar disks, a similar calculation yields $E_{Bd} = M_d c_d^2$ where the subscript $d$ refers to the disk quantities."
+ The ellective sound speeds in (he disks are likely to be considerably larger due to proximity to the massive stars., The effective sound speeds in the disks are likely to be considerably larger due to proximity to the massive stars.
+" Using c, = 10 km |!1 gives Ej;=2.0xLOYMei, ergs for the equipartition magnetic energy that could be stored in a disk.", Using $c_s$ = 10 km $^{-1}$ gives $E_{Bd} = 2.0 \times 10^{45} M_1 c^2_{10}$ ergs for the equipartition magnetic energy that could be stored in a disk.
+ Here M4 is in units of LAL. and ey) is in units of 10 km !., Here $M_1$ is in units of 1 $_{\odot}$ and $c_{10}$ is in units of 10 km $^{-1}$.
+ Magnetic fields may become stronger (han equipartition with the local pressure in the disk or envelope., Magnetic fields may become stronger than equipartition with the local pressure in the disk or envelope.
+ Keplerian shear can wind-up the field so (hat its enerev approaches the eravitational potential. a circumstance discussed in the context of cireumstellar disks by Shu et al. (," Keplerian shear can wind-up the field so that its energy approaches the gravitational potential, a circumstance discussed in the context of circumstellar disks by Shu et al. ("
+2008).,2008).
+ The condition that Alfven speed (or effective sound speed) is equal to the gravitational escape speed sets an upper bound on the fiell-strength., The condition that Alfven speed (or effective sound speed) is equal to the gravitational escape speed sets an upper bound on the field-strength.
+ Thus. where M is the mass enclosed inside radius r.," Thus, $V^2_A = B^2 / 4 \pi \rho \approx GM/r$ where $M$ is the mass enclosed inside radius $r$."
+ In this case. the field could have stored more than 107 eres of magnetic enerev in the volume inhabited by the cluster.," In this case, the field could have stored more than $10^{47}$ ergs of magnetic energy in the volume inhabited by the cluster."
+ The removal of about 40 AL. of stars would have left the circum-cluster gas severely inagnelically over-pressured., The removal of about 40 $_{\odot}$ of stars would have left the circum-cluster gas severely magnetically over-pressured.
+" The release of magnetic energy. in effect a ""magnetic bomb (Matt. Frank. DBlackman 2004. 2006). may have been a major source of energy lor launching"," The release of magnetic energy, in effect a `magnetic bomb' (Matt, Frank, Blackman 2004, 2006), may have been a major source of energy for launching"
+"As with the Millennium Simulation, at each simulation snapshot particle groups were identified on-the-fly according to the friends-of-friends algorithm (FoF), with a linking length parameter of b=0.2 (?)..","As with the Millennium Simulation, at each simulation snapshot particle groups were identified on-the-fly according to the friends-of-friends algorithm (FoF), with a linking length parameter of $b=0.2$ \citep{defw85}."
+" Subsequently, self-bound substructures within these groups were found usingthe algorithm (?).."," Subsequently, self-bound substructures within these groups were found usingthe algorithm \citep{subfind2001}."
+" Finally, the progenitors and descendents of each particle group were found, creating a *merger tree' structure allowing haloes to be tracked over time."," Finally, the progenitors and descendents of each particle group were found, creating a `merger tree' structure allowing haloes to be tracked over time."
+" The merger tree algorithm (and associated halo definition) used is that described in ?,, which was originally designed for use with the semi-analytic galaxy formation model and the Millennium (??).."," The merger tree algorithm (and associated halo definition) used is that described in \cite{harker2006}, which was originally designed for use with the semi-analytic galaxy formation model and the Millennium \citep{helly2003,bowergalform2006}."
+" Since the construction of haloes and mergers trees from Simulatiorf]particle groups in discrete snapshots is essential for the current work, we will now describe this process in more detail."," Since the construction of haloes and mergers trees from particle groups in discrete snapshots is essential for the current work, we will now describe this process in more detail."
+" The preliminary set of haloes consists, at each snapshot, of the FoF particle groups, which in turn contain the self-bound substructures identified by (these include the main body of the halo, plus less massive subhaloes), as well as so-called ""fuzz"" particles not gravitationally bound to any structure in the halo."," The preliminary set of haloes consists, at each snapshot, of the FoF particle groups, which in turn contain the self-bound substructures identified by (these include the main body of the halo, plus less massive subhaloes), as well as so-called “fuzz” particles not gravitationally bound to any structure in the halo."
+" It is well known that the purely spatial nature of the FoF algorithm allows multiple objects to be linked together spuriously, in the sense that although they are close together they might not necessarily be physically connected."," It is well known that the purely spatial nature of the FoF algorithm allows multiple objects to be linked together spuriously, in the sense that although they are close together they might not necessarily be physically connected."
+" A common example is that of two close objects linked with a tenuous bridge of particles, which the FoF algorithms identifies as a single “halo” (see ? for a detailed discussion and comparison of the effect of groupfinders on the measured halo angular momentum and related properties)."," A common example is that of two close objects linked with a tenuous bridge of particles, which the FoF algorithms identifies as a single “halo” (see \citealt{bett07} for a detailed discussion and comparison of the effect of groupfinders on the measured halo angular momentum and related properties)."
+" Including additional physical information — e.g. gravitational binding fromSusrFinp,, and temporal evolution from merger trees — allows the operational definition of a halo to be refined, to match better our physical intuition of what a halo is, and thus, when used in conjunction with semi-analytic models of galaxy formation, to allow better comparisons with both observations and hydrodynamic simulations. """," Including additional physical information – e.g. gravitational binding from, and temporal evolution from merger trees – allows the operational definition of a halo to be refined, to match better our physical intuition of what a halo is, and thus, when used in conjunction with semi-analytic models of galaxy formation, to allow better comparisons with both observations and hydrodynamic simulations. “"
+"splitting"" algorithm is applied to the basic FoF halo catalogues, whereby spuriously linked subhaloes are split off from their original FoF groups and identified as separate haloes in their own right.","splitting” algorithm is applied to the basic FoF halo catalogues, whereby spuriously linked subhaloes are split off from their original FoF groups and identified as separate haloes in their own right."
+ A subhalo is split off from its original FoF parent if it satisfies at least one of the following conditions: (1) The distance between the subhalo centre and the parent centre is more than twice the half-mass radius of the parent; or (2) the subhalo still has more than 7596 of the mass it had when it was last identified as a separate halo., A subhalo is split off from its original FoF parent if it satisfies at least one of the following conditions: (1) The distance between the subhalo centre and the parent centre is more than twice the half-mass radius of the parent; or (2) the subhalo still has more than $75\%$ of the mass it had when it was last identified as a separate halo.
+ This yields halo catalogues containing more objects than the corresponding FoF catalogues., This yields halo catalogues containing more objects than the corresponding FoF catalogues.
+ ? showed that these *merger tree haloes' are a great improvement on both the simple FoF groups and groups found from a simple spherical overdensity criterion., \cite{bett07} showed that these `merger tree haloes' are a great improvement on both the simple FoF groups and groups found from a simple spherical overdensity criterion.
+" Merger trees for these haloes are constructed by tracking the particles that constitute the subhaloes between each snapshot, starting at early times and continuing to redshift z=0."," Merger trees for these haloes are constructed by tracking the particles that constitute the subhaloes between each snapshot, starting at early times and continuing to redshift $z=0$."
+ Each halo or subhalo can have at most one descendent in a later snapshot., Each halo or subhalo can have at most one descendent in a later snapshot.
+ The most bound 10% of a subhalo’s mass (or 10 most bound particles if that is more massive) is located in the next snapshot., The most bound $10\%$ of a subhalo's mass (or $10$ most bound particles if that is more massive) is located in the next snapshot.
+" Occasionally, these particles might no longer reside in a subhalo in the next snapshot: the subhalo might have temporarily dropped below FIND’’s detection limit, or might be passing through a high density region and be interpreted as unbound matter around that density peak."," Occasionally, these particles might no longer reside in a subhalo in the next snapshot: the subhalo might have temporarily dropped below 's detection limit, or might be passing through a high density region and be interpreted as unbound matter around that density peak."
+" In practice therefore, the next five snapshots are scanned to find the earliest time when these particles are again in subhaloes."," In practice therefore, the next five snapshots are scanned to find the earliest time when these particles are again in subhaloes."
+ The descendent subhalo is then identified as the subhalo containing the largest number of those most bound particles., The descendent subhalo is then identified as the subhalo containing the largest number of those most bound particles.
+ The descendent of a halo is identified as the halo whose most massive substructure (i.e. the main self-bound halo component) is the descendent of its own most massive substructure., The descendent of a halo is identified as the halo whose most massive substructure (i.e. the main self-bound halo component) is the descendent of its own most massive substructure.
+ Itis possible that a subhalo's mass ends up distributed between two (or more) subhaloes in a subsequent snapshot., It is possible that a subhalo's mass ends up distributed between two (or more) subhaloes in a subsequent snapshot.
+" While one will be identified as the descendent, the other will be left as a separate “orphan” object without a progenitor."," While one will be identified as the descendent, the other will be left as a separate “orphan” object without a progenitor."
+" This situation is known as a de-merger, and is a physical effect separate to the splitting of groups described above — here, sets of particles physically end up in separate objects as the simulation evolves."," This situation is known as a de-merger, and is a physical effect separate to the splitting of groups described above – here, sets of particles physically end up in separate objects as the simulation evolves."
+" The end result of this process is a catalogue of haloes (groups of self-bound substructures) identified at each snapshot, with at most one descendent and one or more progenitors."," The end result of this process is a catalogue of haloes (groups of self-bound substructures) identified at each snapshot, with at most one descendent and one or more progenitors."
+" Each halo identified at z—O is the root of its own tree, which branches into many progenitor haloes at preceding output times."," Each halo identified at $z=0$ is the root of its own tree, which branches into many progenitor haloes at preceding output times."
+" In this paper, we study the evolution of properties of individual haloes that at z=0 have a mass corresponding roughly to that of the Milky Way."," In this paper, we study the evolution of properties of individual haloes that at $z=0$ have a mass corresponding roughly to that of the Milky Way."
+" After identifying an appropriate halo at z=0, we track its evolution back by finding its most massive progenitor at the preceding snapshot, then finding the most massive of halo's progenitors, and so on."," After identifying an appropriate halo at $z=0$, we track its evolution back by finding its most massive progenitor at the preceding snapshot, then finding the most massive of halo's progenitors, and so on."
+" It is important to note that, just like the halo definition and galaxy formation model, the halo merger tree algorithm is not by any means unique — even within a given N-body simulation."," It is important to note that, just like the halo definition and galaxy formation model, the halo merger tree algorithm is not by any means unique – even within a given $N$ -body simulation."
+ The halo merger trees used for the “MPA” semi-analytic models of the Millennium Simulation (e.g. ? and ?)) in fact track the binding-energy-weighted massin the subhaloes., The halo merger trees used for the “MPA” semi-analytic models of the Millennium Simulation (e.g. \citealt{mill2005} and \citealt{2007MNRAS.375....2D}) ) in fact track the binding-energy-weighted massin the subhaloes.
+" Other methods that use splitting/stitching algorithms similar to the one used here include those by e.g. ??,, ?;; see also ?.."," Other methods that use splitting/stitching algorithms similar to the one used here include those by e.g. \cite{2008MNRAS.386..577F,2009MNRAS.394.1825F}, \cite{2009ApJ...701.2002G}; see also \cite{2006ApJ...647..763M}."
+ ? provide a recent detailed study of halo definition and merger tree algorithms., \cite{2009A&A...506..647T} provide a recent detailed study of halo definition and merger tree algorithms.
+ Various properties of the haloes are computed at each output time., Various properties of the haloes are computed at each output time.
+" Properties are computed in thecentre-of-momentum frame of each halo, and in physical rather than comoving coordinates."," Properties are computed in thecentre-of-momentum frame of each halo, and in physical rather than comoving coordinates."
+" Each halo consists of a set of N, particles, with each particle i having mass m;= mp, position x; and velocity v;."," Each halo consists of a set of $\Np$ particles, with each particle $i$ having mass $m_i =
+\mp$ , position $\vv{x}_i$ and velocity $\vv{v}_i$ ."
+" The halo mass is therefore M,=»;*Npm; N,m,.", The halo mass is therefore $\Mh= \sum_{i=1}^{\Np}m_i = \Np\mp$ .
+ The halo centre is taken, The halo centre is taken
+Particle acceleration remains a mystery in solar flares. as well as in cosmic ravs from the extrasolar svstems.,"Particle acceleration remains a mystery in solar flares, as well as in cosmic rays from the extrasolar systems."
+ Observationallv. (he energetic particles always show a single or double power-law spectral behavior.," Observationally, the energetic particles always show a single or double power-law spectral behavior."
+ Three mechanisms have been proposed (seeMilleretal.1997:Aschwanden2002.[or reviews).. i.e.. the direct ciurent (DC) electric field acceleration (e.g..Sturrock1963).. turbulence (or stochastic) acceleration (e.g..Miller 1998).. and shock acceleration (e.g..eBlandford&Ostriker1978).," Three mechanisms have been proposed \citep[see][for reviews]{mill97, asch02}, i.e., the direct current (DC) electric field acceleration \citep[e.g.,][]{stur68}, turbulence (or stochastic) acceleration \citep[e.g.,][]{ml98}, , and shock acceleration \citep[e.g.,][]{blan78}."
+. MagneticC» reconnection. asthe effective," Magnetic reconnection, asthe effective"
+was based on short scans of a nuniber of strong unresolved sources before aud after each svuthesis observation of1.,was based on short scans of a number of strong unresolved sources before and after each synthesis observation of.
+" Experience has shown that this procedure is correct to <1"" in position aud ~h ylus in flux deusity (IIunstead 1991).", Experience has shown that this procedure is correct to $<$ $''$ in position and $\sim$ rms in flux density \cite{hun91}) ).
+ The AIPS task was used to fit a euussian to cach iuage at the known position of1., The AIPS task was used to fit a gaussian to each image at the known position of.
+ In a few cases Where failed to converge (sequence ΗΠΙΟΥΣ 16591810. 1659185123. 16591822. 16591823. 165918285. 165918529) because of à marginal detection. we used the task to measure the flux density.," In a few cases where failed to converge (sequence numbers 16594810, 16594813, 16594822, 16594823, 16594828, 16594829) because of a marginal detection, we used the task to measure the flux density."
+ By iutegratiug over two concentric reeious centred on WE were able ο correct for non-zero or sloping backgrounds., By integrating over two concentric regions centred on we were able to correct for non-zero or sloping backgrounds.
+ Tn order to minimise svstematic errors we then refined the flux calibration bv usiug to measure peak flux deusities for three reference sources SL. S2 aud οὐ (Fie. 13) ," In order to minimise systematic errors we then refined the flux calibration by using to measure peak flux densities for three reference sources S1, S2 and S3 (Fig. \ref{fig1}) )"
+in the field ofI., in the field of.
+ All three were found to be unresolved. aud to show the same trends with epoch. so we adopted the sununed fux density for scaling the flux density of1. on the assuuptiou that the MOST calibration was correct on average over the monitoriug period.," All three were found to be unresolved, and to show the same trends with epoch, so we adopted the summed flux density for scaling the flux density of, on the assumption that the MOST calibration was correct on average over the monitoring period."
+ The rms scatter of the correction factors was 5.2*4.. consistent with expectation.," The rms scatter of the correction factors was $\sim$, consistent with expectation."
+ The nus noise level in the vicinity of Was estimated from the statistics of regions ~500 pixels to the south and west of[., The rms noise level in the vicinity of was estimated from the statistics of regions $\sim$ 500 pixels to the south and west of.
+ We have adopted the 11011 ruis nolse level (n 1àJv/bei) as the flux deusityv error for cach epoch. aud for setting a lait for the nou-doetectious.," We have adopted the mean rms noise level (in mJy/beam) as the flux density error for each epoch, and for setting a limit for the non-detections."
+ The final corrected flux deusities and adopted. errors are tabulate in Table 1.. while Table 2 eives the positions and flux densities of the three reference sources.," The final corrected flux densities and adopted errors are tabulated in Table \ref{tbl1}, while Table \ref{tbl2} gives the positions and flux densities of the three reference sources."
+ The truncated Juliau date (TID) is defined here as TJD=JD—2100000 aud specified at mid-obscrvation.," The truncated Julian date (TJD) is defined here as ${\rm TJD} = {\rm
+JD} - 2400000$ and specified at mid-observation."
+ Note the fall iu ius noise between 199[ aud 1995. resulting frou the installation of new pre-uupli&ers as part of the hardware upgrade of MOST.," Note the fall in rms noise between 1994 and 1995, resulting from the installation of new pre-amplifiers as part of the hardware upgrade of MOST."
+ Figure d shows a deep image of the fiele of obtained by co-adding selected. images from the full dataset (see Table 13)., Figure \ref{fig1} shows a deep image of the field of obtained by co-adding selected images from the full dataset (see Table \ref{tbl1}) ).
+ The backeround noise level of 0.3 1Jvbeam riis includes contributions from weak sources. low level Galactic ciission aud telescope artifacts.," The background noise level of 0.3 mJy/beam rms includes contributions from weak sources, low level Galactic emission and telescope artifacts."
+ There Is no evidence for any prominent extended radio emission which might be associated with1. iu contrast to a similar montage of thefield around 1," There is no evidence for any prominent extended radio emission which might be associated with, in contrast to a similar montage of thefield around ,"
+ There Is no evidence for any prominent extended radio emission which might be associated with1. iu contrast to a similar montage of thefield around 10," There is no evidence for any prominent extended radio emission which might be associated with, in contrast to a similar montage of thefield around ,"
+ There Is no evidence for any prominent extended radio emission which might be associated with1. iu contrast to a similar montage of thefield around 10.," There is no evidence for any prominent extended radio emission which might be associated with, in contrast to a similar montage of thefield around ,"
+ There Is no evidence for any prominent extended radio emission which might be associated with1. iu contrast to a similar montage of thefield around 10..," There is no evidence for any prominent extended radio emission which might be associated with, in contrast to a similar montage of thefield around ,"
+reasonable if (he clearing timescale has been greater (han 5 \ivrs.,reasonable if the clearing timescale has been greater than 5 Myrs.
+ Furthermore. we propose that the low fraction of sources with A-band excesses found by ?.anclsample is less abnormal if a significant [fraction of sources are older (han previously believed.," Furthermore, we propose that the low fraction of sources with $K$ -band excesses found by \citet[][and confirmed from our larger sample]{alba07} is less abnormal if a significant fraction of sources are older than previously believed."
+ We measure the stellar mass functions for stus within our completeness limits and find a slope of D=—1.0940.18. in good agreement with the potentially universal value estimated bv ?..," We measure the stellar mass functions for stars within our completeness limits and find a slope of $\Gamma = -1.09 \pm 0.13$, in good agreement with the potentially universal value estimated by \citet{krou02}."
+ This is the deepest and most accurate measurement of (he mass function in Cvg OD2. building on previously measurements (hat have found either steeper (?) or shallower (7?) slopes.," This is the deepest and most accurate measurement of the mass function in Cyg OB2, building on previously measurements that have found either steeper \citep{knod00} or shallower \citep{mass91} slopes."
+ A steepening of the mass function is observed at high masses. as found by 2.. which we attribute to the demise of the highest mass members of the older generations in the region.," A steepening of the mass function is observed at high masses, as found by \citet{kimi07}, which we attribute to the demise of the highest mass members of the older generations in the region."
+ An estimation of the total size of the region suggests that the stellar density drops off faster than estimated by ?.., An estimation of the total size of the region suggests that the stellar density drops off faster than estimated by \citet{knod00}.
+ This. combined with our mass finetion estimate. reveals a total mass of ~3x10 AL...," This, combined with our mass function estimate, reveals a total mass of $\sim 3 \times 10^4$ $_{\odot}$."
+ This makes Cve OB? similar in mass to many of the recently discovered massive clusters in our Galaxy such as the Arches and Quintuplet clusters., This makes Cyg OB2 similar in mass to many of the recently discovered massive clusters in our Galaxy such as the Arches and Quintuplet clusters.
+ However. we note that comparing other properties of (hese massive star forming regions. such as their central densities. reveals differences that may hint at different formation mechanisms.," However, we note that comparing other properties of these massive star forming regions, such as their central densities, reveals differences that may hint at different formation mechanisms."
+ Past studies of Cvg OB? have clearly been complicated by an uncertain star formation history aud a laree spatial extent in a crowded region of the Galactic Plane., Past studies of Cyg OB2 have clearly been complicated by an uncertain star formation history and a large spatial extent in a crowded region of the Galactic Plane.
+ These are likely to be inherent properües of massive star forming regions in galactie spiral armis where star formation continues for longer periods and across larger areas. propagated through entire GAICs by continued rejuvenation and feedback (risgering.," These are likely to be inherent properties of massive star forming regions in galactic spiral arms where star formation continues for longer periods and across larger areas, propagated through entire GMCs by continued rejuvenation and feedback triggering."
+ These findings support the need to fully understand the full star formation history of a region before if its observed properties are to be correctly interpreted., These findings support the need to fully understand the full star formation history of a region before if its observed properties are to be correctly interpreted.
+ Cvgnus OB? is clearly a very important region to study. nol onlv because of its proximily and size. bul also because of the insights it will give us into ealactic-scale star formation and (the evolution of GMCs.," Cygnus OB2 is clearly a very important region to study, not only because of its proximity and size, but also because of the insights it will give us into galactic-scale star formation and the evolution of GMCs."
+ We thank Phil Lucas and Yvonne Unruh for careful reading of this paper ancl helpful discussions., We thank Phil Lucas and Yvonne Unruh for careful reading of this paper and helpful discussions.
+ We also thank the anonvmous referee for helpful suggestions Chat have improved (his work., We also thank the anonymous referee for helpful suggestions that have improved this work.
+ This research has made use of data from theObservatory (operated bv the Smithsonian Astrophysical Observatory on behalf of NASA) obtained. [rom the Data Archive., This research has made use of data from the (operated by the Smithsonian Astrophysical Observatory on behalf of NASA) obtained from the Data Archive.
+ This publication makes use of data products from IPILAS (carried oul al the Isaac Newton Telescope. INT. operated on the island of La Palma by the Isaac Newton Group). 2MASS (a joint project of the University of Massachusetts and the Infrared Processing and Analvsis Center / California Institute of Technology. funded by NASA and ihe NSF). and UKIDSS (supported by the UIKATC and CASU).," This publication makes use of data products from IPHAS (carried out at the Isaac Newton Telescope, INT, operated on the island of La Palma by the Isaac Newton Group), 2MASS (a joint project of the University of Massachusetts and the Infrared Processing and Analysis Center / California Institute of Technology, funded by NASA and the NSF), and UKIDSS (supported by the UKATC and CASU)."
+ JJD was funded by NASA, JJD was funded by NASA
+Two degeneracies are suggested by these considerations.,Two degeneracies are suggested by these considerations.
+ Iu the first. illustrated by solutious 220a aud 220a. the present distauce to the third body is substantially larger in 220a.," In the first, illustrated by solutions 220a and $^\ast$, the present distance to the third body is substantially larger in 220a."
+ That is because the third body had larger initial velocity away from the LG., That is because the third body had larger initial velocity away from the LG.
+ The present. tidal field ou the LG is stronger in 220a*. but that does not much matter because in these solutious the tidal field at low redshift has little ellect.," The present tidal field on the LG is stronger in $^\ast$, but that does not much matter because in these solutions the tidal field at low redshift has little effect."
+ The freedom to adjust the present distance of the third body has uot been explored in other solutious., The freedom to adjust the present distance of the third body has not been explored in other solutions.
+ A second degeneracy that has not been explored at all is the relation between the third body mass AM and its iuitial distauce ry from the LC., A second degeneracy that has not been explored at all is the relation between the third body mass $M_3$ and its initial distance $r_3$ from the LG.
+ The condition of near balance of gravity aud the counter term for the third body is approximately GCMpc+Ma)/r7Cpra. where p is the mean mass deusity in the backgrouud homogeneous cosmology.," The condition of near balance of gravity and the counter term for the third body is approximately $G(M_{\rm LG}+M_3)/r_3^2 \sim G\rho r_3$, where $\rho$ is the mean mass density in the background homogeneous cosmology."
+" The tidal field ou the LC. scales as AL,r3 +Ma).", The tidal field on the LG scales as M_3/r_3^3 +M_3).
+ It the third body is inore massive than the LG the dynamical iuodel is expected to allow the mass of the third body to scale with its initial distauce as , If the third body is more massive than the LG the dynamical model is expected to allow the mass of the third body to scale with its initial distance as r_3^3.
+This may eureo in the assessment of how the third massive body in this aualvsis might5 be related to the present clistribution of mass around the LG., This may figure in the assessment of how the third massive body in this analysis might be related to the present distribution of mass around the LG.
+ Figure 2. compares orbits iu soluious 220a aud 220b with the initial couditious shown iu Figure 1.., Figure \ref{figure:2} compares orbits in solutions 220a and 220b with the initial conditions shown in Figure \ref{figure:1}.
+ The comoving coordi1ales a'e centered on the center of inass of the four-bocly system auc oriented as iu equation (3.1) ., The comoving coordinates are centered on the center of mass of the four-body system and oriented as in equation \ref{eq:LMCvelocity}) ).
+" The ¢istinet dilfereuce between the solutions is seeu [or example iu the position of the massive t""rd bocv. plotted as the dot dash-line. well to the right of the LG in the w-cirection in 220a. anc well to te left in 220b (iu the first aud fourth panels in the figure)."," The distinct difference between the solutions is seen for example in the position of the massive third body, plotted as the dot dash-line, well to the right of the LG in the $x$ -direction in 220a, and well to the left in 220b (in the first and fourth panels in the figure)."
+" But one may also notice the sit""larity «X relative motions within the LC.", But one may also notice the similarity of relative motions within the LG.
+ The stinilarity of solutions within te LG is more clearly seen iu Figure 3.., The similarity of solutions within the LG is more clearly seen in Figure \ref{figure:3}.
+ In the panel ou the left the comoving coordinates a'e referred to the ceuter of mass of the three LG galaxies. but with the offset that puts the present positio of the MW at the origin.," In the panel on the left the comoving coordinates are referred to the center of mass of the three LG galaxies, but with the offset that puts the present position of the MW at the origin."
+ The pauel on the right slows paths of the heliocentric angular position ofthe LMC in ealactic coordinates., The panel on the right shows paths of the heliocentric angular position of the LMC in galactic coordinates.
+ The paths in both, The paths in both
+Tt is not known what the exact temperature profile of the tube is.,It is not known what the exact temperature profile of the tube is.
+ ΔΕΑ ealeulated the thermal structure of the eas funnel. but obtained values for the temperature (hat were too low to explain the observations.," \citeauthor{martin} calculated the thermal structure of the gas funnel, but obtained values for the temperature that were too low to explain the observations."
+ WIC adopted an ad hoc temperature profile that was able to reproduce the observations., \citeauthor{hart..94} adopted an ad hoc temperature profile that was able to reproduce the observations.
+ In the absence of a reliable temperature estimate Lor (he tube. we assume here a temperature profile similar to that of ILLIC. shown in Figure 2..," In the absence of a reliable temperature estimate for the tube, we assume here a temperature profile similar to that of \citeauthor{hart..94}, shown in Figure \ref{fig2}."
+ The mean temperature is ~8000 Ix. The temperature does not vary greatly inside (he tube. cropping near (he disk.," The mean temperature is $\sim 8000$ K. The temperature does not vary greatly inside the tube, dropping near the disk."
+ As we show in the following sections. Alfvenic heating does not have a strong dependence on temperature and. thus. (he exact profile is not of any great importance.," As we show in the following sections, Alfvenic heating does not have a strong dependence on temperature and, thus, the exact profile is not of any great importance."
+ In Table 2.. we show the magnetospheric model parameters adopted in our calculations. which arethe same as those used by LIC.," In Table \ref{tab1}, we show the magnetospheric model parameters adopted in our calculations, which arethe same as those used by \citeauthor{hart..94}. ."
+terms of observed 0.3.10 keV luminosity as a traction of the Eddineton limit.,terms of observed 0.3–10 keV luminosity as a fraction of the Eddington limit.
+ They (see also Barnardetal. 2003)) jwe applied this to identify black hole candidates in M3I., They (see also \citealt{Bea2003}) ) have applied this to identify black hole candidates in M31.
+ We have performed a similar PDS analvsis ofULX-1.. both from the 2000 March observation and from the 2004 July observations relpdslong)).," We have performed a similar PDS analysis of, both from the 2000 March observation and from the 2004 July observations \\ref{pdslong}) )."
+ The PDSs can be described bx a power law with a slope of about. —1. merging into the photon counting noise at 7 IIz.," The PDSs can be described by a power law with a slope of about $-$ 1, merging into the photon counting noise at $^{-3}$ Hz."
+ This is consistent with a Type D PDS for an extragalactic X-ray. binary. implying (hat the Edcdington accretion rate of (he source was in the range 0.1.1.0 both in 2000 March and in 2004 July.," This is consistent with a Type B PDS for an extragalactic X-ray binary, implying that the Eddington accretion rate of the source was in the range 0.1–1.0 both in 2000 March and in 2004 July."
+ Therefore. equating our estimate of the 0.310 keV luminosity in 2004 July of ~6x10 ((83.3) with 0.1 Liga. we infer the upper limit of compact object mass of ~ 40 M...," Therefore, equating our estimate of the 0.3–10 keV luminosity in 2004 July of $\sim 6 \times 10^{38}$ 3.3) with 0.1 $_{\rm Edd}$, we infer the upper limit of compact object mass of $\sim$ 40 $_\sun$."
+ If the X-ravs we observe from are beamecl. (hen the (rue Iuminosity is lower. aud the compact object mass is also lower.," If the X-rays we observe from are beamed, then the true luminosity is lower, and the compact object mass is also lower."
+ On the other hand. equating the highest luminosity we measure (3xI0 1)) with Lia. we infer a lower limit of 20 MN... again assuming (hat there is no beaming and sslavs sub-Ecddington at all times.," On the other hand, equating the highest luminosity we measure $\times 10^{39}$ ) with $_{\rm Edd}$, we infer a lower limit of 20 $_\sun$, again assuming that there is no beaming and stays sub-Eddington at all times."
+ However. the latter assumption need not apply strictly. since jb is quite possible for an accreting object (o exceed its Eddineton limit for a short (ime.," However, the latter assumption need not apply strictly, since it is quite possible for an accreting object to exceed its Eddington limit for a short time."
+ We have investigated the high state spectra ofULX-1.. fist using an absorbed blackbody model (Mukaietal.2003:Kone2004).," We have investigated the high state spectra of, first using an absorbed blackbody model \citep{Mea2003,KDY2004}."
+. When applied to the 2004 July data. we obtain an high absorbingcolumn (Nyj~2x107! 7)) and an absorbed 0.3.7 keV huninosity of ~4x109AN|. (," When applied to the 2004 July data, we obtain an high absorbingcolumn $\sim 2 \times 10^{21}$ ) and an absorbed 0.3–7 keV luminosity of $\sim 4 \times 10^{38}$. ("
+"Note that the 3x10!"" vvalue claimed by Nongetal.(2004) is the unabsorbed luminosity in the 0.3.1 keV band).",Note that the $\times 10^{40}$ value claimed by \citet{KDY2004} is the unabsorbed luminosity in the 0.3–7 keV band).
+" The A-rayv absorbing column inferred. by this fit is well above the Galactic column of ~1010 aand the total absorption through the disk of MIOLI (~6x10°"" 7)).", The X-ray absorbing column inferred by this fit is well above the Galactic column of $\sim 10^{20}$ and the total absorption through the disk of M101 $\sim 6 \times 10^{20}$ ).
+ Η this X-ray column was due to normal interstellar matter. this translates to an expected color excess Egν of ~0.33 (Bohlinetal.1978).. or a visual extinction Ay of ~1 magnitude.," If this X-ray column was due to normal interstellar matter, this translates to an expected color excess $_{B-V}$ of $\sim$ 0.33 \citep{Bea1978}, or a visual extinction $_V$ of $\sim$ 1 magnitude."
+ In contrast. the optical counterpart has a D—V color of —0.15z0.12. which is consistent with that of a D supereiant wilh little reddening (see Figure 4 of Ixuntzetal. 2005)).," In contrast, the optical counterpart has a $-$ V color of $-$ $\pm$ 0.12, which is consistent with that of a B supergiant with little reddening (see Figure 4 of \citealt{Kea2005}) )."
+ In fact. an assumed value of Ay=0.4 will make (he optical counterpart intrinsically bluer than any early (vpe stars.," In fact, an assumed value of $_V$ =0.4 will make the optical counterpart intrinsically bluer than any early type stars."
+ Moreover. there is a suggestion of clay-to-day variability of dduring (he interval July 512. 2004 in the spectral fits by Noneetal. (2004)..," Moreover, there is a suggestion of day-to-day variability of during the interval July 5–12, 2004 in the spectral fits by \citet{KDY2004}. ."
+ Therefore. if this model is correct. a significant. fraction of the X-ray absorber is likely to be located," Therefore, if this model is correct, a significant fraction of the X-ray absorber is likely to be located"
+ up s The svstem of equations (2)). (3.1)). (:2)A). »)). acdinits the following self-similar solution:Per qo," ^2 c_s The system of equations \ref{1}) ), \ref{2r}) ), \ref{3r}) ), \ref{45r}) ), admits the following self-similar solution:, d^0."
+", ]lere the constant factors are: and vy and Qy are given by equations (32)). (3: ))."," Here the constant factors are: and $v_0$ and $\Omega_0$ are given by equations \ref{v1}) ), \ref{Omega1}) )."
+ llaving derived (he self-similar solutions. we now prove that (he assumptions mace in order to simplify equations (4)) ancl (7)) are consistent.," Having derived the self-similar solutions, we now prove that the assumptions made in order to simplify equations \ref{2}) ) and \ref{3}) ) are consistent."
+ Let us first consider the angular momentum equation (7))., Let us first consider the angular momentum equation \ref{3}) ).
+ Because pez=pq —const., Because $\rho c_s^2=p_{\rm out}=$ const.
+ in the boundary laver and angular momentum flix J —const., in the boundary layer and angular momentum flux $\dot J=$ const.
+ throughout the entire accretion, throughout the entire accretion
+there is a better understaudiug of how best to treat attnosphere-tuterior coupling. we will want to extend the formalism preseuted here to include exchange terms.,"there is a better understanding of how best to treat atmosphere-interior coupling, we will want to extend the formalism presented here to include exchange terms."
+ Although may drag mechanisms could be at work iu a planets atinosphere. here we review three main pliysica processes that have bee1 ideitified as possibly produciug signiicant drag on lot Jupiters: vertical sjear instabilities (Cioodiran2009:Li&Coodman2010). shocss etal.2010:Rauscher&Menou 2010).. aud maeetic drag (Pernaetal.20102).," Although many drag mechanisms could be at work in a planet's atmosphere, here we review three main physical processes that have been identified as possibly producing significant drag on hot Jupiters: vertical shear instabilities \citep{Goodman,Li2010}, shocks \citep{DobbsDixon2010,RM10}, and magnetic drag \citep{Perna2010a}."
+. While mumeric‘al simulatious are able ο Caxture large-scaleinstabilities?.. (1e computationa limits on resolution for global circulation moclels could elininate any ---istabilities that are triggerecMD on sub-grid scales.," While numerical simulations are able to capture large-scale, the computational limits on resolution for global circulation models could eliminate any instabilities that are triggered on sub-grid scales."
+ Aclditionally. mode stlal asstme hydrostatie equilibrium have 1ο formal vertica acceleration (vertical veOCILY exists. Ut results [rom the continuity eclallol) aid as such. these moclels are uuable to capture vertical shear istabilities. which nu be parameterized.," Additionally, models that assume hydrostatic equilibrium have no formal vertical acceleration (vertical velocity exists, but results from the continuity equation) and as such, these models are unable to capture vertical shear instabilities, which must be parameterized."
+ As al attempt to determine wrat kind of stb-erid parameterization nee ye iucluded in bot Jupiter eloyal circulation moclels. Li&Goodau(2010) perforin a linear analysis auc non-linear sunulatLOS ol lunited regious of a |ot Jupiter atmosphere (two-dimensional. wi inite exteuts in longitcde al altitude).," As an attempt to determine what kind of sub-grid parameterization needs to be included in hot Jupiter global circulation models, \citet{Li2010} perform a linear analysis and non-linear simulations of limited regions of a hot Jupiter atmosphere (two-dimensional, with finite extents in longitude and altitude)."
+ I1 their siuittlatious they see weak trausieut shocks. b he main source of kinetic euerey dissipation is through recurriiο lxelviu-Helmholtz instabilities.," In their simulations they see weak transient shocks, but the main source of kinetic energy dissipation is through recurring Kelvin-Helmholtz instabilities."
+ Although there is not vet lear prescripion lor how to do so. it will be important to determiue how to connect. processes in localized anc elobal simulations to correctly trea the full range of atiuospherie scales on hot Jupiters.," Although there is not yet a clear prescription for how to do so, it will be important to determine how to connect processes in localized and global simulations to correctly treat the full range of atmospheric scales on hot Jupiters."
+" ""Trausonic low is present in lost iumnerical siijiulatious of hot Jupiter atinospheres.", Transonic flow is present in most numerical simulations of hot Jupiter atmospheres.
+ While transient shocks may exist iu the atmosphere. many 11oclels contain a shock (or Lydratlic jump?)) as a steady feature (Dobbs-Dixon&Lju2008:Showmanetal.2008.2009:Dobbs-Di2010:Rauscher&Menou 2010).," While transient shocks may exist in the atmosphere, many models contain a shock (or hydraulic ) as a steady feature \citep{DobbsDixon2008,Showman2008,S09,DobbsDixon2010,RM10}."
+". This is a localized 'eejon of strong dissipation aud ieatiug that has a markedly different. spatia st""ucture than the radiative heatiug patteru. often ap»eariug on the night side of the planet. a fac tha μιαν strouey allect the APE generatio1."," This is a localized region of strong dissipation and heating that has a markedly different spatial structure than the radiative heating pattern, often appearing on the night side of the planet, a fact that may strongly affect the APE generation."
+ Tje primitive equation models of Showmaneal.(2008.2009) aud Rauscher&Menou(2010) do not iuclude sliock-capturiug schemes ancl te |eating associated with this feature is Cousistert withl adiabatic heating of the downward [low hat is indced by the stroug couvergeuce. ratle ‘than from any," The primitive equation models of \citet{Showman2008,S09} and \citet{RM10} do not include shock-capturing schemes and the heating associated with this feature is consistent with adiabatic heating of the downward flow that is induced by the strong convergence, rather than from any"
+are indeed the most abundant perturber.,are indeed the most abundant perturber.
+ Thus. our discussion below will predominantly focus on the 2000-8000 K data.," Thus, our discussion below will predominantly focus on the 2000–8000 K data."
+ Although the data for inverse processes are related via the detailed-balance relation and thus redundant. little space can be saved by omitting one of the two and so we present data for both. thus saving the end user the need to make such calculations themselves.," Although the data for inverse processes are related via the detailed-balance relation and thus redundant, little space can be saved by omitting one of the two and so we present data for both, thus saving the end user the need to make such calculations themselves."
+ lt is worth noting that the excitation and 10n-pair. production processes were calculated directly from the cross-sections. while dexcitation and mutual-neutralisation processes were calculated via the detailed-balance relation rather than direct calculation from the deexcitation cross-sections.," It is worth noting that the excitation and ion-pair production processes were calculated directly from the cross-sections, while dexcitation and mutual-neutralisation processes were calculated via the detailed-balance relation rather than direct calculation from the deexcitation cross-sections."
+ As for Lt (2).. for the temperatures most relevat for cool stars around 6000 K. it is readily seen that the rate coefficients for 1on-pair production are remarkably large. particularly from the first excited 5 state. here +s.," As for Li \citep{2003A&A...409L...1B}, for the temperatures most relevant for cool stars around 6000 K, it is readily seen that the rate coefficients for ion-pair production are remarkably large, particularly from the first excited $s$ state, here $s$."
+ Some rates between close-togetherneighbouring levels. such as 4s3d and 4d-4f are also seen to be very large.," Some rates between close-together neighbouring levels, such as $4s \rightarrow 3d$ and $4d \rightarrow 4f$ are also seen to be very large."
+ At low temperatures. the rate coefficients for these transitions. with small energy thresholds become by far the largest.," At low temperatures, the rate coefficients for these transitions with small energy thresholds become by far the largest."
+ In? it was shown that narrow orbiting and Feshbach resonances are often present in the cross-sections at low energy (see Figs., In \cite{2010PhRvA..81c2706B} it was shown that narrow orbiting and Feshbach resonances are often present in the cross-sections at low energy (see Figs.
+ 5. 8 and 9 of that paper).," 5, 8 and 9 of that paper)."
+ To fully resolve these features requires calculations with very small steps in collision energy. which is extremely time consuming.," To fully resolve these features requires calculations with very small steps in collision energy, which is extremely time consuming."
+ Such detailed calculations were performed for the 3s—3p cross-sections in order to resolve a significant fraction of the orbiting resonances., Such detailed calculations were performed for the $3s\rightarrow 3p$ cross-sections in order to resolve a significant fraction of the orbiting resonances.
+ This further allowed pure background cross-sections. Le. with no resonances. to be constructed for this transition.," This further allowed pure background cross-sections, i.e. with no resonances, to be constructed for this transition."
+ From these calculations it was estimated that the contribution of the resonances to the rate coefficient was roughly at 500 K. at 2000 K. at 5000 K and at 8000 K. The rate coefficients for 3s—3p presented here are calculated including the contribution of orbiting resonances.," From these calculations it was estimated that the contribution of the resonances to the rate coefficient was roughly at 500 K, at 2000 K, at 5000 K and at 8000 K. The rate coefficients for $3s\rightarrow 3p$ presented here are calculated including the contribution of orbiting resonances."
+ In all other cases. the cross-section calculations were performed using an energy grid with steps much greater than the widths of the resonances.," In all other cases, the cross-section calculations were performed using an energy grid with steps much greater than the widths of the resonances."
+ Consequently most resonances are missed and those resonances that are by chance sampled are not resolved., Consequently most resonances are missed and those resonances that are by chance sampled are not resolved.
+ Missed resonances. will lead to an underestimate of the contribution of the rate coefficient. while unresolved resoances will generally lead to an overestimate of the contribution since the widths of the resonances are significantly overestimated. and this may lead to some cancellation of errors.," Missed resonances will lead to an underestimate of the contribution of the rate coefficient, while unresolved resoances will generally lead to an overestimate of the contribution since the widths of the resonances are significantly overestimated, and this may lead to some cancellation of errors."
+ In any case. the error in the rate coefficients from this source of uncertainty should not exceed the magnitude of the corrections giver above for 3s—3p. 1.9. of order10%.. which ts significantly less than the uncertainties from other sources discussed below.," In any case, the error in the rate coefficients from this source of uncertainty should not exceed the magnitude of the corrections given above for $3s\rightarrow 3p$, i.e. of order, which is significantly less than the uncertainties from other sources discussed below."
+ The accuracy of these rate coefficients is determined by the accuracy of the cross-sections near the threshold., The accuracy of these rate coefficients is determined by the accuracy of the cross-sections near the threshold.
+ At 7=2000 K the values of the rate coetticients vary from 10731 to 107? em?/s. more than 20 orders of magnitude.," At $T=2000$ K the values of the rate coefficients vary from $10^{-31}$ to $10^{-9}$ $^3/$ s, more than 20 orders of magnitude."
+ At 8000 K the variation is from 1077 to 107? eni)/s. 13 orders of magnitude.," At $T=8000$ K the variation is from $10^{-22}$ to $10^{-9}$ $^3/$ s, 13 orders of magnitude."
+ This reflects the fact that at the collision energies of interest the nonadiabatie transition probabilities also vary by similar factors. and that many of them have extremely small values.," This reflects the fact that at the collision energies of interest the nonadiabatic transition probabilities also vary by similar factors, and that many of them have extremely small values."
+ These small transition probabilities are sensitive to variations in both the input quantum chemical data (potentials and couplings and the method employed for solving the coupled channel equations in the quantum scattering In addition. transitions can occur in symmetries other than !X: for example. in the singlet system the p-. d-. and f-states can populate the ionic state via rotational couplings between [I and X states. as well as excite other states due to both rotational and radial nonadiabatic. couplings m Π. A. and other symmetries.," These small transition probabilities are sensitive to variations in both the input quantum chemical data (potentials and couplings) and the method employed for solving the coupled channel equations in the quantum scattering In addition, transitions can occur in symmetries other than $^1 \Sigma^+$: for example, in the singlet system the $p$ -, $d$ -, and $f$ -states can populate the ionic state via rotational couplings between $\Pi$ and $\Sigma$ states, as well as excite other states due to both rotational and radial nonadiabatic couplings in $\Pi$, $\Delta$, and other symmetries."
+ The same is true for the triplet system with larger (at least by a factor of 3) initial statistical probabilities., The same is true for the triplet system with larger (at least by a factor of 3) initial statistical probabilities.
+ For example. the cross-sections for 3s-3p excitation and their sensitivity to the input quantum chemical data are discussed at length in ?.. where we have shown that the results may vary by a factor of about 100.," For example, the cross-sections for $3s \rightarrow 3p$ excitation and their sensitivity to the input quantum chemical data are discussed at length in \cite{2010PhRvA..81c2706B}, where we have shown that the results may vary by a factor of about 100."
+" Similarly. from our investigations 1t is possible to make estimates of how much rate coefficients might vary for all transitions. and thus we define a ""fluctuation factor"" representing an estimate of this variation and which ts somewhat analogous to an estimate of the uncertainty."," Similarly, from our investigations it is possible to make estimates of how much rate coefficients might vary for all transitions, and thus we define a “fluctuation factor” representing an estimate of this variation and which is somewhat analogous to an estimate of the uncertainty."
+ The probable range for a given rate coefficient can be estimated to be from the value given in Table | multiplied by the minimum value of the fluctuation factor. up to the same value multiplied by the maximum value of the fluctuation factor. noting that from the arguments discussed above. particularly contributions from other symmetries. 1t is generally more likely that the rate coefficient would be increased rather than decreased.," The probable range for a given rate coefficient can be estimated to be from the value given in Table \ref{tab:rates} multiplied by the minimum value of the fluctuation factor, up to the same value multiplied by the maximum value of the fluctuation factor, noting that from the arguments discussed above, particularly contributions from other symmetries, it is generally more likely that the rate coefficient would be increased rather than decreased."
+ We have assembled estimates of these fluctuation factors for all transitions in Table 2.., We have assembled estimates of these fluctuation factors for all transitions in Table \ref{tab:errors}.
+ A commol feature of the quantum scattering calculations. particularly for very adiabatic transitions. is that the sensitivity of the cross-sections to the input quantum chemistry data usually decreases with increasing collision energy. and thus cross-sections are often most uncertai very close to threshold. becoming less uncertain with increasing collision energy.," A common feature of the quantum scattering calculations, particularly for very adiabatic transitions, is that the sensitivity of the cross-sections to the input quantum chemistry data usually decreases with increasing collision energy, and thus cross-sections are often most uncertain very close to threshold, becoming less uncertain with increasing collision energy."
+ This means that rate coefficients are usually most uncertain at low temperatures. becoming less uncertain at higher temperatures.," This means that rate coefficients are usually most uncertain at low temperatures, becoming less uncertain at higher temperatures."
+ Thus. in some cases we provide a range of fluctuation factors corresponding to the lower and higher ends of the temperature range considered.," Thus, in some cases we provide a range of fluctuation factors corresponding to the lower and higher ends of the temperature range considered."
+ The rate coefficient for 3s—3p excitation is roughly 10! em?/s at 2000 K. somewhere in the middle of the range for all transitions.," The rate coefficient for $3s \rightarrow 3p$ excitation is roughly $10^{-21}$ $^3/$ s at 2000 K, somewhere in the middle of the range for all transitions."
+ Thus. it is quite natural to expect that fluctuation. factors for transitions. with. small. transition probabilties can be much larger than 100. while the fluctuation coefficients for the transitions determined by large transition probabilities should be close to unity.," Thus, it is quite natural to expect that fluctuation factors for transitions with small transition probabilties can be much larger than 100, while the fluctuation coefficients for the transitions determined by large transition probabilities should be close to unity."
+ It is worth making clear that the smallest rate coefhcients usually have the largest fluctuation factors: even if these rates were multiplied by such a large fluctuation factor they would not become dominant or substantially change the general pattern of relative strengths between different transitions., It is worth making clear that the smallest rate coefficients usually have the largest fluctuation factors: even if these rates were multiplied by such a large fluctuation factor they would not become dominant or substantially change the general pattern of relative strengths between different transitions.
+ Nevertheless. for astrophysical applications 1t," Nevertheless, for astrophysical applications it"
+Ry2.4 kpc even though the base model has Ry=2.0 kpe.,$\rd = 2.4$ kpc even though the base model has $\rd = 2.0$ kpc.
+ This (vpe of variation in wwith the fitted racial range is commonly found in external galaxies. and may go some wav lo explaining the variation in reported scale lengths for the Milkv. Way (Sackett 1997).," This type of variation in with the fitted radial range is commonly found in external galaxies, and may go some way to explaining the variation in reported scale lengths for the Milky Way (Sackett 1997)."
+ Clearly MOND prefers a compact Milky Was. consistent with the CODE data (Dinnev 11997: Drimunel Spergel 2001) and the rather high observed microlensing optical depth (Popowski 22005) that requires (hat barvonic mass dominate within the solar circle (Binney Evans 2001: Dissantz. Debattista. Gerhard 2004).," Clearly MOND prefers a compact Milky Way, consistent with the COBE data (Binney 1997; Drimmel Spergel 2001) and the rather high observed microlensing optical depth (Popowski 2005) that requires that baryonic mass dominate within the solar circle (Binney Evans 2001; Bissantz, Debattista, Gerhard 2004)."
+ A related test of MOND in the Alilky Way is provided by the vertical support of the stellar ancl gas disks., A related test of MOND in the Milky Way is provided by the vertical support of the stellar and gas disks.
+ In the outskirts of the galaxy. the barvonie components shoulel [lare substantially as Che effective mass scale height comes to be dominated by the quasi-spherical dark matter halo.," In the outskirts of the galaxy, the baryonic components should flare substantially as the effective mass scale height comes to be dominated by the quasi-spherical dark matter halo."
+ In contrast. the potential remains disk-like close to the plane in MOND.," In contrast, the potential remains disk-like close to the plane in MOND."
+ The net effect is (hat MONDian disks will be somewhat thinner. all other things being equal.," The net effect is that MONDian disks will be somewhat thinner, all other things being equal."
+ Recently. Sannchez-Saleedo. Saha. Naravan (2008) have analyzed the thickness of the eas laver of the Milky Wav in the context of MOND.," Recently, Sánnchez-Salcedo, Saha, Narayan (2008) have analyzed the thickness of the gas layer of the Milky Way in the context of MOND."
+ Thev find that MONDian sell-eravity of the disk provides a plausible explanation of the thickness and flaring of the gas laver., They find that MONDian self-gravity of the disk provides a plausible explanation of the thickness and flaring of the gas layer.
+ Ailegrom (2008) points out that the form of the interpolation function as well as (he mass model matters to the details of the support of the gas laver., Milgrom (2008) points out that the form of the interpolation function as well as the mass model matters to the details of the support of the gas layer.
+ It also seems possible that magnetic field support may be non-negligible. so a perfect fit might be difficult to obtain.," It also seems possible that magnetic field support may be non-negligible, so a perfect fit might be difficult to obtain."
+" In comparison. it is necessary {ο invoke a massive (>LOM!AA, ) ring-like dark matter component in addition to the traditional quasi-spherical halo in order to explain (he gas Chickuess in conventional terms (Ixalberla 22007)."," In comparison, it is necessary to invoke a massive $> 10^{11}\;\Msun$ ) ring-like dark matter component in addition to the traditional quasi-spherical halo in order to explain the gas thickness in conventional terms (Kalberla 2007)."
+ For stus. MOND will support a higher vertical velocity dispersion at a given disk thickness.," For stars, MOND will support a higher vertical velocity dispersion at a given disk thickness."
+ This effect in the Milkv. Way is rather subtle until large radii (Nipoti 22007). but might conceivably be detectable with GALA.," This effect in the Milky Way is rather subtle until large radii (Nipoti 2007), but might conceivably be detectable with GAIA."
+ The large velocity dispersions of planetary. nebulae al large radii in lace-on spirals might be an indication of such an ellect (Hlerrmann Ciardullo 2005)., The large velocity dispersions of planetary nebulae at large radii in face-on spirals might be an indication of such an effect (Herrmann Ciardullo 2005).
+ Treating the Milkv Way as we would an external galaxy. il is possible to obtain the rotation curve from the surface density with MOND.," Treating the Milky Way as we would an external galaxy, it is possible to obtain the rotation curve from the surface density with MOND."
+ There is no freedom to adjust the nmass-to-light ratio as in external galaxies., There is no freedom to adjust the mass-to-light ratio as in external galaxies.
+ The result is satisfactory provided (he scale length ol the Milkv. Way is relatively short (2 2.5 κρο). as implied by the CODE data (Binnev," The result is satisfactory provided the scale length of the Milky Way is relatively short (2 — 2.5 kpc), as implied by the COBE data (Binney"
+indicates (hat the peak of the spectral energy. distribution is associated with electrons. αἱ 5—y Where the radiative cooling time equals (he escape lime from the svstem. and (hat the electron index p>2.,"indicates that the peak of the spectral energy distribution is associated with electrons at $\gamma=\gamma_b$ where the radiative cooling time equals the escape time from the system, and that the electron index $p \geq 2$."
+ If the EGRET--observed tix were due to Thomson scattering. the observed. photon index s=à+1 of a blazar would relate to the index p of the electron power law via p—2s1.," If the -observed flux were due to Thomson scattering, the observed photon index $s=\alpha+1$ of a blazar would relate to the index $p$ of the electron power law via $p=2s-1$."
+ However. since INN effects steepen the spectrum. application of this relation results in an electron index Chat is too steep. as can be seen in the lower left panel of figure 2.. where we plot the INN corrected photon index (solid line) as well as the relation s=(p+1)/2 (dashed line).," However, since KN effects steepen the spectrum, application of this relation results in an electron index that is too steep, as can be seen in the lower left panel of figure \ref{fig2}, where we plot the KN corrected photon index (solid line) as well as the relation $s=(p+1)/2$ (dashed line)."
+ In (his figure. (he seed photon energy is 64=5xI0. and the two-point spectral index is caleulated here and throughout the paper using (he Iuminositv of the model spectrum at 100 MeV and 10 GeV. The maximum electron Lorentz [actor is assumed to be 53D>L/eg. since the observations show no evidence of a cut-off.," In this figure, the seed photon energy is $\epsilon_0=5\times 10^{-5}$, and the two-point spectral index is calculated here and throughout the paper using the luminosity of the model spectrum at $100 $ MeV and $10 $ GeV. The maximum electron Lorentz factor is assumed to be $\gamma_2{\cal D}>1/\epsilon_0$, since the observations show no evidence of a cut-off."
+ Under this assumption the observed photon index is not a Iunction of 53D., Under this assumption the observed photon index is not a function of $\gamma_2{\cal D}$.
+ For photon indices s around 2.2 the electron index that would result from Thomson scattering is sleeper than the INN case by about 0.3., For photon indices $s$ around $2.2$ the electron index that would result from Thomson scattering is steeper than the KN case by about $0.3$.
+ In the lower right panel ol ligure 2. we show a histogram of the observed photon idices of all blazars in the third ecalalog (Ilartinan.etal.1999).. except for sources with errors greater than 0.3 in (heir photon index determination aud also excluding the three X-ray selected BL Lacertae (MINAS 421. MINN 501. PINS 2155-304). which. because (μον lack line emission in their spectra. are believed to be predominantly SSC emitters. We see that the distribution peaks al s22.3 (a careful calculation for all (he broad line emitting blazars gives s=2.20£0.05 (Mukherjee.etal. 1997))).," In the lower right panel of figure \ref{fig2} we show a histogram of the observed photon indices of all blazars in the third catalog \citep{hartman99}, except for sources with errors greater than $0.3$ in their photon index determination and also excluding the three X-ray selected BL Lacertae (MKN 421, MKN 501, PKS 2155-304), which, because they lack line emission in their spectra, are believed to be predominantly SSC emitters, We see that the distribution peaks at $s\approx 2.3$ (a careful calculation for all the broad line emitting blazars gives $s=2.20 \pm 0.05$ \citep{mukherjee97}) )."
+ For each of these sources we calculate the electron index p that would result in the observed photon index s. taking into account the correction due to the redshift of the source. and plot the results in (he upper left panel of figure 2..," For each of these sources we calculate the electron index $p$ that would result in the observed photon index $s$, taking into account the correction due to the redshift of the source, and plot the results in the upper left panel of figure \ref{fig2}. ."
+ The results cluster in a wide interval around p=3.2 that includes the value p=3.0., The results cluster in a wide interval around $p=3.2$ that includes the value $p=3.0$.
+ These values are expected from electron distributions with indices p=2.2 and p=2.0. respectively. that steepen by Ap=1 above ο.," These values are expected from electron distributions with indices $p=2.2$ and $p=2.0$, respectively, that steepen by $\Delta p=1$ above $\gamma_b$."
+ Since the rrange is al energies higher Chan the peak of the spectral energy distribution. the observed photon indices are compatible with acceleration at both relativistic (τςetal.—2000:Achterbergetal.2001). and. non-relativistic shocks. which predict p=2.23 and p=2 respectivelv [or the uncooled electron distributions.," Since the range is at energies higher than the peak of the spectral energy distribution, the observed photon indices are compatible with acceleration at both relativistic \citep{kirk00, achterbergetal01} and non-relativistic shocks, which predict $p=2.23$ and $p=2$ respectively for the uncooled electron distributions."
+ Unfortunately. however. the scatter in the observed indices is of the same order as the difference between the predictions. so that the data do not allow us to distinguish between these (wo cases.," Unfortunately, however, the scatter in the observed indices is of the same order as the difference between the predictions, so that the data do not allow us to distinguish between these two cases."
+ Simple electron cooling considerations predict a spectral break of Aa=0.5 in the transition beloreand alter the peak energy έρως OF the observed energy distribution. as a," Simple electron cooling considerations predict a spectral break of $\Delta\alpha=0.5$ in the transition beforeand after the peak energy $\epsilon_{peak}$ of the observed energy distribution, as a"
+the mid-IR luminous NLBRGs in our sample with the missing lobe-dominant quasar population (Fig.,the mid-IR luminous NLRGs in our sample with the missing lobe-dominant quasar population (Fig.
+ 8)., 8).
+" Their median core to lobe ratio is 2.=vl,(core. 5 Gllz)/vL,(1T3 MllIz)=5. while the median 2,=180 lor the quasars."," Their median core to lobe ratio is $R_\mathrm{c}=\nu L_\nu($ core, 5 $)/ \nu L_\nu(178$ $)=5$, while the median $R_\mathrm{c}=180$ for the quasars."
+ Conversely. all of the mil-I luminous sources with {ο>100 are classified as quasars.," Conversely, all of the mid-IR luminous sources with $R_\mathrm{c}>100$ are classified as quasars."
+ We will perform a more detailed statistical analysis when (he Spitzer observations of our [ull sample are complete., We will perform a more detailed statistical analysis when the Spitzer observations of our full sample are complete.
+ For most mid-IR. weak NLRGs. we reject the possibility that the AGN and radio core are viewed in an ‘off state while the radio lobes are still active.," For most mid-IR weak NLRGs, we reject the possibility that the AGN and radio core are viewed in an 'off' state while the radio lobes are still active."
+ A 5 Gllz core is detected in 13/17 of the mic-IR weak radio galaxies (Fig., A 5 GHz core is detected in 13/17 of the mid-IR weak radio galaxies (Fig.
+ 8) and 13/14 of the mid-IR. Iuminous ealaxies., 8) and 13/14 of the mid-IR luminous galaxies.
+ Furthermore. mid-IR. weak and mid-IB. luminous NLRGs have comparable core {ο total luminosity ratios of R.=1—100.," Furthermore, mid-IR weak and mid-IR luminous NLRGs have comparable core to total luminosity ratios of $R_\mathrm{c} = 1-100$."
+ Deeper 5 Gllz observations of the 5 non«detected cores (in ος 28. 153. 172. 315. and 319) will be necessary {ο determine whether or not they are in an olf state (e.g.. I2.« O.1).," Deeper 5 GHz observations of the 5 non-detected cores (in 3C 28, 153, 172, 315, and 319) will be necessary to determine whether or not they are in an off state (e.g., $R_\mathrm{c}<0.1$ )."
+ We find that FR II radio morphology. is not a reliable predictor of nuclear mid-Ih luminosity lor radio galaxies., We find that FR II radio morphology is not a reliable predictor of nuclear mid-IR luminosity for radio galaxies.
+humninosity. Unification theories must be modified to account for an additional population of mid-Ih weak radio galaxies., Unification theories must be modified to account for an additional population of mid-IR weak radio galaxies.
+ Both intrinsic jet power and interaction with the interstellar and intergalactic medium are likely to be important for determining radio morphology., Both intrinsic jet power and interaction with the interstellar and intergalactic medium are likely to be important for determining radio morphology.
+ The break Iuminosity between FR Land FR I radio sources is found to increase with host galaxy optical Iuminosity Bicknell 1996).," The break luminosity between FR I and FR II radio sources is found to increase with host galaxy optical luminosity \citep{ol94,b96}."
+. The existence of radio sources with hybrid FR ΕΕ morphology also points to the importance of environmental effects in determining radio morphology (Gopal-Ixrishna&Willa2000;Gawronskietal.2005).," The existence of radio sources with hybrid FR I/II morphology also points to the importance of environmental effects in determining radio morphology \citep{gkw00,gmk05}."
+. Furthermore. most ET I radio jets are one-sided. relativistic. and narrowly collimated on sub-parsec scales. just like FR Ls. ancl decollimate only on kpe scales (e.g.M87.Junor.Diretta.&Livio.1999).," Furthermore, most FR I radio jets are one-sided, relativistic, and narrowly collimated on sub-parsec scales, just like FR IIs, and decollimate only on kpc scales \citep[e.g. M87,][]{jbl99}."
+. Contrary (o a common nmisconception. not all FR I sources have radiativelv inefficient nuclei.," Contrary to a common misconception, not all FR I sources have radiatively inefficient nuclei."
+ For example. Centaurus A is persuasively argued to have a powerful hidden AGN (Whysong5&Antonucci2004).. and the BLRG 3C 120 is a well-known FR I source.," For example, Centaurus A is persuasively argued to have a powerful hidden AGN \citep{wa04}, and the BLRG 3C 120 is a well-known FR I source."
+ Deep VLA observations of the optically Iuminous. radio-quiet! quasar E13212-643 demonstrate that it has an FR I radio morphology (Dlundell&Rawlines2001).," Deep VLA observations of the optically luminous, 'radio-quiet' quasar E1821+643 demonstrate that it has an FR I radio morphology \citep{br01}."
+. These aud other cases (Antonneci2001) demonstrate that many powerful AGNs are ER Is., These and other cases \citep{ant01} demonstrate that many powerful AGNs are FR Is.
+ The observed variation in radio morphology and a wide range in AGN radio-Ioudness (Ixellernanetal.|1989) do not necessarily require a weak coupling between jet power and accretion power. but may demonstrate that multiple factors are at work.," The observed variation in radio morphology and a wide range in AGN radio-loudness \citep{kss89} do not necessarily require a weak coupling between jet power and accretion power, but may demonstrate that multiple factors are at work."
+ At least [ive parameters may be necessary to theoretically unify all AGN types: black hole mass. black hole," At least five parameters may be necessary to theoretically unify all AGN types: black hole mass, black hole"
+"where, p,V,B, and T are mass density, velocity, magnetic field, and temperature, respectively.","where, $\rho,\mathbf{V}, \mathbf{B}$, and $T$ are mass density, velocity, magnetic field, and temperature, respectively."
+" Pr indicates total pressure, Py+ where is gas pressure."," $P_T$ indicates total pressure, $P_g + B^2/8\pi$, where $P_g$ is gas pressure."
+" Total energy per unit B?/87,volume is describedP, as 6=pV?/2+P,/(y—1)B?/8r with specific heat ratio ΥΞ5/3.", Total energy per unit volume is described as ${\cal E} = \rho V^2/2 + P_g/(\gamma-1) + B^2/8\pi$ with specific heat ratio $\gamma = {5/3}$.
+" g is gravitational acceleration, —GMf£/r?, where G and M are the gravitational constant and solar mass, respectively."," $\mathbf{g}$ is gravitational acceleration, $-GM \hat{\mathbf{r}}/r^2$, where $G$ and $M$ are the gravitational constant and solar mass, respectively."
+ We adopt anisotropic thermal conduction tensor « along magnetic field lines (Yokoyama&Shi-bata 2001).., We adopt anisotropic thermal conduction tensor $\kappa$ along magnetic field lines \citep{yoko01}. .
+ Radiative cooling term R is assumed to be a function of local density and temperature (Suzuki&In-utsuk, Radiative cooling term $R$ is assumed to be a function of local density and temperature \citep{suzu05}.
+a Note that no other ad hoc source terms for heating are 2005)..included in the energy equation., Note that no other ad hoc source terms for heating are included in the energy equation.
+" The vector form of our basic equations are appropriately transformed into spherical coordinate system with (r,0= and Og=0."," The vector form of our basic equations are appropriately transformed into spherical coordinate system with $r,\theta=\pi/2,\phi$ ) and $\partial _\theta=0$."
+" Initial magnetic field is extrapolated using m/2,)potential field approximation so that the open field lines are extended from a kilo Gauss strength to 10 Gauss at the height of 2,000 km."," Initial magnetic field is extrapolated using potential field approximation so that the open field lines are extended from a kilo Gauss strength to 10 Gauss at the height of 2,000 km."
+" Initially, we set an isothermal K) atmosphere."," Initially, we set an isothermal $10^4$ K) atmosphere."
+ We input velocity disturbance in 6 (104direction at the footpoint of the flux tube in order to generate Alfvénn waves., We input velocity disturbance in $\theta$ direction at the footpoint of the flux tube in order to generate Alfvénn waves.
+" We assume white noise power spectrum with (1 km s-!)? in total power, which is rather a good approximation to the observed spectrum (Matsumoto&Kitai2010).."," We assume white noise power spectrum with (1 km $^{-1}$ $^2$ in total power, which is rather a good approximation to the observed spectrum \citep{mats10b}."
+" Since the velocity disturbances are pre-determined, any feedback effects (Grappinetal.2008) of the coronal disturbances on the photospheric motions are ignored."," Since the velocity disturbances are pre-determined, any feedback effects \citep{grap08} of the coronal disturbances on the photospheric motions are ignored."
+ Periodic boundary condition is posed in the @¢ direction., Periodic boundary condition is posed in the $\phi$ direction.
+ Our numerical simulation is based on HLLD scheme (Miyoshi&Kusano2005) that is robust and efficient among the various kind of approximate Riemann solvers., Our numerical simulation is based on HLLD scheme \citep{miyo05} that is robust and efficient among the various kind of approximate Riemann solvers.
+" The solenoidal condition ,V-B=0, is satisfied within a round-off error by using flux-CT method (Téth2000).."," The solenoidal condition $\nabla\cdot\mathbf{B}=0$, is satisfied within a round-off error by using flux-CT method \citep{toth00}."
+" The TVD-MUSCL scheme enables us to archive the second order accuracy in space, while the Runge-Kutta method gives the second order accuracy in time."," The TVD-MUSCL scheme enables us to archive the second order accuracy in space, while the Runge-Kutta method gives the second order accuracy in time."
+" We use min-mod limiter in order to suppress the numerical oscillation around shocks, which is one of the standard techniquein TVD-scheme."," We use min-mod limiter in order to suppress the numerical oscillation around shocks, which is one of the standard techniquein TVD-scheme."
+ Our, Our
+There is compelling evideuce for extensive mixiug in the envelopes of low mass evolved stars which is uot predicted to occur in Classical stellar models.,There is compelling evidence for extensive mixing in the envelopes of low mass evolved stars which is not predicted to occur in classical stellar models.
+ Rotation is frequently invoked. at least," Rotation is frequently invoked, at least"
+"Nova Oph 2009 (= V2672 Oph) was discovered by Ix. Hagaki on August 16.515 UL at mag 0.0 on untiltered CCD images (cf.Nakano.Yamaoka.&Ixadota2000)... al position (92000) a = 19.70. d = 26° 441A"".","Nova Oph 2009 (= V2672 Oph) was discovered by K. Itagaki on August 16.515 UT at mag 10.0 on unfiltered CCD images \citep[cf.][]{nakano}, at position (J2000) $\alpha$ = $^h$ $^m$ $^s$ .70, $\delta$ = $-$ $^\circ$ $^\prime$ $^{\prime\prime}$ .0."
+" This corresponds to Galactic 1738""coordinates {=00.021. b=|02.535. thus V2672 Oph lies very close to the dixtion to the Galactic centre."," This corresponds to Galactic coordinates $l$ =001.021, $b$ =+02.535, thus V2672 Oph lies very close to the direction to the Galactic centre."
+ Spectroscopic confirmation was obtained on August 17.6 UT bv Avanietal.(20))) and on August 17.8 UT by Munarietal.(2009).," Spectroscopic confirmation was obtained on August 17.6 UT by \citet{AMH09}
+ and on August 17.8 UT by \citet{MSO09}."
+. Both reported a very large velocity width for Lla and the pauciv of other emission lines on a red and featureless continuum., Both reported a very large velocity width for $\alpha$ and the paucity of other emission lines on a red and featureless continuum.
+ Munarietal.(2009) noted the great. strength. of interstellar lines. and. clilluse interstellar bands. which indicated a large reddening and therefore accounted. for the red slope of the continuum. of the nova.," \citet{MSO09} noted the great strength of interstellar lines and diffuse interstellar bands, which indicated a large reddening and therefore accounted for the red slope of the continuum of the nova."
+ Munarietal.(2009) also noted how the profile of the emission ines was highly structured. in the form of a broad trapezium with a shiuper Gaussian on top.," \citet{MSO09} also noted how the profile of the emission lines was highly structured, in the form of a broad trapezium with a sharper Gaussian on top."
+ The trapezium had a velocity widtl rol ~11700 aat the base anc ~6900 aab the top. while the sharp Gaussian component had a EAVILIAL of 14001.," The trapezium had a velocity width of $\sim$ 11700 at the base and $\sim$ 6900 at the top, while the sharp Gaussian component had a FWHM of $\sim$ 1400."
+ Thev also noted how the colours. the rapid decline. and the velxdty of the ejecta suggested that V2672 Oph was an outhurst occurring on a massive," They also noted how the colours, the rapid decline, and the velocity of the ejecta suggested that V2672 Oph was an outburst occurring on a massive"
+reveal the values of f. as they did for the first identification of the low-degree modes detected in whole-dise observations of the Sun.,"reveal the values of $\ell$, as they did for the first identification of the low-degree modes detected in whole-disc observations of the Sun."
+ A sixth frequency in the [986 data at —2.506 mHz was separated by //Hz from the fifth which. at a separation of iA. did not fit the asymptotie relation.," A sixth frequency in the 1986 data at $\sim 2.806\,$ mHz was separated by $\mu$ Hz from the fifth which, at a separation of $\frac{3}{4}\Delta\nu$, did not fit the asymptotic relation."
+ However. that relation is valid for linear adiabatic pulsations in spherically symmetric stars. whereas strong magnetic tields and chemical inhomogeneities in roAp stars significantly break the spherical symmetry.," However, that relation is valid for linear adiabatic pulsations in spherically symmetric stars, whereas strong magnetic fields and chemical inhomogeneities in roAp stars significantly break the spherical symmetry."
+ The effects of chemical inhomogeneities were examined by Balmforthetal.(2001)... but of greater interest to us here is the effect of the magnetic field.," The effects of chemical inhomogeneities were examined by \citet{Balmforth01}, but of greater interest to us here is the effect of the magnetic field."
+ To estimate this a singular perturbation approach has been applied to determine the modification to the frequency spectra (Dziembowski&Dziembowski2002:SaioGautschy 2004).," To estimate this a singular perturbation approach has been applied to determine the modification to the frequency spectra \citep{Dziembowski96, Bigot00, Cunha00, Bigot02, Saio04}."
+. Cunha&Gough(2000). found that at some frequencies the magneto-acoustical coupling is particularly strong. resulting in a loss of amplitude from the mode and a frequency shift of several (; Hz.," \citet{Cunha00} found that at some frequencies the magneto-acoustical coupling is particularly strong, resulting in a loss of amplitude from the mode and a frequency shift of several $\mu$ Hz."
+ Cunha(2001) suggested such a mechanism may be responsible for the unusual separation of the sixth frequency of the 1986 data. and predicted an additional mode 34srllz above the fifth frequency that had been missed by the observations of Kurtzetal.(1989).," \citet{Cunha01} suggested such a mechanism may be responsible for the unusual separation of the sixth frequency of the 1986 data, and predicted an additional mode $34\,\mu \mathrm{Hz}$ above the fifth frequency that had been missed by the observations of \citet{Kurtz89}."
+. More recently. Saio.Ryabchikova&Sachkov(2010). suggested that this mode might correspond to a moditied /=3 mode.," More recently, \citet{Saio10} suggested that this mode might correspond to a modified $\ell = 3$ mode."
+ 11217 was again observed photometrically in November and December 2000 as a part of the Whole Earth Telescope (WET) extended coverage campaign CXcov20). with telescopes from eight observatories with apertures ranging from 0.6 to 2.1m. in one of the highest-precision ground-based photometric studies ever undertaken (Kurtzetal.2005)..," 1217 was again observed photometrically in November and December 2000 as a part of the Whole Earth Telescope (WET) extended coverage campaign (Xcov20), with telescopes from eight observatories with apertures ranging from 0.6 to 2.1m, in one of the highest-precision ground-based photometric studies ever undertaken \citep{Kurtz05a}."
+ In a preliminary analysis of the data. the hypothesized “missing” mode was indeed found (Kurtzal. 2002).. with amplitude lower than those of the six modes that had been found already.," In a preliminary analysis of the data, the hypothesized `missing' mode was indeed found \citep{Kurtz02}, with amplitude lower than those of the six modes that had been found already."
+ From a more detailed analysis. two closely spaced modes at this frequency (Kurtzetal.2005) were found.," From a more detailed analysis, two closely spaced modes at this frequency \citep{Kurtz05a} were found."
+ This finding appeared to support the theory of Cunha&Gough (20003., This finding appeared to support the theory of \citet{Cunha00}.
+. Moreover. the greater damping expected from the theory might provide the explanation for the lower amplitude.," Moreover, the greater damping expected from the theory might provide the explanation for the lower amplitude."
+ However. Shibahashi(2003) expressed the concern that the magneto-acoustic coupling could result in damping so strong that the mode would orobably not be exeited. rendering the reason for the existence of the modes an open question.," However, \citet{Shibahashi03} expressed the concern that the magneto-acoustic coupling could result in damping so strong that the mode would probably not be excited, rendering the reason for the existence of the modes an open question."
+ However. no calculation of the excitation of the modes has yet been carried out to justify or allay hat concern. not even by a generalization of the simple. highly idealized. procedure adopted by Balmforth et al. (," However, no calculation of the excitation of the modes has yet been carried out to justify or allay that concern, not even by a generalization of the simple, highly idealized, procedure adopted by Balmforth et al. ("
+2001).,2001).
+ It must be asked why the “missing” modes were missing in he first place. given that they were found in the 2000 data with an amplitude greater than the noise level of the 1986 data.," It must be asked why the `missing' modes were missing in the first place, given that they were found in the 2000 data with an amplitude greater than the noise level of the 1986 data."
+ Kurtz(2005) noted that the amplitudes of the other modes had also varied between the two datasets by several tenths of à mmag. which is greater than the amplitude of the “missing” modes in the 2000 data.," \citet{Kurtz05a} noted that the amplitudes of the other modes had also varied between the two datasets by several tenths of a mmag, which is greater than the amplitude of the `missing' modes in the 2000 data."
+" Such a change in amplitude could well provide the explanation for why the ""missing modes were originally absent.", Such a change in amplitude could well provide the explanation for why the `missing' modes were originally absent.
+ An interesting observation made by Kurtzetal.(2005) was that there was little difference in the sum of power between the two data sets. indicating that. since the observed modes all have about the same inertia. the total pulsation energy may have been conserved. suggesting that nonlinear interactions between modes transfer energy on a timescale shorter than the intrinsic growth times.," An interesting observation made by \citet{Kurtz05a} was that there was little difference in the sum of power between the two data sets, indicating that, since the observed modes all have about the same inertia, the total pulsation energy may have been conserved, suggesting that nonlinear interactions between modes transfer energy on a timescale shorter than the intrinsic growth times."
+ Other roAp stars that have multiple pulsation modes show significant amplitude variations on timescales of a few days. 660435 being one example (Matthews.Wehlau&Kurtz 1987).," Other roAp stars that have multiple pulsation modes show significant amplitude variations on timescales of a few days, 60435 being one example \citep{Matthews87}."
+. On the other hand. pulsation amplitudes of singly periodic roAp stars. such as 33831. (Kurtzetal.1997)... appear to be very stable.," On the other hand, pulsation amplitudes of singly periodic roAp stars, such as 3831 \citep{Kurtz97}, appear to be very stable."
+ However. some roAp stars appear not to exhibit conservation of total pulsation energy.," However, some roAp stars appear not to exhibit conservation of total pulsation energy."
+ For example. power is clearly not being conserved in 2217522. in which a new Tequency appeared between data sets in 1982 and 1989 etal.1991).," For example, power is clearly not being conserved in 217522, in which a new frequency appeared between data sets in 1982 and 1989 \citep{Kreidl91}."
+". A unique aspect of the oscillations of roAp stars is that he pulsation amplitude is modulated with the rotation of the star, as is the apparent magnetic field strength."," A unique aspect of the oscillations of roAp stars is that the pulsation amplitude is modulated with the rotation of the star, as is the apparent magnetic field strength."
+ The modulation is well explained by the oblique pulsator model. in which the o)ulsations are axisymmetric oscillations with the same axis as the magnetic field and inclined to to the rotational axis (Kurtz1982).," The modulation is well explained by the oblique pulsator model, in which the pulsations are axisymmetric oscillations with the same axis as the magnetic field and inclined to to the rotational axis \citep{Kurtz82}."
+". Bagnuloetal.(1995). found the magnetic field in 11217 to be approximately dipolar. with a polar strength of 5,=3.9κ an angle /=137 between the line-of-sight and rotation axis. and angle ;j=150° between the rotation and magnetic axes."," \citet{Bagnulo95} found the magnetic field in 1217 to be approximately dipolar, with a polar strength of $B_p=3.9\,\mathrm{kG}$, an angle $i=137^{\circ}$ between the line-of-sight and rotation axis, and angle $\beta=150^{\circ}$ between the rotation and magnetic axes."
+" More recent studies have confirmed that the field is primarily dipolar. with some variation in the value of the polar field strength. with Ryabchikovaetal.(2007) finding D,=4.4kC and al.(2010) finding B,=4.2kC.."," More recent studies have confirmed that the field is primarily dipolar, with some variation in the value of the polar field strength, with \citet{Ryab97} finding $B_p=4.4\,\mathrm{kG}$ and \citet{Luftinger10} finding $B_p=4.2\,\mathrm{kG}$."
+ In this paper we seek to understand better how the amplitude of pulsation modes varies with time in 11217 through a wavelet analvsis of the 1986 and 2000 data., In this paper we seek to understand better how the amplitude of pulsation modes varies with time in 1217 through a wavelet analysis of the 1986 and 2000 data.
+ Furthermore. since the analysis of the 1986 and 2000 data by Kurtzetal.(1989) and Kurtzetal. did not use explicit statistical weighting. we also calculate weights for these data.," Furthermore, since the analysis of the 1986 and 2000 data by \citet{Kurtz89} and \citet{Kurtz05a} did not use explicit statistical weighting, we also calculate weights for these data."
+ Considering the wide range in the apertures. instrumentation. atmospherie and weather conditions. there is much to be gained by weighting the data appropriately.," Considering the wide range in the apertures, instrumentation, atmospheric and weather conditions, there is much to be gained by weighting the data appropriately."
+ In weighting data. the estimated uncertainties. σι. are used.," In weighting data, the estimated uncertainties, $\sigma_i$, are used."
+ We can assign a weight. a). to the P data point. as «we;=Lay.," We can assign a weight, $w_i$, to the $i^{th}$ data point, as $w_i=1/\sigma_i^x$ ."
+ The value of the exponent. .r. depends on the nature of the noise in the data: for Gaussian noise αν=2. which is commonly used.," The value of the exponent, $x$, depends on the nature of the noise in the data: for Gaussian noise $x=2$, which is commonly used."
+ Handler(2003). found that a smaller exponent was more effective in reducing the noise level for some data sets. so we do not assume .c—2. but rather find the value that minimizes the noise level.," \citet{Handler03} found that a smaller exponent was more effective in reducing the noise level for some data sets, so we do not assume $x=2$, but rather find the value that minimizes the noise level."
+ It must be noted that low-frequency sky-transparency noise was removed from the data by successive removal of low-frequency peaks (generally below mmHz. and always well separated from the oscillation signal at mmHz) to the point that low-frequency noise was at the same amplitude as noise at higher frequencies.," It must be noted that low-frequency sky-transparency noise was removed from the data by successive removal of low-frequency peaks (generally below mHz, and always well separated from the oscillation signal at mHz) to the point that low-frequency noise was at the same amplitude as noise at higher frequencies."
+ This was done to obtain approximately white noise across the frequency spectrum so that least-squares errors could be appropriately estimated for the original analysis of Kurtzetal.(1989) and Kurtzetal.(2005).., This was done to obtain approximately white noise across the frequency spectrum so that least-squares errors could be appropriately estimated for the original analysis of \citet{Kurtz89} and \citet{Kurtz05a}.
+ For our weighting scheme. this means that we do expect .r2.," For our weighting scheme, this means that we do expect $x\approx2$."
+ For neither the 1986 nor the 2000 data do we have uncertainties available. and so we needed to determine them.," For neither the 1986 nor the 2000 data do we have uncertainties available, and so we needed to determine them."
+ To do this. the power in the range of frequencies in which oscillations are present was temporarily removed from the time series using a notch filter.," To do this, the power in the range of frequencies in which oscillations are present was temporarily removed from the time series using a notch filter."
+ Thestandard deviation. o;. in the residual time series was then used to calculate the weights.," Thestandard deviation, $\sigma_i$, in the residual time series was then used to calculate the weights."
+ The timespan around each data point over which o; was calculated was treated as a free parameter., The timespan around each data point over which $\sigma_i$ was calculated was treated as a free parameter.
+ To find the best weighting scheme. we varied the timescale to tind the one that minimized the average noise level at high frequencies withthe exponent 1= 2.," To find the best weighting scheme, we varied the timescale to find the one that minimized the average noise level at high frequencies withthe exponent $x=2$ ."
+ Upon settling on this timescale. the exponent was varied to minimize the noise level further.," Upon settling on this timescale, the exponent was varied to minimize the noise level further."
+ The optimal timescale for both 1986 and 2000 was found to be about 20 minutes. a value presumably related to the typical," The optimal timescale for both 1986 and 2000 was found to be about 20 minutes, a value presumably related to the typical"
+absorber. suggesting a gas rich environment. distributed at all angles.,"absorber, suggesting a gas rich environment, distributed at all angles."
+ As also pointed out bv Sambrunaetal.(1999).. who analvzed a large sample of AGN withτς. large cold gas absorbing columus are rather common in BLRG.," As also pointed out by \citet{sam99}, who analyzed a large sample of AGN with, large cold gas absorbing columns are rather common in BLRG."
+ On the other hand. in the picture outlined above. a larger C; implies larger iron line EW values. which are. instead. nol observed.," On the other hand, in the picture outlined above, a larger $C_{\rm f}$ implies larger iron line EW values, which are, instead, not observed."
+ analvsis of 35 AGN has shown the existence of a variety of X-ray spectra., analysis of 35 AGN has shown the existence of a variety of X-ray spectra.
+ In general. the shape of the spectrum appears smooth in FSRQ and rich of features in DLRG.," In general, the shape of the spectrum appears smooth in FSRQ and rich of features in BLRG."
+ In SSRQ the analvsis is limited by poor statistics., In SSRQ the analysis is limited by poor statistics.
+ ILowever in the only case with hieh signal to noise spectrum. some similarity wilh BLRG is evident.," However in the only case with high signal to noise spectrum, some similarity with BLRG is evident."
+ FSRQ in our sample are characterized by a harder power law aud by (the general absence of reprocessed features., FSRQ in our sample are characterized by a harder power law and by the general absence of reprocessed features.
+" This is in clear agreement with the well consolidated idea that the jel component dominates this class of AGN,", This is in clear agreement with the well consolidated idea that the jet component dominates this class of AGN.
+ In the other AGN. the traces of a non-thermal Doppler component are more elusive and the results controversial and not conclusive vel.," In the other AGN, the traces of a non-thermal Doppler component are more elusive and the results controversial and not conclusive yet."
+ Neither (he spectral slopes nor the hieh energy cutolIs seem to be related to Core Badio Dominance (11) parameter. indicator of (he orientation of the beamecl radiation wilh respect to the light of sight.," Neither the spectral slopes nor the high energy cutoffs seem to be related to Core Radio Dominance $R$ ) parameter, indicator of the orientation of the beamed radiation with respect to the light of sight."
+ However a possible anti-correlation between the reprocessed features (EW. reflection) ancl 22 is found. although not statistically robust.," However a possible anti-correlation between the reprocessed features (EW, reflection) and $R$ is found, although not statistically robust."
+ A FAV=Ls44 anti-correlation (Daldwin effect) is observed in AGN., A $EW-L_{\rm 2-10 }$ anti-correlation (Baldwin effect) is observed in AGN.
+ This could be simply interpreted in terms of jet dilution. if the case of 3C4 373. in which beamed 10n-Chermal and thermal radiation are untanelec. is extrapolated to other sources.," This could be simply interpreted in terms of jet dilution, if the case of 3C 373, in which beamed non-thermal and thermal radiation are untangled, is extrapolated to other sources."
+ However. since Sevlert 1 galaxies strengthen the case for the anticorrelation. other causes apart the jet must exist which could weaken the Fe line EW in non-blazar ACN.," However, since Seyfert 1 galaxies strengthen the case for the anticorrelation, other causes apart the jet must exist which could weaken the Fe line EW in non-blazar AGN."
+ As recently proposed by Miniutti&Fabian(2004).. the Daldwin effect could be explained or example within a light bending model.," As recently proposed by \citet{min04}, the Baldwin effect could be explained for example within a light bending model."
+ Η this model is extended to AGN. the weaker iron lines and larger X-ray [Iuxes observed in BLRG could then be simply explained in terms of larger heights of the primary source above a cold untruncated disk.," If this model is extended to AGN, the weaker iron lines and larger X-ray fluxes observed in BLRG could then be simply explained in terms of larger heights of the primary source above a cold untruncated disk."
+ Incidentally. we note that a reduction of the reprocessed features could be also explained by a mildly relativistic motion of the X-ray source (lor example a hot corona) directed. away [rom the disk (Belobodorov1999a.b:Revnolds&Fabian1997).," Incidentally, we note that a reduction of the reprocessed features could be also explained by a mildly relativistic motion of the X-ray source (for example a hot corona) directed away from the disk \citep{bel9a,bel9b,rey97}."
+. Alternative explanations seem to be less supported by our results., Alternative explanations seem to be less supported by our results.
+ The signilicantlv, The significantly
+"in the DII. model below z = 0.1 is due (o the host galaxy maturing and increasing in huminosity. the contribution from the entire host galaxy must be included in the magnitudes used in the logz-m, plot if the brightening is to be detected.","in the DIR model below z = 0.1 is due to the host galaxy maturing and increasing in luminosity, the contribution from the entire host galaxy must be included in the magnitudes used in the $_{v}$ plot if the brightening is to be detected."
+" Although the Iuminositv of the nucleus may be adequate in determining (he luminosity function of the active nucleus in the CR. model. because of the complex process required to identify AGN galaxies 2005b).. obviously great. care will be required in obtaining the magnitudes of low reclshilt AGN galaxies if thev are to be used in logz-m, plots."," Although the luminosity of the nucleus may be adequate in determining the luminosity function of the active nucleus in the CR model, because of the complex process required to identify AGN galaxies \citep{hao05b}, obviously great care will be required in obtaining the magnitudes of low redshift AGN galaxies if they are to be used in $_{v}$ plots."
+" The most Iuminous radio galaxies and first-ranked cluster galaxies have been compared here to the high luminositv edge of the AGN galaxy. distribution on a logz-m, plot.", The most luminous radio galaxies and first-ranked cluster galaxies have been compared here to the high luminosity edge of the AGN galaxy distribution on a $z$ $_{v}$ plot.
+ It is found that while the radio galaxies and cluster galaxies are good standard candles at all epochs. the luminosity of the AGN galaxies varies signifieantlv from one epoch to another.," It is found that while the radio galaxies and cluster galaxies are good standard candles at all epochs, the luminosity of the AGN galaxies varies significantly from one epoch to another."
+ Compared to the comparison galaxies (he AGN galaxies are found to be super-Juminous at high redshifts. but become sub-Inminous as the redshift decreases.," Compared to the comparison galaxies the AGN galaxies are found to be super-luminous at high redshifts, but become sub-luminous as the redshift decreases."
+ These new results show that below z = 0.3 the rate of Inminosity decrease begins to slow clown and below z = 0.1 the luminosity begins to increase again., These new results show that below $z $ = 0.3 the rate of luminosity decrease begins to slow down and below $z$ = 0.1 the luminosity begins to increase again.
+ Although their apparent super-Iuminous nature ad hieh z can be explained by a higher star formation rate. and the fact that there night have been more raw material around to make galaxies at that epoch. a luminosity increase below 2 = 0.1 is more difficult to explain when these arguments are unlikely explanations.," Although their apparent super-luminous nature at high z can be explained by a higher star formation rate, and the fact that there might have been more raw material around to make galaxies at that epoch, a luminosity increase below $z$ = 0.1 is more difficult to explain when these arguments are unlikely explanations."
+" It is therefore concluded here that the evidence favors the argument that the high redshift AGN ealaxies (quasars) Chat do lie above the mature galaxy. line on a logzan, plot have all been pushed there because of a large intrinsic component in their redshilts ancl not because they have a superimposed non-thermal component that is many magnitudes brighter (han that seen in radio galaxies.", It is therefore concluded here that the evidence favors the argument that the high redshift AGN galaxies (quasars) that do lie above the mature galaxy line on a $z$ $_{v}$ plot have all been pushed there because of a large intrinsic component in their redshifts and not because they have a superimposed non-thermal component that is many magnitudes brighter than that seen in radio galaxies.
+ All AGN galaxies then will be sub-Inuminous to mature galaxies. as predicted in the DIR model.," All AGN galaxies then will be sub-luminous to mature galaxies, as predicted in the DIR model."
+ For a given intrinsic redshift component. all are likely also to be eood standard candles.," For a given intrinsic redshift component, all are likely also to be good standard candles."
+" Finally. ifit (turns out that many faint AGN galaxies at low redshifts have been missed in a particular magnitude range the conclusion that the bright edge of the logzan, plot increases again in (his redshift range may need to be re-evaluated."," Finally, if it turns out that many faint AGN galaxies at low redshifts have been missed in a particular magnitude range the conclusion that the bright edge of the $_{v}$ plot increases again in this redshift range may need to be re-evaluated."
+ Such an ellect might be introduced by the selection elfect discussed by Haoetal.(2005a.b).. but only if the VCVeat contains manv of the SDSS AGN galaxies. as explained in Sec 3.1.," Such an effect might be introduced by the selection effect discussed by \citet{hao05a,hao05b}, but only if the VCVcat contains many of the SDSS AGN galaxies, as explained in Sec 3.1."
+ ] wish to thank two anouvmous referees for suggestions on how this paper might be improved., I wish to thank two anonymous referees for suggestions on how this paper might be improved.
+ I also thank Dr. D. MeDiarmid for helpful comments., I also thank Dr. D. McDiarmid for helpful comments.
+The observations give the power spectrum in angle-angle-redshift space rather than in real space.,The observations give the power spectrum in angle-angle-redshift space rather than in real space.
+ The inverse of the peak positions in the observed power spectrum therefore gives the angular and redshift intervals corresponding to the sonic horizon., The inverse of the peak positions in the observed power spectrum therefore gives the angular and redshift intervals corresponding to the sonic horizon.
+" The peaks in the angular spectrum are proportional to dy(2)/a, and those in the redshift spectrum to djCO/a,.", The peaks in the angular spectrum are proportional to $d_T(z)/a_s$ and those in the redshift spectrum to $d_H(z)/a_s$.
+ dy~IOSMpe is the acoustic horizon comoving size at recombination. dy(z)=(1+dy is the comoving angular distance and dj=c/H(z) is the Hubble distance (see Eq. 6)):," $a_s \sim 105 h^{-1} \mathrm{Mpc}$ is the acoustic horizon comoving size at recombination, $d_T(z) = (1+z) \dang$ is the comoving angular distance and $d_H=c/H(z)$ is the Hubble distance (see Eq. \ref{eq:expHz}) ):"
+" The quantities dy. dj; and a, all depend on the cosmological parameters."," The quantities $d_T$ , $d_H$ and $a_s$ all depend on the cosmological parameters."
+ Figure 17. gives the angular and redshift intervals as a function of redshift for four cosmological models., Figure \ref{fig:hubble} gives the angular and redshift intervals as a function of redshift for four cosmological models.
+ The error bars on the lines for (Qi.O4)=(0.27.0.73) correspond to the expected errors on the peak positions taken from Table 7 for the four-month runs with the packed array.," The error bars on the lines for $(\Omega_M,\Omega_\Lambda)=(0.27,0.73)$ correspond to the expected errors on the peak positions taken from Table \ref{tab:ErrorOnK}
+ for the four-month runs with the packed array."
+ We see that with these uncertainties. the data would be able to measure ww at better than the level.," We see that with these uncertainties, the data would be able to measure $w$ at better than the level."
+ To estimate the sensitivity to parameters describing dark energy equation of state. we follow the procedure explained in. (BlakeandGlazebrook(2003)).," To estimate the sensitivity to parameters describing dark energy equation of state, we follow the procedure explained in \citep{blake.03}."
+". We can introduce the equation of state of dark energy. wil)=wo+Wyz/(o2) by replacing Q, in the definition of dy(z) and ας). CEq. 40))"," We can introduce the equation of state of dark energy, $w(z)=w_0 + w_a\cdot z/(1+z)$ by replacing $\Omega_\Lambda$ in the definition of $d_T (z)$ and $d_H (z)$, (Eq. \ref{eq:dTdH}) )"
+ by: where ος is the present-day dark energy fraction with respect to the critical density., by: where $\Omega_{\Lambda}^0$ is the present-day dark energy fraction with respect to the critical density.
+ Using the relative errors on Kpao and Keay given in Tab. 7..," Using the relative errors on $\koperp$ and $\kopar$ given in Tab. \ref{tab:ErrorOnK},"
+" we can compute the Fisher matrix for five cosmological parameter: (O,,.Qe./1.o.Wy)."," we can compute the Fisher matrix for five cosmological parameter: $(\Omega_m, \Omega_b, h, w_0, w_a)$."
+ Then. the combination of this BAO Fisher matrix with the Fisher matrix obtained for Planck mission. allows us tocompute the errors on dark energy parameters.," Then, the combination of this BAO Fisher matrix with the Fisher matrix obtained for Planck mission, allows us tocompute the errors on dark energy parameters."
+ The Planck Fisher matrix ts obtained for the 8 parameters (assuming a flat universe): Oy. Q;. Jr. wo. Wy. Cs. Hy (spectral index of the primordial power spectrum) and r (optical depth to the last-scatter surface).," The Planck Fisher matrix is obtained for the 8 parameters (assuming a flat universe): $\Omega_m$, $\Omega_b$, $h$, $w_0$, $w_a$, $\sigma_8$, $n_s$ (spectral index of the primordial power spectrum) and $\tau$ (optical depth to the last-scatter surface)."
+ For an optimized project over a redshift range. 0.25.«< 2.75. with a total observation time of 3 years. the packed 400-dish interferometer array has a precision of on wo and on wy.," For an optimized project over a redshift range, $0.25<z<2.75$ , with a total observation time of 3 years, the packed 400-dish interferometer array has a precision of on $w_0$ and on $w_a$."
+ The Figure of Merit. the inverse of the area in the confidence level contours 1s 38.," The Figure of Merit, the inverse of the area in the confidence level contours is 38."
+ Finally. Fig.," Finally, Fig."
+" 18. shows a comparison of different BAO projects. with a set of priors on (Q,,.Q,./1) corresponding to the expected precision on these parameters in early 2010°s."," \ref{fig:Compw0wa} shows a comparison of different BAO projects, with a set of priors on $(\Omega_m, \Omega_b, h)$ corresponding to the expected precision on these parameters in early 2010's."
+ This BAO project based on intensity mapping ts clearly competitive with the current generation of optical surveys such as SDSS-III (SDSS-IH(2008))., This BAO project based on intensity mapping is clearly competitive with the current generation of optical surveys such as SDSS-III \citep{sdss3}.
+. The 3D mapping of redshifted 21 cm emission though is a novel and complementary approach to optical surveys to study the statistical properties of the large scale structures in the universe up to redshifts zx 3., The 3D mapping of redshifted 21 cm emission though is a novel and complementary approach to optical surveys to study the statistical properties of the large scale structures in the universe up to redshifts $z \lesssim 3$ .
+ A radio instrument with large instantaneous field ofview (10-100 deg?) and large bandwidth (2100 MHz) with ~10 aremin resolution, A radio instrument with large instantaneous field ofview (10-100 $^2$) and large bandwidth $\gtrsim 100$ MHz) with $\sim 10$ arcmin resolution
+ inodel (7). includes a range of surface eravitics and effective temperatures and incorporates thermal electron conduction aud selfirradiation by photons from he compact object., model \citep{heinke06apj} includes a range of surface gravities and effective temperatures and incorporates thermal electron conduction and self-irradiation by photons from the compact object.
+ It also asstunes negligible (<10? C) naguetic fields and a pure hydrogen atmosphere., It also assumes negligible $<$ $^9$ G) magnetic fields and a pure hydrogen atmosphere.
+ The xuwanmeters of the component are the NS mass and radius. the effective temperature in the NS frame. he source distance aud a normalisation factor. which xuanmeterizes the fraction of the surface that is radiating.," The parameters of the component are the NS mass and radius, the effective temperature in the NS frame, the source distance and a normalisation factor, which parameterizes the fraction of the surface that is radiating."
+ Throughout this work. we imposed a lower luit to the nass of the NS of M... as derived by ?..," Throughout this work, we imposed a lower limit to the mass of the NS of $\Msun$ , as derived by \citet{0748:bassa09mnras}."
+ Further. suce NSs in LAINBs do not show evidence. such as spin-powered pulsations or cvclotron spectral features. or a stroug maesnuetie field. we kept the normalisation of the retro star anosphere model fixed to 1. which corresponds to the eutire NS surface ciitting.," Further, since NSs in LMXBs do not show evidence, such as spin-powered pulsations or cyclotron spectral features, for a strong magnetic field, we kept the normalisation of the neutron star atmosphere model fixed to 1, which corresponds to the entire NS surface emitting."
+ We performed 3 fits per observation: first we left the mass of the NS as a ree paramcter diving the fit aud then we fixed its value to the canonical mass of £M. aud to Lat. for comparison., We performed 3 fits per observation: first we left the mass of the NS as a free parameter during the fit and then we fixed its value to the canonical mass of $\Msun$ and to $\Msun$ for comparison.
+ We extrapolated the fiux of the thermal component to the energy rauge of 0.01-100 keV to estimate the thermal bolometric flux., We extrapolated the flux of the thermal component to the energy range of 0.01-100 keV to estimate the thermal bolometric flux.
+ The fits were acceptable. with reducedντι.\2.. ~ 11 for 250.150 d.o.£..," The fits were acceptable, with reduced, $\sim$ 1 for 250–450 d.o.f.,"
+ for all the observations., for all the observations.
+ The inclusion of a power-law component improved the goodness of the ft only for Obs 1., The inclusion of a power-law component improved the goodness of the fit only for Obs 1.
+ The vost-fit results are shown in Table 2 fora distance of 7.1 kpe., The best-fit results are shown in Table \ref{tab:epic-rgs-fit} for a distance of 7.1 kpc.
+ The errors on the temperature aud radius of the NS are significantly simaller when the mass of the NS is fixed., The errors on the temperature and radius of the NS are significantly smaller when the mass of the NS is fixed.
+ The value of sshows siguificaut differences amoug observations., The value of shows significant differences among observations.
+ This clearly affects the accuracy with which we cau measure the value of the NS radius aud its effective temperature. and shows that a single observation is not sufficicut to inpose constraints ou the mass and radius of tle NS.," This clearly affects the accuracy with which we can measure the value of the NS radius and its effective temperature, and shows that a single observation is not sufficient to impose constraints on the mass and radius of the NS."
+ Iu order to obtain the best possible coustraints for the mass and radius of the NS with the απο of observations. we next fitted the EPIC and RGS spectra of all the observations simultaueoushv.," In order to obtain the best possible constraints for the mass and radius of the NS with the set of observations, we next fitted the EPIC and RGS spectra of all the observations simultaneously."
+ We first asstuned that the value of the absorption docs not chanec annone observations and tied the values ofNyy... and lass and radius of the NS for all the observations.," We first assumed that the value of the absorption does not change among observations and tied the values of, and mass and radius of the NS for all the observations."
+ Since the coutribution of the power-law componcut was ouly significant iu the first observation(see Sect. 3.2.2)).," Since the contribution of the power-law component was only significant in the first observation(see Sect. \ref{sec:epic-rgs}) ),"
+QSO emission in the UV part of the spectrum.,QSO emission in the UV part of the spectrum.
+ We cannot exclude however. the possibility of scattered radiation from the AGN itself as the origin of this excess.," We cannot exclude however, the possibility of scattered radiation from the AGN itself as the origin of this excess."
+ It is interesting nevertheless. that 2MASS.33. one of the objects that shows the UV excess. has narrow emission lines in addition to the broad wings of the Lla.," It is interesting nevertheless, that 3, one of the objects that shows the UV excess, has narrow emission lines in addition to the broad wings of the $\rm H\alpha$."
+ The flux ratios of the narrow emission line components place this source in the region of Transition Objects in the diagnostic diagram of 2.. suggesting some level of star-formation.," The flux ratios of the narrow emission line components place this source in the region of Transition Objects in the diagnostic diagram of \cite{Kewley2001}, suggesting some level of star-formation."
+ In summary. the red 2MASSt sources studied ⋡⋡in this. paper are reddened QSOs. whileMM a. large fraction (7/10⊲∡ in proiTable 3)) have samples far-IR properties consistent with star-formation in the host galaxy.," In summary, the red 2MASS sources studied in this paper are dust reddened QSOs, while a large fraction (7/10 in Table \ref{tab_restframe}) ) have far-IR properties consistent with star-formation in the host galaxy."
+ In this section we first explore possible systematic differencesin the IR SEDs of the red 2MASS QSOs with optically selected ones., In this section we first explore possible systematic differences in the IR SEDs of the red 2MASS QSOs with optically selected ones.
+ For this comparison we use the Palomar-Green (PG) QSO sample ?.. for which a rich set of multiwavelength observations is available.," For this comparison we use the Palomar-Green (PG) QSO sample \cite{Schmidt1983}, for which a rich set of multiwavelength observations is available."
+ We considerin particular the of sub-sample64 PG QSOs with mid- and far-IR observations (5.— 2007/m in the ISO archive analysed and presented by ?.., We consider in particular the sub-sample of 64 PG QSOs with mid- and far-IR observations $\rm 5-200\mu m$ ) in the ISO archive analysed and presented by \cite{Haas2003}.
+ From that sub-sample only sources with SDSS optical photometry (total of 26) are used here., From that sub-sample only sources with SDSS optical photometry (total of 26) are used here.
+ The near-IR data for these sources are from the 2MASS survey., The near-IR data for these sources are from the 2MASS survey.
+ In order to avoid uncertainties associated with the photometrie redshift estimates. in this section we consider2MASS QSOs with spectroscopic redshifts only.," In order to avoid uncertainties associated with the photometric redshift estimates, in this section we consider 2MASS QSOs with spectroscopic redshifts only."
+ The median redshift of the PG QSO subsample used here is z 0.161. lower than that of 2MASS QSOs. +0.5.," The median redshift of the PG QSO subsample used here is $\approx0.161$ , lower than that of 2MASS QSOs, $z\approx0.5$."
+ Figure S compares the rest-frame UV-to-IR SEDs of the 2MASS QSOs with those of the optically selected PG QSOs., Figure \ref{fig_sed_all} compares the rest-frame UV-to-IR SEDs of the 2MASS QSOs with those of the optically selected PG QSOs.
+ The SEDs are normalised to the rest-frame 12j/m flux density. as this wavelength is proposed to be nearly unbiased census of the AGN power (e.g.21.," The SEDs are normalised to the rest-frame $\rm 12\mu m$ flux density, as this wavelength is proposed to be nearly unbiased census of the AGN power \citep[e.g.][]{Spinoglio1989}."
+ As expected dust reddening makes the UV/optical SEDs of 2MASS QSOs steeper than PG QSOs., As expected dust reddening makes the UV/optical SEDs of 2MASS QSOs steeper than PG QSOs.
+ Also. there is evidence that the far-IR dustproperties of the two differ.," Also, there is evidence that the far-IR properties of the two samples differ."
+ The normalised rest-frame far-IR flux density of at least 4 out of 8 2MASS sources spectroscopic redshifts is elevated compared to the median SED of PG QSOs., The normalised rest-frame far-IR flux density of at least 4 out of 8 2MASS sources spectroscopic redshifts is elevated compared to the median SED of PG QSOs.
+ Our SED modeling (see Table 3)) and recent studies of optically selected QSOs (e.g22)... suggest that the far-IR emission of these monsters is dominated by star-formation.," Our SED modeling (see Table \ref{tab_restframe}) ) and recent studies of optically selected QSOs \citep[e.g][]{Haas2003, Schweitzer2006}, suggest that the far-IR emission of these monsters is dominated by star-formation."
+ Under this assumption. - ↕∶↥≝⋯⊾∁∃↾∣↧∁⇂⊾∁↑⊲⋯⊾∁⊰∣⋯∖↖⊰↾∣↧⋡∙∣↾∸∣↾∣∁⋡∙∣⊰↾↭∣∏∁⊐∖↧↱∖⊊⊊⊖⊊⊙⊰ have. on average. higher. star-formation. than optically. selected QSOs.," Under this assumption, Figure \ref{fig_sed_all} therefore shows that at least some 2MASS QSOs have, on average, higher star-formation than optically selected QSOs."
+ This is further demonstrated in Figure 6 which plots 60a luminosity against 12jim luminosity (see figure caption for details on the determination of the luminosities).," This is further demonstrated in Figure \ref{fig_ll}
+ which plots $\rm 60\, \mu m$ luminosity against $\rm 12 \, \mu m$ luminosity (see figure caption for details on the determination of the luminosities)."
+ For given 12sam luminosity. some 2MASS QSOs are more luminous at 60jum.," For given $\rm 12 \, \mu m$ luminosity, some 2MASS QSOs are more luminous at $\rm 60\, \mu m$."
+ The assessment of the significance of this excess is limited by the small sample size., The assessment of the significance of this excess is limited by the small sample size.
+ Nevertheless. we use survival statistics as implemented in the ASURV package (2?) to estimate the null hypothesis probability that the distribution of the luminosity ratio logeL.(1259m)/Εν(θα) of 2MASS and PG QSOs is drawn from the same parentpopulation.," Nevertheless, we use survival statistics as implemented in the ASURV package \citep{Isobe1986, Lavalley1992} to estimate the null hypothesis probability that the distribution of the luminosity ratio $\log \rm \nu L_\nu (12\mu m) / \nu L_\nu (60\mu m)$ of 2MASS and PG QSOs is drawn from the same parentpopulation."
+ Using the Gehan's generalised Wilcoxon test we estimate this probability to be 3 per cent., Using the Gehan's generalised Wilcoxon test we estimate this probability to be 3 per cent.
+ The, The
+where D is in Mpe. Όρο IS 3n fas du days aud ko=3.5697-101! is the conversion factor between the quantities.,"where $D$ is in Mpc, $v_{phot}$ is in $^{-1}$, $t$ is in days and $k = 3.5697 \cdot 10^{14}$ is the conversion factor between the quantities."
+" Plotting f versus (/c4,,4. D aud fy cau bo determined by linear regression."," Plotting $t$ versus $\theta / v_{phot}$, $D$ and $t_0$ can be determined by linear regression."
+ The results for SN 2002ap are plotted im Figs.5. 9 aud 10.," The results for SN 2002ap are plotted in Figs.8, 9 and 10."
+ Fie.8 contains the plot of the temperature variation (from dereddenued (2BWo data. via Eq.5).," Fig.8 contains the plot of the temperature variation (from dereddened $(B-V)_0$ data, via Eq.5)."
+ Fig.9 shows 1e angular size calculated from Eq.l., Fig.9 shows the angular size calculated from Eq.4.
+ It is appareut that je expansion of the photosphere stopped at about JD we152317. so the expansion phase did not last lone.," It is apparent that the expansion of the photosphere stopped at about JD 2452317, so the expansion phase did not last long."
+ Note. iat. as the referee poiuted out. this does not mean that jephotospherie phase has euded. because the spectra clearly indicate that the SN is still in the photospheric hase (see Fie.5).," Note, that, as the referee pointed out, this does not mean that the phase has ended, because the spectra clearly indicate that the SN is still in the photospheric phase (see Fig.5)."
+ The EPM was applied only for data obtained before this epoch., The EPM was applied only for data obtained before this epoch.
+ The fitting of Eq.6 to the conibined data is preseuted in Fie.10., The fitting of Eq.6 to the combined data is presented in Fig.10.
+" The linear regressiou resulted iu D = 6.7 £0.5 Mpe aud ty = 2152301.3 41.0 as the distance and the moment of explosion. respectively,"," The linear regression resulted in $D$ = 6.7 $\pm 0.5$ Mpc and $t_0$ = 2452304.3 $\pm 1.0$ as the distance and the moment of explosion, respectively."
+ It is clear froin Fie.lO that the relation becomes non-linear when the expansion phase ends. thus. at least oue of the assuniptious of the EPM became invalid at that time.," It is clear from Fig.10 that the relation becomes non-linear when the expansion phase ends, thus, at least one of the assumptions of the EPM became invalid at that time."
+ The quoted uncertainties are formal (random) errors due to the fitting procedure., The quoted uncertainties are formal (random) errors due to the fitting procedure.
+ These results show remarkably good agreement with previous ones. determined via other methods.," These results show remarkably good agreement with previous ones, determined via other methods."
+ The EPM-distance is almost perfectly equal to the photometric distance of M71 when the Schlegeletal.1905. recdclenine correction is applied (4—G.S Mpc. see Sect.1).," The EPM-distance is almost perfectly equal to the photometric distance of M74 when the \cite{sfd}
+ reddening correction is applied $d = 6.8$ Mpc, see Sect.1)."
+ The moment of explosion agrees within the uncertainties with the one derived by Mazzalietal.2002. (fy=21522303.1. Jan. 28. 2002) from model computations.," The moment of explosion agrees within the uncertainties with the one derived by \cite{mazz} $t_0 = 2452303.4$, Jan. 28, 2002) from model computations."
+ It also worths noting that the derived explosion date is ouly 1 dav prior to the discovery of SN 2002ap (see Sect.1)., It also worths noting that the derived explosion date is only 1 day prior to the discovery of SN 2002ap (see Sect.1).
+ As mentioned above. difference between the assuiiptious of the EPM. and the true physical state of the SN atmosphere may lead to significant uncertaiutv in the inferred distance.," As mentioned above, difference between the assumptions of the EPM and the true physical state of the SN atmosphere may lead to significant uncertainty in the inferred distance."
+ Despite the good agreement betweeu the new LPM-distance and other distance estimates. all these may suffer from quite large systematic errors.," Despite the good agreement between the new EPM-distance and other distance estimates, all these may suffer from quite large systematic errors."
+" Receutly, Leonardctal.2003. preseuted a thorough[um test of the EPAL aud Cepleid-based distances of SN 1999011 and NGC 1637."," Recently, \cite{leo3} presented a thorough test of the EPM- and Cepheid-based distances of SN 1999em and NGC 1637."
+ They coucluded that the Cepheid-based distance is 50 ereater than the EPME-based. distance. which meaus that there is serious disagreement between hese two measurement methods.," They concluded that the Cepheid-based distance is 50 greater than the EPM-based distance, which means that there is serious disagreement between these two measurement methods."
+ They also showed that. at present. the Cepheid-based distance is more favourable han he EPM-based one. thus. the EPALEdistances are rlaeued oat leas systematic error.," They also showed that, at present, the Cepheid-based distance is more favourable than the EPM-based one, thus, the EPM-distances are plagued by at least systematic error."
+ This leads oa GU - 70% total (systematic plus random) uncertaintyv for he EPA, This leads to a 60 - 70 total (systematic plus random) uncertainty for the EPM.
+ Wha are the possible sources of such lüeh uncertainty?, What are the possible sources of such high uncertainty?
+ Leonardetal.2003Pa mention the dilution actor aud the radial velocity measurement technique as he most probable ones., \cite{leo3} mention the dilution factor and the radial velocity measurement technique as the most probable ones.
+ The need for a precise dilution actor based on model atinosphliere tuned for the particular SN has been emphasized. bx Eastinanetal.1996.., The need for a precise dilution factor based on model atmosphere tuned for the particular SN has been emphasized by \cite{eastman}. .
+ ILoxwever. their dilutiou factors may also suffer fro nucertainics. because indepeudent caleulatious bv other eroups have lead to significantly differeut factors (sce Leonardetal.2003/— for extensive discussion aud references).," However, their dilution factors may also suffer from uncertainies, because independent calculations by other groups have lead to significantly different factors (see \cite{leo3}
+ for extensive discussion and references)."
+ These xoblenmis are probably more eublianced for the present study. due to the atmospheric properties of lypernovac.," These problems are probably more enhanced for the present study, due to the atmospheric properties of hypernovae."
+ Since there was no detailed model atmosphere for SN 2002ap at our disposal. we had to assiuune ¢=1. which may lead up to 50% systematic error in the distance (as it was mentioned in Section 3.1).," Since there was no detailed model atmosphere for SN 2002ap at our disposal, we had to assume $\zeta = 1$, which may lead up to 50 systematic error in the distance (as it was mentioned in Section 3.1)."
+ This error can be even larger if the atmosphere of ον 2002ap. were very different from a Type II SN atinosphere so that the dilution actor Was quite differeut from those of Eastinanetal.1996., This error can be even larger if the atmosphere of SN 2002ap were very different from a Type II SN atmosphere so that the dilution factor was quite different from those of \cite{eastman}.
+. Ou the other haud. the application of a wpcrnova model of Mazzali.Iwainoto&Nomoto.2000 (constructed for SN 19970f) may somewhat reduce this arge πουΤαήν.," On the other hand, the application of a hypernova model of \cite{maz97} (constructed for SN 1997ef) may somewhat reduce this large uncertainty."
+ As it is explained above. we assigned the model effective cluperature (via Eq.5. by comparing the observed aud xedieted B.V ‘olor indices) to thebolometric fux data.," As it is explained above, we assigned the effective temperature (via Eq.5, by comparing the observed and predicted $B-V$ color indices) to the flux data."
+" This technique nav have the advantage of using he CVOZR “bolometric” fluxes tha are mteerated over a long wavelength reenne. thus. if is not sensitive to he chosen waveleneth aud waveleneth-depeudent dilution “actors,"," This technique may have the advantage of using the $UVOIR$ “bolometric” fluxes that are integrated over a long wavelength regime, thus, it is not sensitive to the chosen wavelength and wavelength-dependent dilution factors."
+ Ou the other haud. the wnacertaities of the derived bolometric fixes are probably iu the same order of maguitude as that of the dilution factor.," On the other hand, the uncertainties of the derived bolometric fluxes are probably in the same order of magnitude as that of the dilution factor."
+ The systematic errors caused by the measurement method of radial velocities are also quite laree for SN 2002ap because of the very broad lines audstrong blends iu the spectra. especiallyat the carly phases.," The systematic errors caused by the measurement method of radial velocities are also quite large for SN 2002ap because of the very broad lines andstrong blends in the spectra, especiallyat the early phases."
+with and the mean surface density inside the radius c can be written as (in (see Golse&Ixneib (2002))).,with and the mean surface density inside the radius $x$ can be written as with (see \citet{GolKne02}) ).
+" The deflection angle a. convergence & aud shear > (urn out as where &,=SPPMA. wilh Xa=CD./(xGDjD,.)."," The deflection angle $\mathbf{\alpha}$, convergence $\kappa$ and shear $\gamma$ turn out as where $\kappa_s= \delta_c \rho_c r_s\Sigma_\mathrm{crit}^{-1}$, with $\Sigma_\mathrm{crit}
+\equiv \mathfrak{c}^{2} D_{s} / (4 \pi G D_{d} D_{ds})$."
+ By integrating the deflection angle. the potential ο”) can be found: with," By integrating the deflection angle, the potential $\varphi(x)$ can be found: with"
+3.6 and 4.5 ddetections. we find a total MLR of up to 0.9-2.4x10 7 Ἐν ΠΟΠ is likely an overestimate due to (he inclusion of objects photometrically scattered to their relatively red positions in the IR CMD.,"3.6 and 4.5 detections, we find a total MLR of up to $\times$ $^{-3}$ $^{-1}$, which is likely an overestimate due to the inclusion of objects photometrically scattered to their relatively red positions in the IR CMD."
+ For the region of the galaxy covered by the IIST/WEDPC? study of Dolphin(2000)... we find of the initial mass in stars lormed over the last Gyr is currently being returned into the ISAT via mass-loss from AGB stars. assuming hall of the total mass lost by an ACB star occurs in the dust enshrouded phase.," For the region of the galaxy covered by the HST/WFPC2 study of \citet{dol00}, we find of the initial mass in stars formed over the last Gyr is currently being returned into the ISM via mass-loss from AGB stars, assuming half of the total mass lost by an AGB star occurs in the dust enshrouded phase."
+ This value is somewhat higher than expected. eiven a total of is estimated to be returned from the entire stellar population 1994).. and our data are not sensitive to the contribution from supernovae and AGBs that have already died.," This value is somewhat higher than expected, given a total of is estimated to be returned from the entire stellar population \citep{ken94}, and our data are not sensitive to the contribution from supernovae and AGBs that have already died."
+ The prescription we use may systematically overestimate (he MLBs., The prescription we use may systematically overestimate the MLRs.
+ We find the distribution of MLBs and bolometrie Iuminosities of AGBs and RSGs are very similar to those in the LAIC and SAIC (vanLoonetal.1999.2006b). ancl the empirically derived maximum MLR of the LMC (vanLoonetal.1999). is in excellent agreement with our data.," We find the distribution of MLRs and bolometric luminosities of AGBs and RSGs are very similar to those in the LMC and SMC \citep{van99,van06b} and the empirically derived maximum MLR of the LMC \citep{van99} is in excellent agreement with our data."
+ Finally. we show that the AGB stars. which (race the intermediate age stellar population. provide the dominant contribution to the luminosity in the TRAC bands. while studies at shorter wavelengths can preferentially detect either much younger sources (in (he optical) or older age populations (short-wavelength near-IR).," Finally, we show that the AGB stars, which trace the intermediate age stellar population, provide the dominant contribution to the luminosity in the IRAC bands, while studies at shorter wavelengths can preferentially detect either much younger sources (in the optical) or older age populations (short-wavelength near-IR)."
+ We thank Martin Groenewegen for his conments on a previous version of this manuscript and for providing his models of mass-losing AGBs prior to publication., We thank Martin Groenewegen for his comments on a previous version of this manuscript and for providing his models of mass-losing AGBs prior to publication.
+ We thank Jacco van Loon and Maria-IRosa Cioni for their valuable comments and suggestions. which contributed significantlv to (he presentation of these data.," We thank Jacco van Loon and Maria-Rosa Cioni for their valuable comments and suggestions, which contributed significantly to the presentation of these data."
+ We also thank the anonvmous referee. [or their prompt and careful reading of the manuscript and their valuable comments., We also thank the anonymous referee for their prompt and careful reading of the manuscript and their valuable comments.
+ E. D. S. acknowledges partial support from à NASA LTSARP grant NAG 5-9221 and the University of Minnesota., E. D. S. acknowledges partial support from a NASA LTSARP grant NAG 5-9221 and the University of Minnesota.
+" This work is based in part on observations made with the Spitzer Space Telescope. which is operated bv the Jet Propulsion Laboratory, California Institute of Technology under NASA contract 1407."," This work is based in part on observations made with the Spitzer Space Telescope, which is operated by the Jet Propulsion Laboratory, California Institute of Technology under NASA contract 1407."
+ Support for this work was provided bv NASA through Contract Numbers 1256406 and 1215746 issued bv JPL/Caltech to the University of Minnesota., Support for this work was provided by NASA through Contract Numbers 1256406 and 1215746 issued by JPL/Caltech to the University of Minnesota.
+ This research has made use of NASA's Astrophysics Data System Bibliographic Services and ihe NASA/IPAC Extragalactic Database (NED) which is operated bv the Jet. Propulsion Laboratory. California Institute of Technology. under contract with the National Aeronautics and Space Administration.," This research has made use of NASA's Astrophysics Data System Bibliographic Services and the NASA/IPAC Extragalactic Database (NED) which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration."
+Observatoire de Paris and operated by CNRS/INSU (Iustitut National des Sciences de LIUnivers).,Observatoire de Paris and operated by CNRS/INSU (Institut National des Sciences de l'Univers).
+" Due to the cluster location with respect to the plate boundaries. ouly ἃ 2.57 42,5"" square reekπι roughly centered ou the cluster coordinates could iudoeed be searched for objects."," Due to the cluster location with respect to the plate boundaries, only a $^{\circ}\times$ $^{\circ}$ square region roughly centered on the cluster coordinates could indeed be searched for objects."
+ At the cluster redshift. this Inte sca is however large cuough to secure a correct Ivesigation of the overall cluster area (radius of ~ L5 sg \Qrc at least).," At the cluster redshift, this limited scan is however large enough to secure a correct investigation of the overall cluster area (radius of $\sim$ 4.5 $h_{50}^{-1}$ Mpc at least)."
+ The algoriluus invoVOC in the ATAAMTA ou-line mode available at that time for detecting and measuring astronomical sources are sununarized in Slezak (1998)., The algorithms involved in the MAMA on-line mode available at that time for detecting and measuring astronomical sources are summarized in Slezak (1998).
+" Basically, objecs appear to be defined as a set of connected pixels witi iuteusitv higher than typically a 2 sigma threshold alxwe the local sky backerounud aud hey are described bv their flux (the s1ii of backeround-subtracted pixel values. which lead Oo what is called rereatter a “plate” inagnitude). their area aud their elliptical shape pariuneters."," Basically, objects appear to be defined as a set of connected pixels with intensity higher than typically a 2 sigma threshold above the local sky background and they are described by their flux (the sum of background-subtracted pixel values, which lead to what is called hereafter a “plate” magnitude), their area and their elliptical shape parameters."
+" This apprch ds efficient iu ost cases,", This approach is efficient in most cases.
+ Ilowever. one must expec that fjs simple vision model fails for exwded fields where the probability o get bleued objects inercases drastically when low detectio1 thresholds are applied.," However, one must expect that this simple vision model fails for crowded fields where the probability to get blended objects increases drastically when low detection thresholds are applied."
+ The ealactic atitude of ABCC 196 is bp=3.17 and some blends with stars are indeed included im the :35.511 individual sources listed in he on-line catalogue we (were provided with (c.," The galactic latitude of ABCG 496 is $b_{\rm II} \simeq -34.4^{\circ}$ and some blends with stars are indeed included in the 35,541 individual sources listed in the on-line catalogue we were provided with (cf."
+ the final visual check of the galaxv candidates which is described low)., the final visual check of the galaxy candidates which is described below).
+ Tje astrometric reduction of the whole catalogue was oxforiued with respect to 98 stars of the PPM star catalogue (Roeser Bastian 1991) spread over the plate using a ον polyποια]. fitting., The astrometric reduction of the whole catalogue was performed with respect to 98 stars of the PPM star catalogue (Roeser Bastian 1991) spread over the plate using a $^{\rm rd}$ -order polynomial fitting.
+ The residuals of he fit vielding the instrmucutal constants were sialler han 0.20 arcsecond auxL the astrometry of our catalogue indecd appears to be very good. as coufirmed by eross-checking galaxy coordinates with the literature or the APAI database (ποσα average separaion: 0.7 arcsecoud) and by our multiobject fibre spectroscopy follow-up where the galaxies were always fou o be very close to he expected poSitios.," The residuals of the fit yielding the instrumental constants were smaller than 0.20 arcsecond and the astrometry of our catalogue indeed appears to be very good, as confirmed by cross-checking galaxy coordinates with the literature or the APM database (mean average separation: 0.7 arcsecond) and by our multi-object fibre spectroscopy follow-up where the galaxies were always found to be very close to the expected positions."
+ The CCD observations needed to calibrate accurately he Sclunidt pate were not available a hat time., The CCD observations needed to calibrate accurately the Schmidt plate were not available at that time.
+ Hence. a xelininarv shotometzic calibration of hese photographic data has been done using galaxies with known total due imiagnuituce.," Hence, a preliminary photometric calibration of these photographic data has been done using galaxies with known total blue magnitude."
+ The use of catalogued stars was rejected. since such üueh-surface brightucss ojects suffer from OO severe saturation effects (comme both roni the cluulsion itself and from the electronic settigs of the MAMA faciitv|., The use of catalogued stars was rejected since such high-surface brightness objects suffer from too severe saturation effects (coming both from the emulsion itself and from the electronic settings of the MAMA facility).
+ So. LO galaxies available iu he Lyon Extragalacic Database (LEDA. Paturel ct al.," So, 40 galaxies available in the Lyon Extragalactic Database (LEDA, Paturel et al."
+ L997) were selected auLtheir magnitudes compared to their neasured due fluxes providing that no close or overlapping objects wore presen as checked from a sinall scan around cach LEDA ealaxy., 1997) were selected and their magnitudes compared to their measured blue fluxes providing that no close or overlapping objects were present as checked from a small scan around each LEDA galaxy.
+ These 31 undisturbed objects span a 3 uaenitude range. but in a very non uuiforii way.," These 31 undisturbed objects span a 3 magnitude range, but in a very non uniform way."
+ ence. he 3o clipvue routiιο used to compute the best fit Aurther discarded 150 jects.," Hence, the $\sigma$ clipping routine used to compute the best fit further discarded 15 objects."
+" The final ruis on the zero-TOW js unfoxτιlately uot better than 0.1 mag owing οἱ 1) the «uie laree nucertainty quoted for the total nagutude estimate 1i the database along with their irregular distrinition. ii) the limited magnitude range. aud iii) Mfereuces in the 1wolved fux estimates (""plate vs. otal iaenitude)."," The final rms on the zero-point is unfortunately not better than 0.4 mag owing to: i) the quite large uncertainty quoted for the total magnitude estimate in the database along with their irregular distribution, ii) the limited magnitude range, and iii) differences in the involved flux estimates (“plate” vs. total magnitude)."
+ A basic star-galaxy separation has been performed nainvo with respect to a classical surface briehtuess criterion., A basic star-galaxy separation has been performed mainly with respect to a classical surface brightness criterion.
+" As 1sual for glass copies of survey pates. the abilitv of this oulcrion to discriminate drops sharply OY oOjects falter than approximately Y? inaenitude,"," As usual for glass copies of survey plates, the ability of this criterion to discriminate drops sharply for objects fainter than approximately $^{\rm th}$ magnitude."
+ However completeless and purity of anv cataOgue are nost of the time competing eoals., However completeness and purity of any catalogue are most of the time competing goals.
+ So. another test USCC on the cloncation Was then performed 1i orer o reject linear plate flaws. as wel as to pickBN bright elongated. galaxlos first classified as stars «uc O stro18o saturation effects.," So, another test based on the elongation was then performed in order to reject linear plate flaws, as well as to pick bright elongated galaxies first classified as stars due to strong saturation effects."
+ Finally. spurious detectious that occur around very brigh stars (area eyoater thai 10? pixels) due to an incorrect estimate of the local backerouxd were teutatively removed by checking their ocation with respect to these bright objects (the detectk1n processing iucluded no zinoothiug).," Finally, spurious detections that occur around very bright stars (area greater than $^3$ pixels) due to an incorrect estimate of the local background were tentatively removed by checking their location with respect to these bright objects (the detection processing included no smoothing)."
+ Down to he detection nuit. we obtained a list of more than £000 ealaxy caididates over our SRC-J 621 due field to the detection limit.," Down to the detection limit, we obtained a list of more than 4,000 galaxy candidates over our SRC-J 621 blue field to the detection limit."
+ The differential hDunuinosivo distribution of our catalogue of diffuse objects indicaes that this sample is colplete down to the b.= 19.5 magnitude (see Fig.," The differential luminosity distribution of our catalogue of diffuse objects indicates that this sample is complete down to the $b_{\rm
+J}=$ 19.5 magnitude (see Fig."
+ 2j, 2).
+ However. its |rity is less thai the usual level for ligh ealactic laitucle fields (cf.," However, its purity is less than the usual level for high galactic latitude fields (cf."
+ Slezak 1998)., Slezak 1998).
+ Our spectroscopic τιn Πίσας lndicates a contamination level by stars at least three times ligrer (22205€... see Durret 1999).," Our spectroscopic run indeed indicates a contamination level by stars at least three times higher $\simeq$, see Durret 1999)."
+ Such a low suc‘Coss Fate may partly be ascribed to the overall image qlality at the plate corner where, Such a low success rate may partly be ascribed to the overall image quality at the plate corner where
+threshold to be acinittecd to graduate school (i.e.. students who earn equal numbers of A’s ancl Bs in their upper clivision core courses are borderline for most graduate programs).,"threshold to be admitted to graduate school (i.e., students who earn equal numbers of A's and B's in their upper division core courses are borderline for most graduate programs)."
+ We find that iu Physics aud Mathematics no student with SAT; less than roughly 600 was able to attain this level OL mastery., We find that in Physics and Mathematics no student with $_{M}$ less than roughly 600 was able to attain this level of mastery.
+ Figure 11 shows the upper division in-major GPAs of UO physics aud math graduates from ao year period., Figure 11 shows the upper division in-major GPAs of UO physics and math graduates from a 5 year period.
+ The math GPAs were computed using a specilic set of rigorous courses taken by graduate school bound majors., The math GPAs were computed using a specific set of rigorous courses taken by graduate school bound majors.
+ The figure provides at least 1uodest evidence for a minimum cognitive threshold required for physics aud mathematics., The figure provides at least modest evidence for a minimum cognitive threshold required for physics and mathematics.
+ That is. tlie a priori probability that a studeut with SÀT3; « 600 will perform well enough to be adinitted to graduate school in these subjects is extremely low.," That is, the a priori probability that a student with $_{M}$ $<$ 600 will perform well enough to be admitted to graduate school in these subjects is extremely low."
+ Note many poorly performiug studeuts switch majors. aud hence clo uot populate the lower left corner of our graph.," Note many poorly performing students switch majors, and hence do not populate the lower left corner of our graph."
+ A simple but plausible moclel for college performance includes (at least) two factors: 1., A simple but plausible model for college performance includes (at least) two factors: 1.
+ ability (measured by SAT) and 2., ability (measured by SAT) and 2.
+ couscientiousuess or elfort (for simplicity. an uncorrelated random variable. probably normally distributed).," conscientiousness or effort (for simplicity, an uncorrelated random variable, probably normally distributed)."
+ In one version of this model the predicted GPA might depeud linearly on the stun of the two factors. each measurecl in staudard deviation (SD) units. (," In one version of this model the predicted GPA might depend linearly on the sum of the two factors, each measured in standard deviation (SD) units. ("
+This is similar to the GPA predictor depicted in Figure 7.),This is similar to the GPA predictor depicted in Figure 7.)
+" Applying such a moclel to our data. we would conclude that even a student iu the. e.g.. 90th percentile of work ethic lias a low probability of attaining mastery if their SATyy, score is below 600."," Applying such a model to our data, we would conclude that even a student in the, e.g., 90th percentile of work ethic has a low probability of attaining mastery if their $_{M}$ score is below 600."
+ To reiterate. SATa;z GOO seems to be the lowest score at which even a very motivated studeut has a chance for mastery.," To reiterate, $_{M} \approx$ 600 seems to be the lowest score at which even a very motivated student has a chance for mastery."
+ From the data one might guess that only for SAT3; well above TOO do students have more than a 50 percent chance of obtaining GPA > 3.5., From the data one might guess that only for $_{M}$ well above 700 do students have more than a 50 percent chance of obtaining GPA $>$ 3.5.
+" That is. a student with average motivation or couscientiousuess probably needs SAT4, well above 700 to have a high oobability of obtaining mastery."," That is, a student with average motivation or conscientiousness probably needs $_{M}$ well above 700 to have a high probability of obtaining mastery."
+" We were unable to find any similar threshold (either in SAT, or SAT) in other majors. includit[n]ο economics. sociology. history. philosophy. biology. elieuiistry. etc."," We were unable to find any similar threshold (either in $_{R}$ or $_{M}$ ) in other majors, including economics, sociology, history, philosophy, biology, chemistry, etc."
+" For example. Figure 12 is the analogous plot for English. (black) and. History (blue) majors versus SAT, (verbal)."," For example, Figure 12 is the analogous plot for English (black) and History (blue) majors versus $_{R}$ (verbal)."
+" If a hreshold exists it is probably at SAT), of 150 or so.", If a threshold exists it is probably at $_{R}$ of 450 or so.
+ For the total SAT-takine population. an SATa; score of 600 is about 75th percentile.," For the total SAT-taking population, an $_{M}$ score of 600 is about 75th percentile."
+ For he overall population. it might be roughly 85th percentile (the result depends on assumptions concerning the pool of test takers versus the general population).," For the overall population, it might be roughly 85th percentile (the result depends on assumptions concerning the pool of test takers versus the general population)."
+ SAT4; of 750 is roughly 98th ;»ercentile for the total SAT-takine population., $_{M}$ of 750 is roughly 98th percentile for the total SAT-taking population.
+ The University of Oregon is home to Clark Honors College (CHC). oue of the oldest houors colleges in the country.," The University of Oregon is home to Clark Honors College (CHC), one of the oldest honors colleges in the country."
+ Students must apply separately to CHC. αμα aciuission is competitive: enrolled students have average SAT of 1310 aucl unweighted high school CPA of 3.9.," Students must apply separately to CHC, and admission is competitive: enrolled students have average SAT of 1340 and unweighted high school GPA of 3.9."
+ In terms of, In terms of
+Co-orbital Lindblad resonances ean be ignored in investigations of the eccentricities of the orbits of shepherd satellites and narrow planetary rings because the surface density. at the orbit of the shepherd satellite is negligible.,Co-orbital Lindblad resonances can be ignored in investigations of the eccentricities of the orbits of shepherd satellites and narrow planetary rings because the surface density at the orbit of the shepherd satellite is negligible.
+" But because the ratio of vertical thickness to orbital radius is so much larger in protoplanetary disks. 5h/r~0.04. than im planetary vines. h/r<10"". even a massive planet may not clear a very clean gap so interactions ab co-orbital Lindblad resonances could be significant."," But because the ratio of vertical thickness to orbital radius is so much larger in protoplanetary disks, $h/r\sim 0.04$, than in planetary rings, $h/r\lesssim 10^{-6}$, even a massive planet may not clear a very clean gap so interactions at co-orbital Lindblad resonances could be significant."
+ As shown bv equation(14).. these interactions damp (he planets orbital eccentricitv.," As shown by equation, these interactions damp the planet's orbital eccentricity."
+ Attempts (o estimate the rate at which first order co-orbital Lindblad resonances damp eccentricity. are described by and Artvmowiez(1993a)., Attempts to estimate the rate at which first order co-orbital Lindblad resonances damp eccentricity are described by \cite{WAR88} and \cite{ART93b}.
+. The latter concludes that. in a disk of uniform surface density. they damp eccentricity about three times faster than external first order Lindblad resonances excite il.," The latter concludes that, in a disk of uniform surface density, they damp eccentricity about three times faster than external first order Lindblad resonances excite it."
+ An apsidal resonance is a first order inner Lindblad resonance with /=0 and m=1., An apsidal resonance is a first order inner Lindblad resonance with $l=0$ and $m=1$.
+ Its pattern speed is equal to the planets apsidal precession rate: Quy=OQ)—Ayzm.," Its pattern speed is equal to the planet's apsidal precession rate; $\Omega_{0,1}=\Omega_p -\kappa_p =\dot \pomega_p$."
+ As a co-orbital first order Lindblad resonance. it contributes to eccentricity damping.," As a co-orbital first order Lindblad resonance, it contributes to eccentricity damping."
+ Application of the standard. Lindblad resonance torque formula to the apsidal resonance led Ward.&(1998.2000) to conclude that an apsidal resonance damps eccentricity much [faster than external first order Lindblad resonances excite it.," Application of the standard Lindblad resonance torque formula to the apsidal resonance led \cite{WAH98,WAH00} to conclude that an apsidal resonance damps eccentricity much faster than external first order Lindblad resonances excite it."
+ We have described (ος types of resonant planet-disk interactions that damp eccentricitv and one (hat excites il., We have described three types of resonant planet-disk interactions that damp eccentricity and one that excites it.
+ In what follows we argue that eccentricity can grow. provided its inilial value exceeds some minimal threshold and the planet clears a sulliciently clean gap., In what follows we argue that eccentricity can grow provided its initial value exceeds some minimal threshold and the planet clears a sufficiently clean gap.
+ A minimal initial eccentricity is necessary for (he nonlinear saturation of corotation torques., A minimal initial eccentricity is necessary for the nonlinear saturation of corotation torques.
+ This involves a reduction of the surface density gradient at the resonance position., This involves a reduction of the surface density gradient at the resonance position.
+ A clean eap renders inellective damping bv first order co-orbital Lindblad resonances., A clean gap renders ineffective damping by first order co-orbital Lindblad resonances.
+ Although damping due to the apsidal resonance might appear to be the most potent of all. it is likely to be unimportant.," Although damping due to the apsidal resonance might appear to be the most potent of all, it is likely to be unimportant."
+ Apsidal waves have such long wavelengths in protoplanetary disks that they do not propagate., Apsidal waves have such long wavelengths in protoplanetary disks that they do not propagate.
+ Thus torques αἱ apsidal resonances are much smaller than predicted bv the standaid Lindblad resonance torque formula., Thus torques at apsidal resonances are much smaller than predicted by the standard Lindblad resonance torque formula.
+ Apsidal resonances are special because (the doppler shifted. forcing frequency cetines more gradually from the epievelic Irequencey. wilh distance away from an apsidal resonance than it does at other. Lindblad resonances., Apsidal resonances are special because the doppler shifted forcing frequency detunes more gradually from the epicyclic frequency with distance away from an apsidal resonance than it does at other Lindblad resonances.
+ This is why the standard. Lindblad resonance torque formula predicts an apsidal torque, This is why the standard Lindblad resonance torque formula predicts an apsidal torque
+When. ?. audiscovered the first. planet around a main. sequence star. its. small orbital. separation. came as à surprise. to the astronomical. community.. who had expected to find. planets similarD. to those in. our solar svstem.,"When \cite{mayor} discovered the first planet around a main sequence star, its small orbital separation came as a surprise to the astronomical community, who had expected to find planets similar to those in our solar system."
+" uaThis motivated:. work into. how and where these ""hot Jupiters:. D.form and how they might. migrate. to their. current separations. (seefor.example?7)"," This motivated work into how and where these `hot Jupiters' form and how they might migrate to their current separations \citep[see for example][]{mig1,mig2}."
+ Alany close-in. planets are now known and intriguinelv⊀⋠⋠ some planet properties seem to be correlated., Many close-in planets are now known and intriguingly some planet properties seem to be correlated.
+ ο suggest that planet mass is correlated with orbital period. while ? lind that surface gravity may be correlated. with orbital period.," \cite{mazeh} suggest that planet mass is correlated with orbital period, while \cite{pjw} find that surface gravity may be correlated with orbital period."
+ μονο correlations have been re-plotted by. various authors as more planets have been discovered. (for⋅example ??7)..," These correlations have been re-plotted by various authors as more planets have been discovered \citep[for example][]{update1,update2}. ."
+ Both distributions are reproducedI in our Figure. 1.. usingM the sample of transitingUS planets from⋅ this: work (see Section 2.1)).," Both distributions are reproduced in our Figure \ref{correlations}, using the sample of transiting planets from this work (see Section \ref{sampleage}) )."
+ lt has been suggested that these correlations might be . . . . ⋜↧↓⋅∢⊾⊳∖⊔↓↿∪⇂↿↓↥⋖⋅⊔∐⋏∙≟↓⋅⋜∐↓∪⊔⊔↓∢⊾≼⇍↓↕⋜⋯↓⊳∖⊔↓↿∖∢⊾⊳⋏∙≟⊳∙↗⋜↧⊔∠⇂↓⋅⋖⋅⇂⋖⋅↓⋅∢⊾↓⊔∙⋖⋅≱∖⋅ therein): or that they are a result of⋅ evaporation≜ of⋅ close-in≜ planets. such as that reported to be ongoing ..in ILD209458b by ο (seealso??).. ? develop," It has been suggested that these correlations might be a result of the migration mechanism (e.g. \cite{migr} and references therein) or that they are a result of evaporation of close-in planets, such as that reported to be ongoing in HD209458b by \cite{vidal} \citep[see also][]{benjaffel,evapagain}."
+edlevel! a simpleiml modelib of exoplanctpel evaporationapors. driven. by X-rayOPECEUV SUPradiationECT and considered the current evaporation: rates of. known planets., \cite{lecav} developed a simple model of exoplanet evaporation driven by X-ray/EUV radiation and considered the current evaporation rates of known planets.
+ He found∙ that known ralanets will not be losing MENsignificant mass at today's21 ∔↸⊀⊀irradiation rates., He found that known planets will not be losing significant mass at today's irradiation rates.
+ ⋏7. . pointed that N-rav/)IUV BEENirradiation will: be much stronger around outvoung stars and showed that. in principle. the mass distribution of close-in exoplanets could be significantly. modified by evaporation.," \cite{penzXevap} pointed out that X-ray/EUV irradiation will be much stronger around young stars and showed that, in principle, the mass distribution of close-in exoplanets could be significantly modified by evaporation."
+" In this. paper we use a simple. model of⋅ evaporation⊀ to investigate.: whether the trends in. planet properties⋠⋠ identified"" bv ο7. and 2? can be accounted for⋅ by. past evaporation.", In this paper we use a simple model of evaporation to investigate whether the trends in planet properties identified by \cite{mazeh} and \cite{pjw} can be accounted for by past evaporation.
+. We ⇁⋅find good agreement between the distribution∢⊀⊀ of⋅ known planets and the cut-olf⋅ predicted. by the evaporation. model. and we take this as evidence that a population of closc-in exoplanets has been lost to evaporation.," We find good agreement between the distribution of known planets and the cut-off predicted by the evaporation model, and we take this as evidence that a population of close-in exoplanets has been lost to evaporation."
+ We⇁ select only transiting↔ planets for⋅ our studs because we require the planetary. radius in order to consider surface eravity ancl evaporation rates., We select only transiting planets for our study because we require the planetary radius in order to consider surface gravity and evaporation rates.
+ We selected: planets. from ↾↓∖↓∐⋅∣⊲⇀⊲⇀∖∪↓≻⇂⋜⋃⊔⋅⇂⊳∖∣⊲⇀⊲↓↥≼⇍, We selected planets from The Exoplanets on 15 June 2008.
+∙∖⇁≼↛↓∢⋟↓≻⋖⋅∠⇂↕⋜↧↓∪⊔↓⋅↱≻↼∣⊔⊔∢⊾⇉∪∪↖∖⋡↓∖⋡∢⋅∙∖⇁ ∙ . ≻⋜⊔⋜⋯↓⋖⋅↿∢⊾↓⊳∖∖∖⋎⋖⋅↓∢⊾⊳∖∢⊾↓∢⋅≼∼∩⋅∠⇂⇂↓∪⊔↓⇂↓↥∢⊾⋖⋅↓⊔⇍∙∖≼↛↓∪↓≻⋖⋅∠⊔⋜↧⋜↧⊔∠⇂⊳∖∖⊽↓↥∢⊾⊓⋅7 ]. : necessary. ∙⊽∙⋠⊀original discovery⊀nu papers.," Key parameters were selected from the encyclopedia and, where necessary, original discovery papers."
+ opThese parameters. allow he calculation of the energy. incident upon the planet. per unit time. and the total potential energy. of the planet.," These parameters allow the calculation of the energy incident upon the planet per unit time, and the total potential energy of the planet."
+ Stellar; quantities2. are shown in --Table 1.. ancl planetary ∙ ∙⋠∢⊳∖ ⋅≻⋅∙ ≻↓∪↓≻⋖⊾↓↿⊓⊾⊳∖↓⊔↓⋜↧∣⋡↓⋖⊾−⇂∪↓↿↓⊔⋅⊳∖⋜⋯↓↓≻↓⋖⊾∪⇂↓≻↓⋜⋯⋖⋅⇂⊳∖⊔⊳∖⋯⇂↓⊔⇂↓⊔⊳∖ ⋅ . : workX.," Stellar quantities are shown in Table \ref{table2}, and planetary properties in Table \ref{table1} for the sample of planets used in this work."
+ We: used as many of the known transiting↔ svstems as xossible. although some were removed due to the absenceof⋅ data required∙∢∙ in⊀ our∙ analysis.," We used as many of the known transiting systems as possible, although some were removed due to the absenceof data required in our analysis."
+ ---“Phe ColtoT ο...systems ancl Lupus-PR-3b were removed due to lack of distance, The CoRoT systems and Lupus-TR-3b were removed due to lack of distance
+Our model already: provides an accuracy of about 1( Ol serturbative scales aud 105€ on the transition aud lielly ionlinear scales (this is the typical accuracy checked on 3D simulations iu 11).,Our model already provides an accuracy of about $1\%$ on perturbative scales and $10\%$ on the transition and highly nonlinear scales (this is the typical accuracy checked on 3D simulations in \citet{Valageas2011e}) ).
+ Figure { explicitly shows that for 100«(1000 the one-loop correction to the linear )ower is required to reach this accuracy., Figure \ref{fig_Pl_1H} explicitly shows that for $100<\ell<1000$ the one-loop correction to the linear power is required to reach this accuracy.
+ Unfortunately. or (>200 this is not sufficient since higher-order ternis or nou-perturbative corrections cannot be neglected aud he power spectrmu is scusitive to the transition scales. described by the interpolation (151).," Unfortunately, for $\ell > 200$ this is not sufficient since higher-order terms or non-perturbative corrections cannot be neglected and the power spectrum is sensitive to the transition scales, described by the interpolation \ref{P-tang-2}) )."
+ This suggests that or f>500 a percent accuracy is probably bevond the reach of systematic perturbative approaches. which caunot oeo bevond shell crossing (?).. and requires dedicated numerical simulations.," This suggests that for $\ell > 500$ a percent accuracy is probably beyond the reach of systematic perturbative approaches, which cannot go beyond shell crossing \citep{Valageas2011a}, and requires dedicated numerical simulations."
+ We now study the coutvibutious frou various ferius to the lensing bispectrmu, We now study the contributions from various terms to the lensing bispectrum.
+ From Eq.(13)) we now have three contributions. associated with the 3-halo. 2-halo aud 1-halo terms. the 3-halo contribution being ideutified with the perturbative coutribution.," From \ref{Bk-halo}) ) we now have three contributions, associated with the 3-halo, 2-halo and 1-halo terms, the 3-halo contribution being identified with the perturbative contribution."
+ We plot our results for the convergence bispectruni for equilateral configurations iu Fie. 5..," We plot our results for the convergence bispectrum for equilateral configurations in Fig. \ref{fig_Beq_1H},"
+ showing the various contributions iu addition to the full model curve that was already. shown in Fig. 2.., showing the various contributions in addition to the full model curve that was already shown in Fig. \ref{fig_Beq}.
+ As recalled in Sect. 2.1..," As recalled in Sect. \ref{3D-power},"
+ the 3- term is identified with the perturbative prediction and in this paper we use standard perturbation theory at one-loop order for this contribution., the 3-halo term is identified with the perturbative prediction and in this paper we use standard perturbation theory at one-loop order for this contribution.
+ As for the 3D yspectrmm (7).. and contrary to the power spectrum. his eives a contribution that becomes uceligible on small scales. so that it is nof necessary to use a resunuiuatioln scheme or to add a uouperturbative cutoff to ensure a good Μπο] behavior.," As for the 3D bispectrum \citep{Valageas2011e}, and contrary to the power spectrum, this gives a contribution that becomes negligible on small scales, so that it is not necessary to use a resummation scheme or to add a nonperturbative cutoff to ensure a good $\ell$ behavior."
+" Ou the other haucl. contrary to the power ""pectruni we eal sec hat eoie to L-loop order now makes a strong improvement over the trec-order result."," On the other hand, contrary to the power spectrum we can see that going to 1-loop order now makes a strong improvement over the tree-order result."
+ This cature was already noticed for the 3D density field (??)..," This feature was already noticed for the 3D density field \citep{Sefusatti2010,Valageas2011e}."
+ Iu fact. σοιας this l-loop perturbative contribution with the 2-halo aud l-halo texius is sufBcieut to obtain a σου match to the simulations. from the quasilinear to the lughly noulinear scales.," In fact, combining this 1-loop perturbative contribution with the 2-halo and 1-halo terms is sufficient to obtain a good match to the simulations, from the quasilinear to the highly nonlinear scales."
+ This sugeests that higher orders of perturbation theory do not significantly coutribute to the düspectrumni and that we already have a reasonablv Successful model., This suggests that higher orders of perturbation theory do not significantly contribute to the bispectrum and that we already have a reasonably successful model.
+ For these equilateral triangles the 2-halo terni is also subdonüuaut ou all scales (by a factor ~10 , For these equilateral triangles the 2-halo term is also subdominant on all scales (by a factor $\sim 10$ 
+We wish to thank L.Feretti [or critical readiug of the paper. F.Froutera lor the contribution to the PDS data analysis aud the referee J.Nevalaineu for the very useful suggestions.,"We wish to thank L.Feretti for critical reading of the paper, F.Frontera for the contribution to the PDS data analysis and the referee J.Nevalainen for the very useful suggestions."
+We wish to thank L.Feretti [or critical readiug of the paper. F.Froutera lor the contribution to the PDS data analysis aud the referee J.Nevalaineu for the very useful suggestions.M,"We wish to thank L.Feretti for critical reading of the paper, F.Frontera for the contribution to the PDS data analysis and the referee J.Nevalainen for the very useful suggestions."
+We wish to thank L.Feretti [or critical readiug of the paper. F.Froutera lor the contribution to the PDS data analysis aud the referee J.Nevalaineu for the very useful suggestions.MO,"We wish to thank L.Feretti for critical reading of the paper, F.Frontera for the contribution to the PDS data analysis and the referee J.Nevalainen for the very useful suggestions."
+We wish to thank L.Feretti [or critical readiug of the paper. F.Froutera lor the contribution to the PDS data analysis aud the referee J.Nevalaineu for the very useful suggestions.MOD,"We wish to thank L.Feretti for critical reading of the paper, F.Frontera for the contribution to the PDS data analysis and the referee J.Nevalainen for the very useful suggestions."
+where © denotes ensemble average over density fluctuations.,where $\langle\rangle$ denotes ensemble average over density fluctuations.
+ The diffusion coefficient describing this random refraction of waves is where κι=utc]3v. and uo;=Qo/go., The diffusion coefficient describing this random refraction of waves is where $k_\star={2u_0\omega_{pe}}/{3 {\mr v}_{Te}^2}$ and $u_0=\Omega_0/q_0$.
+" Assuming a modest level of turbulence ón/n=107. with a characteristic speed close to the speed of sound in the corona wo=Qo/go~107 em s! and long wavelength 1/qu~10+ em. this spetral ""Siffusion D, substantially changes the spectrumof Langmuir waves in a few collisional times."," Assuming a modest level of turbulence ${\delta n}/n= 10^{-3}$, with a characteristic speed close to the speed of sound in the corona $u_0=\Omega_0/q_0\sim 10^7$ cm $^{-1}$ and long wavelength $1/q_0 \sim 10^{4}$ cm, this spetral diffusion $D_{k}$ substantially changes the spectrumof Langmuir waves in a few collisional times."
+" Figure 4. shows the evolution of the electron distribution function and the spectral energy ""Sensity of Langmuir waves.", Figure \ref{fig:diffusion} shows the evolution of the electron distribution function and the spectral energy density of Langmuir waves.
+ Langmuir waves diffuse in k- space. from an initial value of κWpe/Vp. SO in time the spectrum of Langmuir waves becomes broader as time increases.," Langmuir waves diffuse in $k$ -space, from an initial value of $k\sim\omega_{pe}/{\mr v}_b$, so in time the spectrum of Langmuir waves becomes broader as time increases."
+" The absorption of Langmuir waves by electrons above 20v;, leads to electron acceleration.", The absorption of Langmuir waves by electrons above $20 {\mr v}_{Te}$ leads to electron acceleration.
+ As in the constant gradient case. the integrated electron flux for the fluctuating density case changes considerably at energies above ~20 keV. For density fluctuations with wavenumber g~c/vs the diffusive treatment ts no longer valid. and we must individually consider the interactions between Langmuir waves and other modes in the plasma.," As in the constant gradient case, the integrated electron flux for the fluctuating density case changes considerably at energies above $\sim 20$ keV. For density fluctuations with wavenumber $q \sim \omega_{pe}/{\mr v}_b$ the diffusive treatment is no longer valid, and we must individually consider the interactions between Langmuir waves and other modes in the plasma."
+ In particular. we are interested in those modes which may be excited by the high level of Langmuir waves present due to the beam-plasma interaction.," In particular, we are interested in those modes which may be excited by the high level of Langmuir waves present due to the beam-plasma interaction."
+ Following the treatments in uniform (?) and. inhomogeneous (??) plasmas. we now consider the three wave process. /; 9. where / and / are Langmuir waves and s is a sound wave.," Following the treatments in uniform \citep{2001PhPl....8.3982Z}
+ and inhomogeneous \citep{2002PhRvE..65f6408K,2011ApJ...727...16Z} plasmas, we now consider the three wave process, $ l^{\prime}$; $l\rightarrow l^{\prime}+s$ , where $l$ and $l'$ are Langmuir waves and $s$ is a sound wave."
+ This process efficiently scatters electron beam generated waves with wavenumber k. into secondary Langmuir waves with wavenumber k=—k.," This process efficiently scatters electron beam generated waves with wavenumber ${\bf k}$, into secondary Langmuir waves with wavenumber ${\bf k}^{\prime}\approx {\bf -k}$."
+ The decay of a Langmuir wave generates a ton-sound wave with q=2k., The decay of a Langmuir wave generates a ion-sound wave with ${\bf q}\approx 2{\bf k}$.
+" Every scattering decreases the absolute value of the Langmuir wave number by ki,=2Vin.[niNT*3T;/T.(Ap).", Every scattering decreases the absolute value of the Langmuir wave number by $k^{*}_d=2\sqrt{m_e/m_i}\sqrt{1+3T_i/T_e}/(3\lambda_{De})$.
+ Hence. repeated wave-wave processes cause Langmuir waves to appear at lower κ (larger phase velocities).," Hence, repeated wave-wave processes cause Langmuir waves to appear at lower $k$ (larger phase velocities)."
+ As in the previous cases. Langmuir waves can be absorbed by the tail of the distribution. leading to acceleration of non-thermal electrons.," As in the previous cases, Langmuir waves can be absorbed by the tail of the distribution, leading to acceleration of non-thermal electrons."
+ The simultaneous evolution of both ion-sound waves and Langmuir waves has been already simulated (???.e.g.) and does produce the effect outlined in the previous paragraph.," The simultaneous evolution of both ion-sound waves and Langmuir waves has been already simulated \citep[e.g.]{2001PhPl....8.3982Z,2002PhRvE..65f6408K,2011ApJ...727...16Z}
+ and does produce the effect outlined in the previous paragraph."
+ Here we limit ourself to à more simple but quite illuminating situation with a fixed level of density fluctuations produced by ion-sound waves., Here we limit ourself to a more simple but quite illuminating situation with a fixed level of density fluctuations produced by ion-sound waves.
+ We assume the level of ton-sound waves to be thermal. The results are presented in Fig. 5..," We assume the level of ion-sound waves to be thermal, The results are presented in Fig. \ref{fig:random_nonlin}."
+" As in the cases of constant density gradient and random density fluctuations. we see a time evolving plateau distribution up to ~60vry, at t= Is. The nonlinear decay term has a strong impact on the distribution above 50v;, due to the high level of waves for/60."," As in the cases of constant density gradient and random density fluctuations, we see a time evolving plateau distribution up to $\sim 60{\mr v}_{Te}$ at $t=1$ s. The nonlinear decay term has a strong impact on the distribution above $50{\mr v}_{Te}$ due to the high level of waves for."
+ As the wave-wave terms are particularly effective in shifting the Langmuir wave energy to low &. making Langmuir Waves More Isotropic. the number of accelerated electrons at energies above 100 keV also increases.," As the wave-wave terms are particularly effective in shifting the Langmuir wave energy to low $k$, making Langmuir waves more isotropic, the number of accelerated electrons at energies above $100$ keV also increases."
+" Below 100 keV. the effect in the time integrated spectrum is similar to the previous cases,"," Below $100$ keV, the effect in the time integrated spectrum is similar to the previous cases."
+" The above results show the significant effect that Langmuir waves have on the corresponding electron distribution function,", The above results show the significant effect that Langmuir waves have on the corresponding electron distribution function.
+ A constant density gradient. or random long-wavelength density fluctuations. produce enhancement of the electron distribution in. the range 30-60 νο.," A constant density gradient, or random long-wavelength density fluctuations, produce enhancement of the electron distribution in the range 30-60 ${\mr v}_{Te}$."
+" On the other hand. the Wave-wave interactions with short wavelength fluctuations g~2k show a more pronounced effect at velocities v.—60vy,"," On the other hand, the wave-wave interactions with short wavelength fluctuations $q \sim 2k$ show a more pronounced effect at velocities ${\mr v}\simeq 60{\mr v}_{Te}$."
+ Unlike the case of a uniform plasma. the time averaged electron spectrum is noticeably different from collisional relaxation.," Unlike the case of a uniform plasma, the time averaged electron spectrum is noticeably different from collisional relaxation."
+ As the time averaged spectrum can be directly compared to Hard X-ray observations. the role of Langmuir waves could be observable.," As the time averaged spectrum can be directly compared to Hard X-ray observations, the role of Langmuir waves could be observable."
+" The resulting HXR spectrum above ~20 keV is higher than expected from collisional relaxation. and if interpreted as thick-target relaxation. the number of energetic electrons injected is overestimated,"," The resulting HXR spectrum above $\sim 20$ keV is higher than expected from collisional relaxation, and if interpreted as thick-target relaxation, the number of energetic electrons injected is overestimated."
+ Given the evolution of the electron distribution. f(v.£). we can calculate the effective energy change rate.," Given the evolution of the electron distribution, $f({\mr v},t)$, we can calculate the effective energy change rate."
+" Evidently the actual energy loss/gain rate is given by the kinetic equations (2--5)). but a single particle evolution often used in the literature (e.g.?) the effective energy loss rate (VE/df),j; can be defined as the one satisfying a continuity equation: where f(E.t)=mv(E)f(v(E).t) is the energy distribution function."," Evidently the actual energy loss/gain rate is given by the kinetic equations \ref{eqk1}- \ref{alphaL}) ), but a single particle evolution often used in the literature \citep[e.g.][]{2009A&A...508..993B}
+ the effective energy loss rate $\langle dE/dt\rangle_{eff}$ can be defined as the one satisfying a continuity equation: where $f(E,t)=m{\mr v}(E)f({\mr v}(E),t)$ is the energy distribution function."
+" The effective energy loss/gain rate satisfying this equation is If the evolutionD of particles is governed by collisional only. the effective loss rate is simply (OE/Ot),,;=—Vxn/(2E) (Fig. 6))."," The effective energy loss/gain rate satisfying this equation is If the evolution of particles is governed by collisional only, the effective loss rate is simply $\langle{\partial E}/{\partial t}\rangle_{eff} =-\Gamma\sqrt{m^3/(2E)}$ (Fig. \ref{fig:dE_dt}) )."
+" The evolving Langmuir wave turbulence notably changes (OE/Of),,;. decreasing the effective energy loss at energies of a few keV. increasing the loss rate in the intermediate range near. 10 keV. and leading to acceleration and/or reduced energy loss rate at tens of keV and above."," The evolving Langmuir wave turbulence notably changes $\langle{\partial E}/{\partial t}\rangle_{eff}$, decreasing the effective energy loss at energies of a few keV, increasing the loss rate in the intermediate range near 10 keV, and leading to acceleration and/or reduced energy loss rate at tens of keV and above."
+ This effective energy change rate increases the flux of electrons above 20 keV by a factor of a few to a few tens (shown in Figures 3--5))., This effective energy change rate increases the flux of electrons above 20 keV by a factor of a few to a few tens (shown in Figures \ref{fig:Grad}- \ref{fig:random_nonlin}) ).
+ We emphasize that the acceleration of energetic electrons is due to spectral evolution of Langmuir waves., We emphasize that the acceleration of energetic electrons is due to spectral evolution of Langmuir waves.
+ The Langmuir wave generation in a uniform plasma only weakly changes the time integrated HXR spectrum and. hence. is not likely to be detected in current HXR measurements.," The Langmuir wave generation in a uniform plasma only weakly changes the time integrated HXR spectrum and, hence, is not likely to be detected in current HXR measurements."
+ The important exception is the generation of waves by spatially localized electron beams (e.g. ?).. when electrons are accelerated in bunches rather as a continuous stream as expected in some acceleration models.," The important exception is the generation of waves by spatially localized electron beams \citep[e.g.][]{2011A&A...529A.109H}, , when electrons are accelerated in bunches rather as a continuous stream as expected in some acceleration models."
+ The acceleration of energetic, The acceleration of energetic
+ Warner(1995) LOIM. ~10% (t5) (t4)! , \citet{tome} $^{-4}$ $^{45}$ $_2$ $_3$ 
+Low-iuass clwarl irregular galaxies provide au importait testing ground for several Cuncdamental questions about star formation. star fo‘nation rates (SFRs). ealetie evolution. and cosmology.,"Low-mass dwarf irregular galaxies provide an important testing ground for several fundamental questions about star formation, star formation rates (SFRs), galactic evolution, and cosmology."
+ Due. iu large part. to atteripts to uncerstaid the possible evoltionary connectious between the cwarls with negligible or extr‘emely low presetit SERs (he dSph aud dE galaxies. hereafter dE galaxies) and the dwarfs with obviois slgus of preset star fornation (the cdLIrrs. blue compact dwarfs. HIE galaxies. hereafter dl galaxies). many theorists are turing thelr attentiou to the problem of dwarl galaxy evolutiou.," Due, in large part, to attempts to understand the possible evolutionary connections between the dwarfs with negligible or extremely low present SFRs (the dSph and dE galaxies, hereafter dE galaxies) and the dwarfs with obvious signs of present star formation (the dIrrs, blue compact dwarfs, HII galaxies, hereafter dI galaxies), many theorists are turning their attention to the problem of dwarf galaxy evolution."
+ Envirouiueual effects are trning oit to bea key paameter (e.e.. van den Bergh 199Lb: Ανρί et 11999: Moore et 11999: Cuedcdin 2000: Mayer et 22001a.b: Carraro et 22001).," Environmental effects are turning out to be a key parameter (e.g., van den Bergh 1994b; Klypin et 1999; Moore et 1999; Gnedin 2000; Mayer et 2001a,b; Carraro et 2001)."
+ Aun especially interesing question in liis rega ‘dis the elfect o reionization on the suppression of dwarf galaxy. formation (Efstathiou 1992: Babul Rees 1992> Quinn. atv. Efstathiou 1996: Ba‘kana Loeb 1999: Bullock et 22000).," An especially interesting question in this regard is the effect of reionization on the suppression of dwarf galaxy formation (Efstathiou 1992; Babul Rees 1992; Quinn, Katz, Efstathiou 1996; Barkana Loeb 1999; Bullock et 2000)."
+ By comparing he properties of dwarf galaxies iu different environimmenus we may be able to isolate key euvironiental variables (e.g.. local deusity. companionship. grouον». ecluster inembershliip) in order to corstrain these theories.," By comparing the properties of dwarf galaxies in different environments we may be able to isolate key environmental variables (e.g., local density, companionship, group cluster membership) in order to constrain these theories."
+ The Sculptor group is the closest group of galaxies beyond οir Local Group., The Sculptor group is the closest group of galaxies beyond our Local Group.
+ Originally thought to be at a distauce o. 2.5 Mpe (de Vaucouleurs 1975). the maii ujernmbers of the group are 5 spiral galaxies (NGC 55. NGC 217. NGC 253. NGC 300) aud NGC 7793) with distances raleine between 1.6 and 3.1 Mpc (Puche Carignan 1988 αιd references therein) and a dwarf spial (NCC. 15) at a distance of Lt Mpc.," Originally thought to be at a distance of 2.5 Mpc (de Vaucouleurs 1975), the main members of the group are 5 late-type spiral galaxies (NGC 55, NGC 247, NGC 253, NGC 300, and NGC 7793) with distances ranging between 1.6 and 3.4 Mpc (Puche Carignan 1988 and references therein) and a dwarf spiral (NGC 45) at a distance of 4.4 Mpc."
+ Additionally. NGC 2[. another dwarf spiral. lies close to (GC [25 in radial velocity aud projected location and is uswally included. as a group inember (e.g.. Giuricin et 22000) although some distauce estimates put it as ar away as Ll Mpe (Puche Cariguau 1955).," Additionally, NGC 24, another dwarf spiral, lies close to NGC 45 in radial velocity and projected location and is usually included as a group member (e.g., Giuricin et 2000) although some distance estimates put it as far away as 11 Mpc (Puche Carignan 1988)."
+ Jerjeu. Freeman. BiuggeliMOD (1998) studied surface brightuess fluctuations in 5 Sculptor group dEs aud determiued a spread in cdistauces of dE members of the group from 1.7 to 1. Ape (but note that Jerjen Rejkuba 2001 claim that some of these distances shouk be revised).," Jerjen, Freeman, Binggeli (1998) studied surface brightness fluctuations in 5 Sculptor group dEs and determined a spread in distances of dE members of the group from 1.7 to 4.4 Mpc (but note that Jerjen Rejkuba 2001 claim that some of these distances should be revised)."
+ À, A
+radiation facility at Lawrence Berkeley National Laboratory in California were used to test the accuracy of our computations (see Sect. 3.1.2)).,radiation facility at Lawrence Berkeley National Laboratory in California were used to test the accuracy of our computations (see Sect. \ref{pi_exp}) ).
+ PI cross sections were computed up to 100 Ryd for each ton., PI cross sections were computed up to 100 Ryd for each ion.
+ The length gauge was used at low energies. and AUTOSTRUCTURE automatically switched to velocity gauge at higher energies.," The length gauge was used at low energies, and AUTOSTRUCTURE automatically switched to velocity gauge at higher energies."
+ Because of our use of x-averaged relativistic orbitals. it was not possible to use acceleration gauge.," Because of our use of $\kappa$ -averaged relativistic orbitals, it was not possible to use acceleration gauge."
+ The electron energy mesh was interpolated onto the photon energy mesh using a 4-point Lagrange interpolation., The electron energy mesh was interpolated onto the photon energy mesh using a 4-point Lagrange interpolation.
+ Photoexcitatiolautoionization resonances were included for each system. over a more restricted energy range due to the finer energy mesh required for comparison to experimental measurements.," Photoexcitation-autoionization resonances were included for each system, over a more restricted energy range due to the finer energy mesh required for comparison to experimental measurements."
+ In all cases. we utilized the radial scaling parameters associated with the (N+ |)-electron ton. as tests showed that using the target scaling parameters led to very poor agreement with experimental PI cross sections.," In all cases, we utilized the radial scaling parameters associated with the $N+1$ )-electron ion, as tests showed that using the target scaling parameters led to very poor agreement with experimental PI cross sections."
+ Sealing parameters for continuum orbitals were taken to be equal to those of the highest principal quantum number bound orbitals with the same orbital angular momentum for s. p. and d orbitals. and equal to unity otherwise.," Scaling parameters for continuum orbitals were taken to be equal to those of the highest principal quantum number bound orbitals with the same orbital angular momentum for $s$, $p$, and $d$ orbitals, and equal to unity otherwise."
+ Total and partial final-state resolved radiative recombination rate coefficients. were computed from the direct PI cross sections using detailed balance., Total and partial final-state resolved radiative recombination rate coefficients were computed from the direct PI cross sections using detailed balance.
+ The CI expansions of the (N+D -electron tons were augmented to include one-electron additions to the target’s ground and mixing configurations., The CI expansions of the $N+1$ )-electron ions were augmented to include one-electron additions to the target's ground and mixing configurations.
+ Since it was not possible to use acceleration gauge at high energies. the rate coefficients may be inaccurate at the highest temperatures.," Since it was not possible to use acceleration gauge at high energies, the rate coefficients may be inaccurate at the highest temperatures."
+ However. we have tested that this does not affect the accuracy at phototonized plasma temperatures (see Sect. 3.2.1)).," However, we have tested that this does not affect the accuracy at photoionized plasma temperatures (see Sect. \ref{rr_error}) )."
+ The rate coefficients were determined over the temperature range (10!— 10) K. where z is the charge.," The rate coefficients were determined over the temperature range $^1-10^7$ $z^2$ K, where $z$ is the charge."
+ We used a fixed maximum principal quantum number of 1000. and a maximum orbital angular momentum of 4 for these calculations.," We used a fixed maximum principal quantum number of 1000, and a maximum orbital angular momentum of 4 for these calculations."
+ For/>4. the analytic hydrogenic radial integrals from the recurrence relations of ? were employed.," For $l> 4$, the analytic hydrogenic radial integrals from the recurrence relations of \citet{burgess64} were employed."
+ In photoronized nebulae such as PNe. low-temperature DR is the dominant form of recombination for many species.," In photoionized nebulae such as PNe, low-temperature DR is the dominant form of recombination for many species."
+ We therefore computed DR for An=Q core excitations. neglecting An>| excitations. which are usually negligible at the temperatures of interest (~ 10 K).," We therefore computed DR for $\Delta n=0$ core excitations, neglecting $\Delta n\geq 1$ excitations, which are usually negligible at the temperatures of interest $\sim$ $^4$ K)."
+" However. for DR of Se and Se (forming Se"" and Se. respectively). we allowed core excitations into the 5s and 5p orbitals since the lowest 45° pf 58s and 457 44p 585p - 2.3) levels lie below the lowest 457 4p‘ 44d levels in energy for these ions."," However, for DR of $^+$ and $^{2+}$ (forming $^0$ and $^+$, respectively), we allowed core excitations into the $5s$ and $5p$ orbitals since the lowest $s^2$ $p^k$ $s$ and $s^2$ $p^k$ $p$ $k=2$ ,3) levels lie below the lowest $s^2$ $p^k$ $d$ levels in energy for these ions."
+ Due to the intensive computational requirements of DR calculations. we ignored the (small) contributions to the Se and ὃς” DR rate coefficients from excitations into the 5d. 5f. and ὃς orbitals.," Due to the intensive computational requirements of DR calculations, we ignored the (small) contributions to the $^+$ and $^{2+}$ DR rate coefficients from excitations into the $5d$, $5f$, and $5g$ orbitals."
+ In the case of low-temperature DR. it is critical to have accurate energies for low-energy autoionizing states. in order to determine whether they lie just above the ionization threshold and contribute to DR. or just below threshold in which case they do not.," In the case of low-temperature DR, it is critical to have accurate energies for low-energy autoionizing states, in order to determine whether they lie just above the ionization threshold and contribute to DR, or just below threshold in which case they do not."
+ Unfortunately. the energies of these levels have not been spectroscopically determined for Se or other n--capture elements. thereby limiting the accuracy of our calculated DR rate coefficients.," Unfortunately, the energies of these levels have not been spectroscopically determined for Se or other -capture elements, thereby limiting the accuracy of our calculated DR rate coefficients."
+ To alleviate this source of uncertainty. we used experimental target energies whenever possible (?)..," To alleviate this source of uncertainty, we used experimental target energies whenever possible \citep{badnell06a}."
+ When experimental energies are incomplete for a given term. we assumed the theoretical level splitting.," When experimental energies are incomplete for a given term, we assumed the theoretical level splitting."
+ We did not attempt to adjust the energies of high-energy terms or configurations that had not been experimentally measured. as done by ?..," We did not attempt to adjust the energies of high-energy terms or configurations that had not been experimentally measured, as done by \citet{badnell06a}."
+ We augmented the CI expansions of the structure calculations by including all relevant core excitations for the target ion. allowing for promotions out of the 4s and 4p subshells.," We augmented the CI expansions of the structure calculations by including all relevant core excitations for the target ion, allowing for promotions out of the $4s$ and $4p$ subshells."
+ One-electron additions to all target configurations were added to the (N+1)-electron ion Cl expansions to account for final states for the Rydberg electron to radiate into the core., One-electron additions to all target configurations were added to the $N+1$ )-electron ion CI expansions to account for final states for the Rydberg electron to radiate into the core.
+ We determined the DR rate coefficients over the temperature range (10!—109° K for each state within the ground configuration., We determined the DR rate coefficients over the temperature range $^1-10^7$ $z^2$ K for each state within the ground configuration.
+ Captures into Rydberg states up to 2=1000 and /=200 were computed. with />9 treated hydrogenically for all ions except Se. in which />1l states were treated hydrogenically to allow the calculation to converge.," Captures into Rydberg states up to $n=1000$ and $l=200$ were computed, with $l\geq 9$ treated hydrogenically for all ions except $^{5+}$, in which $l\geq 11$ states were treated hydrogenically to allow the calculation to converge."
+ The photoionization and photoexcitation cross sections from AUTOSTRUCTURE were processed with the ADASPI and ADASPE codes. which respectively produce the and output files (see?.foradetaileddescription)...," The photoionization and photoexcitation cross sections from AUTOSTRUCTURE were processed with the ADASPI and ADASPE codes, which respectively produce the and output files \citep[see][for a detailed description]{summers05}."
+ To simplify the presentation of these results. we used the post-processing to add the direct PL and autoionization cross sections.," To simplify the presentation of these results, we used the post-processing to add the direct PI and photoexcitation-autoionization cross sections."
+" PI cross sections from 0 to 100 Rydbergs for Se""-Se*~ are presented as supplementary data files available at the CDS. and are illustrated in Fig."," PI cross sections from 0 to 100 Rydbergs for $^0$ $^{5+}$ are presented as supplementary data files available at the CDS, and are illustrated in Fig."
+ 1. for PI out of the ground state of each ion., \ref{allpi} for PI out of the ground state of each ion.
+ The photoionization plus photoexcitation-autoionization cross sections (files labeled “seg+_XXPITOT.dat.” g=0— 5) are provided for all states in the ground configuration of each ion.," The photoionization plus photoexcitation-autoionization cross sections (files labeled $q$ XPITOT.dat,” $q=0-5$ ) are provided for all states in the ground configuration of each ion."
+ In those files. column | gives the photon energy relative to the ground state and column 2 the photor energy relative to the initial state.," In those files, column 1 gives the photon energy relative to the ground state and column 2 the photon energy relative to the initial state."
+ Columns 3 and 4 present the photo-electron energy relative to the ground and initial states. respectively. and in column 5 the cross section is given in Mb.," Columns 3 and 4 present the photo-electron energy relative to the ground and initial states, respectively, and in column 5 the cross section is given in Mb."
+ All energies are in Rydberes., All energies are in Rydbergs.
+ The data for different states are delimited by comment lines (denoted by a sign in the first column). followed by three numbers indicating the energy of that level relative to the ground state (or the total energy. in the case of the ground state). the total energy of the target state. and the," The data for different states are delimited by comment lines (denoted by a sign in the first column), followed by three numbers indicating the energy of that level relative to the ground state (or the total energy, in the case of the ground state), the total energy of the target state, and the"
+the Space Telescope Science Institute through the staudarcd pipeline.,the Space Telescope Science Institute through the standard pipeline.
+ Subsequent reductions were carried out at the Main Astronomical Observatory of the Ukrainian Academy of Sciences aud the University of Virginia using andSTSDAS®., Subsequent reductions were carried out at the Main Astronomical Observatory of the Ukrainian Academy of Sciences and the University of Virginia using and.
+ Cosmic rays were removed aud the images in each filler were combined using the CRREJ routine., Cosmic rays were removed and the images in each filter were combined using the CRREJ routine.
+ We found that all subexposures in a giveu filter coregistered to better than  0.2 pixels., We found that all subexposures in a given filter coregistered to better than $\sim$ 0.2 pixels.
+ Fie., Fig.
+ l shows the mosaic V ---uage of NGC 2366., \ref{Fig1} shows the mosaic $V$ image of NGC 2366.
+ We shall follow the notation of in labeling the three main H regious., We shall follow the notation of \citet{Dr00} in labeling the three main H regions.
+ The two brightest oues. I aud IL are on the PC chip. while we see a small part of IHE in the bottom right corner of the WFE2 frame.," The two brightest ones, I and II, are on the PC chip, while we see a small part of III in the bottom right corner of the WF2 frame."
+ The exteuclec low-surlace-brightuess body of the galaxy is visible in all WE aud PC frames., The extended low-surface-brightness body of the galaxy is visible in all WF and PC frames.
+ Several background galaxies can be seen in the field of NGC 2366. the brightest of them being the spiral galaxy to the north of region IL.," Several background galaxies can be seen in the field of NGC 2366, the brightest of them being the spiral galaxy to the north of region II."
+ The superior spatial resolution of the HST/WEPC?2 images combined with the proximity of NGC 2366 permits to resolve the galaxy. into iudividual stars aud study its stellar populatious by means of CALDs., The superior spatial resolution of the /WFPC2 images combined with the proximity of NGC 2366 permits to resolve the galaxy into individual stars and study its stellar populations by means of CMDs.
+ We used the DAOPHOT package in LRAF for point-spread-Duuction (PSFE)-fitting photometry., We used the DAOPHOT package in IRAF for point-spread-function (PSF)-fitting photometry.
+ The PSFs are derived separately for each of the PC aud WE fields aud in each filter uxing the brightest isolated stars in each image., The PSFs are derived separately for each of the PC and WF fields and in each filter using the brightest isolated stars in each image.
+ The background level iu the PC field was measurect in an annulus with radii [ aud 6 pixels (07118 aud 0/228) around each source and. subtracted., The background level in the PC field was measured in an annulus with radii 4 and 6 pixels 18 and 28) around each source and subtracted.
+ The sky level in the WE fields was measured in an annulus with radii 3 and 5 pixels (07330 and Q'551)., The sky level in the WF fields was measured in an annulus with radii 3 and 5 pixels 30 and 51).
+ The PSF-fittiug stellar photometry was done by adopting the zero points of and a 2-pixel fittiug radius for the PC [rames and a 1.5-pixel fitting radius for the WF frames., The PSF-fitting stellar photometry was done by adopting the zero points of \citet{Ho95b} and a 2-pixel fitting radius for the PC frames and a 1.5-pixel fitting radius for the WF frames.
+ The detectability limit is set to the 3e threshold above the local background. noise., The detectability limit is set to the $\sigma$ threshold above the local background noise.
+ To couvert iustrumental magnitudes with au aperture radius of 2 or 1.5 pixels to the magnitudes correspondiug to the Holtzmanοἱal.(1995a) calibrating aperture radius of 07525. (respectively 11 aud 5 pixels for the PC1 and WE frames). we need to derive aperture correctious.," To convert instrumental magnitudes with an aperture radius of 2 or 1.5 pixels to the magnitudes corresponding to the \citet{Ho95a} calibrating aperture radius of 5 (respectively 11 and 5 pixels for the PC and WF frames), we need to derive aperture corrections."
+ For this. we compared PSF-fitted maguitudes of bright isolated stars with the magnitudes of the same stars measured with the aperture photometry technique within an 0755 aperture.," For this, we compared PSF-fitted magnitudes of bright isolated stars with the magnitudes of the same stars measured with the aperture photometry technique within an 5 aperture."
+ For the PC frames we obtained the corrections Vas(0755) — Vay = -0.63 mag and {μωρο ο.) — Igi = -O.77 mae., For the PC frames we obtained the corrections $V_{\rm ap}$ 5) – $V_{\rm fit}$ = –0.63 mag and $I_{\rm ap}$ 5) – $I_{\rm fit}$ = –0.77 mag.
+ For the WE frames the correspoudiug corrections are 20.39 mae in both V. and /., For the WF frames the corresponding corrections are –0.39 mag in both $V$ and $I$.
+ Our derived corrections are very similar to those obtained by Ποσαιetal...(1995a)., Our derived corrections are very similar to those obtained by \citet{Ho95a}.
+. Stars with a sharpness out of the 21.0 — 41.0 rauge in both PC and WF frames were elininated, Stars with a sharpness out of the –1.0 – +1.0 range in both PC and WF frames were eliminated
+It is widely known that the star formation rate (SFR) of ealaxies depends on the density oftheir euviromnoeut. and specifically that galaxies within galaxy clusters tend to be forming stars at a lower rate than comparable field ealaxies.,"It is widely known that the star formation rate (SFR) of galaxies depends on the density of their environment, and specifically that galaxies within galaxy clusters tend to be forming stars at a lower rate than comparable field galaxies."
+ Siuce star formation is fueled by eas. this leads to the conclusion that galaxies iu clusters teud to have less gas from which to form stars than their counterparts in less deuse euvironnments (?).," Since star formation is fueled by gas, this leads to the conclusion that galaxies in clusters tend to have less gas from which to form stars than their counterparts in less dense environments ."
+. Several iiechanisnis have been proposed to explain the observed trend of star formation with environment., Several mechanisms have been proposed to explain the observed trend of star formation with environment.
+ It is well known that spiral galaxies tend to be blur. and therefore forming stars at a hieher rate. than elliptical aud S0 galaxies. the so-called morphology- relation(2).. and that spiral galaxies teud to be rarer iu the centers of galaxy clusters than carly type ealaxies(27).," It is well known that spiral galaxies tend to be bluer, and therefore forming stars at a higher rate, than elliptical and S0 galaxies, the so-called morphology-density relation, and that spiral galaxies tend to be rarer in the centers of galaxy clusters than early type galaxies."
+. It is clear. therefore. that morphological transformation. driven by interactions such as mergers or multiple weak eravitational encounters with other satellites?).. likely plays a role in driviug the observed quenching of star formation iu clusters.," It is clear, therefore, that morphological transformation, driven by interactions such as mergers or multiple weak gravitational encounters with other satellites, likely plays a role in driving the observed quenching of star formation in clusters."
+ However. this caunot completely explain the trend as it has been observed that even among ealaxies with the same bulee to disk ratio. field) galaxies are forming stars at higher rates than cluster galaxies(?).," However, this cannot completely explain the trend as it has been observed that even among galaxies with the same bulge to disk ratio, field galaxies are forming stars at higher rates than cluster galaxies."
+. Cas may be removed frou galaxies passing through the deuse iutergalactie medinu of a cluster by means of rani pressure stripping (RPS)., Gas may be removed from galaxies passing through the dense intergalactic medium of a cluster by means of ram pressure stripping (RPS).
+ As originally euvisaged. in this mechaunisu the interstellay medium (SAD) of a satellite galaxy is lost due to interaction with the dense eas of the host halo through which it is movingη," As originally envisaged, in this mechanism the interstellar medium (ISM) of a satellite galaxy is lost due to interaction with the dense gas of the host halo through which it is moving."
+", Alternatively. a satellite moving through less dense inaterial may lose just a portion of its diffuse atmosphere of hot halo gas?27)."," Alternatively, a satellite moving through less dense material may lose just a portion of its diffuse atmosphere of hot halo gas."
+. Finally. it is inevitable iu a licrarchical Universe that many cluster galaxies were previously micuibers of lower mass groups of galaxies. aud euvironnienutal effects such as straugulation in these lower density euvironmuenuts nay have begun the quenching of the star formation of these galaxies before they were accreted outo the cluster).," Finally, it is inevitable in a hierarchical Universe that many cluster galaxies were previously members of lower mass groups of galaxies, and environmental effects such as strangulation in these lower density environments may have begun the quenching of the star formation of these galaxies before they were accreted onto the cluster."
+. It has traditional]hSD Όσοι. assumed that these environmental effects only occurred ouce satellites had fallen through the virial radius of a host halo., It has traditionally been assumed that these environmental effects only occurred once satellites had fallen through the virial radius of a host halo.
+ However. recent observational studies of satellite ealaxv SFR versus cluser-eentrie radius have found that the SER rendus depressed relative to the field out to 23 times the virial radius of clusters(?77?).," However, recent observational studies of satellite galaxy SFR versus cluser-centric radius have found that the SFR remains depressed relative to the field out to $2-3$ times the virial radius of clusters."
+. Three explanations have been sugeested for the radial extent of environnieutal effects: that egroup-scale effects iu the locality of the cluster are having a large impact on these cluster galaxies(?).. that many satellite galaxies follow lniehlv elliptical orbits which take them close to the cluster center. where they experieuce strong RPS. aud then back out to large radius(?).. and that the radius at which cluster euvironmental effects begius to become inuportant is further out than the virial radius.," Three explanations have been suggested for the radial extent of environmental effects: that group-scale effects in the locality of the cluster are having a large impact on these cluster galaxies, that many satellite galaxies follow highly elliptical orbits which take them close to the cluster center, where they experience strong RPS, and then back out to large radius, and that the radius at which cluster environmental effects begins to become important is further out than the virial radius."
+ Iu fact. it is already known that using the virial radius as the location at which cuvirommental processes related o the cluster begin is ouly an approximation.," In fact, it is already known that using the virial radius as the location at which environmental processes related to the cluster begin is only an approximation."
+ Iu current structure formation scenarios. dark matter halos are ormed through eravitational collapse when an overdeuse watch of dark matter stops expanding with the universe and collapses.," In current structure formation scenarios, dark matter halos are formed through gravitational collapse when an overdense patch of dark matter stops expanding with the universe and collapses."
+ The surrounding gas then falls iuto this xotential well., The surrounding gas then falls into this potential well.
+ If the scale of the halo is huge enough so that the eravitational dvuamical time of the halo is uuch less than the cooling time of the eas. then the eas will form an accretion shock. where the kinetic cnerey of he infalliug eas is converted iuto thermal energy: heating he eas to the virial temperature of the halo(?).," If the scale of the halo is large enough so that the gravitational dynamical time of the halo is much less than the cooling time of the gas, then the gas will form an accretion shock, where the kinetic energy of the infalling gas is converted into thermal energy, heating the gas to the virial temperature of the halo."
+. These are the so-called “virial shocks.” as they are expected to," These are the so-called “virial shocks,"" as they are expected to"
+the power spectrum Iq. (,the power spectrum Eq. (
+4) (see Secion 4.3).,4) (see Section 4.3).
+ The addition of such a correction with a relative power of about 25% leacks already to significant oscillations in t10 correlation function., The addition of such a correction with a relative power of about 25 leads already to significant oscillations in the correlation function.
+ In this Section. we describe and. analyse the correlation functions and the power spectra for a number of ecometrical tov models of the distribution of clusters of galaxies., In this Section we describe and analyse the correlation functions and the power spectra for a number of geometrical toy models of the distribution of clusters of galaxies.
+ Mocels are geometrical in nature: Le. in these models not the power spectrum is given initially but. the distribution of, Models are geometrical in nature; i.e. in these models not the power spectrum is given initially but the distribution of
+"angular momentun is conserved in the eliteh. the torque is where /,,,; Is the moment of inertia of the solid crust plus any component Chat is coupled to it over a timescale in less than the spin-up time .M. and AQ is the observed spin-up of the crust over A/.","angular momentum is conserved in the glitch, the torque is where $I_{crust}$ is the moment of inertia of the solid crust plus any component that is coupled to it over a timescale in less than the spin-up time $\Delta t$ , and $\Delta\Omega$ is the observed spin-up of the crust over $\Delta t$."
+" In magnitude. (he torque is where Fi, is (he azimuthal component of the force that vortices exert on the lattice during ihe exelitch. and δὲ is the averagee distance of vortices from the rotation axis. approximately the stellar radius."," In magnitude, the torque is where $F_\phi$ is the azimuthal component of the force that vortices exert on the lattice during the glitch, and $R$ is the average distance of vortices from the rotation axis, approximately the stellar radius."
+" The azimuthal force is where / is the force exerted per unit length per vortex. // is the average length of a vortex. and NV, is the total number of vortices that are participating in the spin-up."," The azimuthal force is where $f$ is the force exerted per unit length per vortex, $H$ is the average length of a vortex, and $N_v$ is the total number of vortices that are participating in the spin-up."
+" The number of vortices is where n, is the number of vortices per unit area and AR is the thickness of the pinning region (approximately (he thickness of the crust).", The number of vortices is where $n_v$ is the number of vortices per unit area and $\Delta R$ is the thickness of the pinning region (approximately the thickness of the crust).
+ The number density of vortices in a rotating superlhuicl is where ος is the superfluid rotational velocity. approximately equal to the rotational velocity of the crust ancl & is the quantum of circulation ofa neutron fluid (0.0019 /«).," The number density of vortices in a rotating superfluid is where $\Omega_s$ is the superfluid rotational velocity, approximately equal to the rotational velocity of the crust and $\kappa$ is the quantum of circulation ofa neutron fluid (0.0019$^2$ /s)."
+ Pulling these expressions together gives a constraint for f:, Putting these expressions together gives a constraint for $f$ :
+ methanol masers and the environment in which they form with the Long Baseline Array (LBA)., methanol masers and the environment in which they form with the Long Baseline Array (LBA).
+ the emission., the emission.
+the observed: colors e.,the observed colors ${\bm c}$.
+ 1£ there were no A-correction. then the luminosity in a given band would be determined [rom the observed. apparent brightness by the square of the (cosmology dependent) Luminosity distance — the colors are not necessary.," If there were no $k$ -correction, then the luminosity in a given band would be determined from the observed apparent brightness by the square of the (cosmology dependent) luminosity distance – the colors are not necessary."
+ La practice however. one must. apply a 4- correction: this depends on the spectral type of the galaxy. and hence on its color.," In practice however, one must apply a $k$ -correction; this depends on the spectral type of the galaxy, and hence on its color."
+ As a result. the mappingbetween m and AL depends on z and e.," As a result, the mappingbetween $m$ and $M$ depends on $z$ and ${\bm c}$ ."
+ But it is still true that both A? ancl z are determined by e., But it is still true that both $M$ and $z$ are determined by ${\bm c}$.
+ Therefore. the spectroscopic subsample which was previously used to estimate p(z|e) also allows one to estimate pCGM.z|e).The quantity of interest in the previous section. p(z|e). is simply the integral of pOGOM.z|e) over all AZ.," Therefore, the spectroscopic subsample which was previously used to estimate $p(z|{\bm c})$ also allows one to estimate $p(M,z|{\bm c})$.The quantity of interest in the previous section, $p(z|{\bm c})$, is simply the integral of $p(M,z|{\bm c})$ over all $M$."
+ The quantity of interest here. plAlle). is the integral of pCM.z|e) over all z.," The quantity of interest here, $p(M|{\bm c})$, is the integral of $p(M,z|{\bm c})$ over all $z$."
+ Thus. equation (8)) becomes where the second to last expression writes the integral of pCM.Je) over all z as p(Alle). and the final one writes the integral explicitly as a sum over the objects in the catalog.," Thus, equation \ref{NM}) ) becomes where the second to last expression writes the integral of $p(M,z|{\bm c})$ over all $z$ as $p(M|{\bm c})$, and the final one writes the integral explicitly as a sum over the objects in the catalog."
+ The expression above is the convolution-type estimate of NN CAL): tt does not require à. photometric redshift codo., The expression above is the convolution-type estimate of $N(M)$ ; it does not require a photometric redshift code.
+ llowever. in principle. a photometric redshift code could out put.Z(AM.z|e): the quantity such codes currently output. C(z|e). is the integral of CCV.το) over all AJ.," However, in principle, a photometric redshift code could output ${\cal L}(M,z|{\bm c})$: the quantity such codes currently output, ${\cal L}(z|{\bm c})$, is the integral of ${\cal L}(M,z|{\bm c})$ over all $M$."
+ The relevant weighted sum becomes where Z(A|e) is the integral of ΟΓΑ.z]e) over all z. the sum is over all the objects in the catalog. and the method only works if ZCM]|e)=plMe).," The relevant weighted sum becomes where ${\cal L}(M|{\bm c})$ is the integral of ${\cal L}(M,z|{\bm c})$ over all $z$, the sum is over all the objects in the catalog, and the method only works if ${\cal L}(M|{\bm c}) = p(M|{\bm c})$."
+ Note that the luminosity density (in solar units) can. therefore. be written as The second to last line shows that one requires the average of GL/Vinax(£))) summed over the distribution pOMAA): this is easily. computed from. distributions like those shown in the bottom. panel of Figure 3...," Note that the luminosity density (in solar units) can, therefore, be written as The second to last line shows that one requires the average of $\langle L/V_{\rm max}(L)\rangle$ summed over the distribution $p(M|{\cal M})$; this is easily computed from distributions like those shown in the bottom panel of Figure \ref{MzMphoto}."
+ The final expression writes this as a sum over the observed distribution of colors., The final expression writes this as a sum over the observed distribution of colors.
+ Although the previous section considered the Iuminosity function in a single band. it is clear that the photometric redshift’ codes: could: output CUM.:|e) where M is a set of absolute luminosities (typically. these will be those associated with the various band. passes fron: which the colors e were determined).," Although the previous section considered the luminosity function in a single band, it is clear that the photometric redshift codes could output ${\cal L}({\bm M},z|{\bm c})$, where ${\bm M}$ is a set of absolute luminosities (typically, these will be those associated with the various band passes from which the colors ${\bm c}$ were determined)."
+ Hence. the color magnitude relation. which is really a statement about the joint distribution in two bands. can be estimated by Galaxy scaling relations can be estimated: similarly. if we simply interpret. AZ as being the vector of observables which can include sizes. ete. (," Hence, the color magnitude relation, which is really a statement about the joint distribution in two bands, can be estimated by Galaxy scaling relations can be estimated similarly, if we simply interpret ${\bm M}$ as being the vector of observables which can include sizes, etc. ("
+not. just. Iuminosities).,not just luminosities).
+ In principle. quantities other than colors. (e.g. apparent magnitudes. surface brightness. axis ratios) can plav a role in the photometric redshift determination: this can be incorporated into the formalism simply by using ο to now denote the full set of observables from. which the redshift and other intrisic quantities AY were estimated.," In principle, quantities other than colors (e.g., apparent magnitudes, surface brightness, axis ratios) can play a role in the photometric redshift determination; this can be incorporated into the formalism simply by using ${\bm c}$ to now denote the full set of observables from which the redshift and other intrisic quantities ${\bm M}$ were estimated."
+ If one wishes to use the output from a photo-z. codo. rather than from the spectroscopic subset. one would use having checked that. in the spectroscopic subset. CUM|e;)=pCM|e;).," If one wishes to use the output from a $z$ code, rather than from the spectroscopic subset, one would use having checked that, in the spectroscopic subset, ${\cal L}({\bm M}|{\bm c}_i) = p({\bm M}|{\bm c}_i)$."
+ We showed how previous work on deconvolution algorithms [or making unbiased reconstructions of galaxy distributions and scaling relations (Sheth 2007: Rossi Sheth 2008: IHossi et al., We showed how previous work on deconvolution algorithms for making unbiased reconstructions of galaxy distributions and scaling relations (Sheth 2007; Rossi Sheth 2008; Rossi et al.
+ 2010) could be related to convolution-based moethocks., 2010) could be related to convolution-based methods.
+ Whereas deconvolution based methods. require accurate knowledge of p(c|z). the distribution of the photometric redshift ¢ given the true redshift 2. convolution based methods require accurate knowledge of p(z|C).," Whereas deconvolution based methods require accurate knowledge of $p(\zeta|z)$, the distribution of the photometric redshift $\zeta$ given the true redshift $z$, convolution based methods require accurate knowledge of $p(z|\zeta)$."
+ Since C is derived. from photometry. this may more generally be written às p(z|e). where e is the vector of observed photometrie parameters which were used to estimate the redshift.," Since $\zeta$ is derived from photometry, this may more generally be written as $p(z|{\bm c})$, where ${\bm c}$ is the vector of observed photometric parameters which were used to estimate the redshift."
+ In both cases. p(z|e) and p(£|z) are calibrated from a sample in which z is known. and are then used in a larger sample where z is not available.," In both cases, $p(z|{\bm c})$ and $p(\zeta|z)$ are calibrated from a sample in which $z$ is known, and are then used in a larger sample where $z$ is not available."
+ Lf the smaller training set has the same selection limits as the larger dataset (e.g. both have the same magnitude limit) then both approaches are valid.," If the smaller training set has the same selection limits as the larger dataset (e.g., both have the same magnitude limit) then both approaches are valid."
+ We illustrated our arguments with measurements in the SDSS (Figures 1- 4))., 	We illustrated our arguments with measurements in the SDSS (Figures \ref{pzzeta}- \ref{NMconv}) ).
+ We also showed what acelitional information must be output. from photometric recdshift codes if their. results are to be used in à convolution-like approach to. provide unbiased. estimates of galaxy sealing relations., We also showed what additional information must be output from photometric redshift codes if their results are to be used in a convolution-like approach to provide unbiased estimates of galaxy scaling relations.
+ In particular. we argued that only if the redshift distribution output bv a photo-z algorithm. Z(z|e). has thesame shape as p(z]e). can the algorithm be said to be unbiased.," In particular, we argued that only if the redshift distribution output by a $z$ algorithm, ${\cal L}(z|{\bm c})$ , has thesame shape as $p(z|{\bm c})$ , can the algorithm be said to be unbiased."
+ Only inthis case its output (available for the full sample) can be used in place of pte) Ovhich is typically available for à small subset)., Only inthis case its output (available for the full sample) can be used in place of $p(z|{\bm c})$ (which is typically available for a small subset).
+ The safest wav to accomplish this is for the training set to be, The safest way to accomplish this is for the training set to be
+QSOs would be luminous submillimeter sources. but Cais is ruled out by our observations.,"QSOs would be luminous submillimeter sources, but this is ruled out by our observations."
+ If they were absorbed (o a similar degree as the N-rayv. absorbed QSOs. they would outinunber (his population substantially at the flux limits probed by the hard-spectrum.osa! survey. and certainlv would be found in huge numbers just below this [lux limit in the deeper survevs (hat are now being carried out with andChandra.," If they were absorbed to a similar degree as the X-ray absorbed QSOs, they would outnumber this population substantially at the flux limits probed by the hard-spectrum survey, and certainly would be found in large numbers just below this flux limit in the deeper surveys that are now being carried out with and."
+ This is not the case (Alainierietal.2002:Page2003).. and hence the earlier evolutionary. phases niust be characterised by considerably higher levels of X-ray absorption Vj>107*em.7).," This is not the case \citep{mainieri02,page03}, and hence the earlier evolutionary phases must be characterised by considerably higher levels of X-ray absorption $\nh > 10^{23} {\rm cm}^{-2}$ )."
+ Since the great majority. of optically selected. QSOs show little or no N-rav. absorption 1993).. the earlier phases must also be heavily obscured in the restframe optical/UV.," Since the great majority of optically selected QSOs show little or no X-ray absorption \citep{yuan98}, the earlier phases must also be heavily obscured in the restframe optical/UV."
+ Consequently. AGN in the earliest stages of their evolution will be diffieult to detect in anv waveband. except for the far-IR/submillimeter where they will be strong sources.," Consequently, AGN in the earliest stages of their evolution will be difficult to detect in any waveband, except for the far-IR/submillimeter where they will be strong sources."
+ At present. the small number of optically faint. hard. X-ray. detected submillimeter sources reported by Ivisonetal.(2002) and. Alexanderetal.(2003) are probably the best candidates for AGN in (his initial phase because {μον are more heavily obscured (han our Roseal-selected objects. alihough we can anticipate that future. high-throughput. A-ray observatories could. reveal ihe buried AGN components in many more submillimeter sources.," At present, the small number of optically faint, hard X-ray detected submillimeter sources reported by \citet{ivison02} and \citet{alexander03} are probably the best candidates for AGN in this initial phase because they are more heavily obscured than our -selected objects, although we can anticipate that future high-throughput X-ray observatories could reveal the buried AGN components in many more submillimeter sources."
+ In contrast. the X-ray absorbed QSOs [ound in the hard-spectrum. survey may be caught al the earliest evolutionary stage reaclily accessible to conventional X-ray. aud optical surveys: flourishing and vouthful. they may only just be emerging from the obscuring gaseous cocoons of {heir birth.," In contrast, the X-ray absorbed QSOs found in the hard-spectrum survey may be caught at the earliest evolutionary stage readily accessible to conventional X-ray and optical surveys; flourishing and youthful, they may only just be emerging from the obscuring gaseous cocoons of their birth."
+ The James Clerk Maxwell Telescope is operated on behalf of the Particle Physics aud Astronomy Research Council of the United. Ixingdom. the Netherlands Organisation [or scientific Research and the National Research Council of Canada.," The James Clerk Maxwell Telescope is operated on behalf of the Particle Physics and Astronomy Research Council of the United Kingdom, the Netherlands Organisation for Scientific Research and the National Research Council of Canada."
+ FIC acknowledge financial support by the Spanish Ministerio de Ciencia v iaa. under grant. ESP2003-00812.," FJC acknowledge financial support by the Spanish Ministerio de Ciencia y a, under grant ESP2003-00812."
+There are more (han 200 photometric svstems that have been used in astronomy therein)..,There are more than 200 photometric systems that have been used in astronomy \citep[][and references therein]{bessell05}.
+ The amount of information about à star that can be extracted from photometry is hiehlv dependent on the choice of photometric filters., The amount of information about a star that can be extracted from photometry is highly dependent on the choice of photometric filters.
+ Early on. Stronuneren introduced the σου system with the aim of characterizing (he main stellar atmospheric parameters and reddening (Stvomeren1951.1956).," Early on, Strömmgren introduced the $uvby$ system with the aim of characterizing the main stellar atmospheric parameters and reddening \citep{stromgren51,stromgren56}."
+. StvGmuueren’s filters have widths of the order of 200A.. and are thus considered intermediate-band.," Strömmgren's filters have widths of the order of 200, and are thus considered intermediate-band."
+ Other svstems with similar anc, Other systems with similar and
+ Other svstems with similar ancl, Other systems with similar and
+The entire 6 cimi field is covered by V aud I CCD data collected at the Canada-Frauce-Hawaii Telescope (CFUT) with the UIISI& camera (see Wilson et al. (,The entire 6 cm field is covered by V and I CCD data collected at the Canada-France-Hawaii Telescope (CFHT) with the UH8K camera (see Wilson et al. (
+2001) aud Isaiscr et al. (,2001) and Kaiser et al. (
+2001) for a description of these data).,2001) for a description of these data).
+ The approximate limits in the two bands are V—25.5 mae aud I~21.5 mag OY point sources (Vega - maguitudoe)., The approximate limits in the two bands are $\sim$ 25.5 mag and $\sim$ 24.5 mag for point sources (Vega - magnitude).
+ The CCD field is 30%«30° around the ROSAT ultra deep TRI aud 6 cin fick center., The CCD field is $^{\prime} \times 30^{\prime}$ around the ROSAT ultra deep HRI and 6 cm field center.
+ Moreover. a mosaic of I'-baud images (frou the Omega-Prime camera ou the Calw Alto 3.5-1n telescope) covering the feld in a 1011”uniforii fashion with gaps in between is also available (see Leliuaun et al. (," Moreover, a mosaic of $^{\prime}$ -band images (from the Omega-Prime camera on the Calar Alto 3.5-m telescope) covering the field in a non-uniform fashion with gaps in between is also available (see Lehmann et al. ("
+2000) and Schunidt et al. (,2000) and Schmidt et al. (
+1998) for a sununary of the opticd available data in the Lockiau Tole).,1998) for a summary of the optical available data in the Lockman Hole).
+ The liauiting maguitudes of the I images is K'~ 20.2 imag., The limiting magnitudes of the $^{\prime}$ images is $^{\prime} \sim$ 20.2 mag.
+ The typical photometric error on the V. I aud IX’ maeuitudes is —0.1 mae.," The typical photometric error on the V, I and $^{\prime}$ magnitudes is $\sim$ 0.1 mag."
+ For the optical identification of the 6 ci radio sources. we used the likelihood ratio technique described bv Sutherlaud Sauuders (1992).," For the optical identification of the 6 cm radio sources, we used the likelihood ratio technique described by Sutherland Saunders (1992)."
+ The mean off-set between he radio aud optical positions estimated in Sect., The mean off-set between the radio and optical positions estimated in Sect.
+ 3.2 has con subtracted from the radio positions to compute the yositional offset1., 3.2 has been subtracted from the radio positions to compute the positional offset.
+ For a given optical candidate with uaeuitude auc positional offsetr frou the radio source vosition. we calculate the probabilityp that the true source Hes in anu infinitesimal box drí2. and in a uaeuitude interval dii/2: on the asstuuption that the xositional offsets are independent of theoptical properties às given by: where is the probability distribution function of the positional errors. assumed to be equal iu the two coordinates. with and is the expected distribution as a function of magnitude of the optical counterparts.," For a given optical candidate with magnitude and positional offset from the radio source position, we calculate the probability that the true source lies in an infinitesimal box /2, and in a magnitude interval /2; on the assumption that the positional offsets are independent of the properties is given by: where is the probability distribution function of the positional errors, assumed to be equal in the two coordinates, with and is the expected distribution as a function of magnitude of the optical counterparts."
+ The likelihood ratio is defined as the ratio between the probability that the source is the correct identification and the correspouding probability for a backerouud. unrelated object: where is the surface density of background objects with magnitude2.," The likelihood ratio is defined as the ratio between the probability that the source is the correct identification and the corresponding probability for a background, unrelated object: where is the surface density of background objects with magnitude."
+ The preseuce or absence of other optical candidates for the sue radio source provides to that coutaimed inLA., The presence or absence of other optical candidates for the same radio source provides to that contained in.
+ Thus a selfconsistent formula has been developed for the reliability of cach individual object. taking this information iuto account (Sutherland Saunders. 1992).," Thus a self-consistent formula has been developed for the reliability of each individual object, taking this information into account (Sutherland Saunders, 1992)."
+" The reliability Rel; for object boiug the correct identification is: where the sum is over the set of all candidates for this particular source and Q is the probability that the optical counterpart of the source is brighter than the magnitude limit of the optical catalogue (Q—[77qim)din),", The reliability $Rel_j$ for object being the correct identification is: where the sum is over the set of all candidates for this particular source and Q is the probability that the optical counterpart of the source is brighter than the magnitude limit of the optical catalogue $Q = \int^{m_{lim}} q(m)~dm $ ).
+ Iu order to derive an estimate for we have first counted all objects iu the optical catalogue within a fixed radius around cach source (£fofal(m]))., In order to derive an estimate for we have first counted all objects in the optical catalogue within a fixed radius around each source ).
+ This distribution has then been background subtracted aud normalized to construct the distribution function of ‘real’ identifications: Mj is the sun over alb the magnitude bius of the distribution. £c. is the total nuniber of objects iu the real(ia) distribution.," This distribution has then been background subtracted and normalized to construct the distribution function of `real' identifications: where $\Sigma_i$ is the sum over all the magnitude bins of the distribution, $i.e.$ is the total number of objects in the $real(m)$ distribution."
+ Tn order to maximize the statistica senificauce of the over deusitv due to the presence of the optical counterparts. we have adopted for tle radius he minimal size Which assures that most of the possible counterparts are ducluded within such radius.," In order to maximize the statistical significance of the over density due to the presence of the optical counterparts, we have adopted for the radius the minimum size which assures that most of the possible counterparts are included within such radius."
+ Ou the basis of the positional uucertaimties shown in Fig. 6..," On the basis of the positional uncertainties shown in Fig. \ref{F6},"
+ we ive adopted a radius of 2 arcsec., we have adopted a radius of 2 arcsec.
+ Fig., Fig.
+ 10. shows the resulting distribution (dotted. line) together with the expected distribution of background objects (solid line). ἐς objects unrelated to the radio sources.," \ref{F10} shows the resulting distribution (dotted line) together with the expected distribution of background objects (solid line), $i.e$ objects unrelated to the radio sources."
+ The smooth curve fitted to 16 distribution (dot dot dot dashed line) has been jormalized according to Eq. (, The smooth curve fitted to the distribution (dot dot dot dashed line) has been normalized according to Eq. (
+9) aud then used as input iu 1e likelihood. caleulation.,9) and then used as input in the likelihood calculation.
+ The total uuuber of objects in je distribution is [8.58 d 8.7., The total number of objects in the distribution is 48.8 $\pm$ 8.7.
+ Since the ΠΟ MX radio sources is 623. this corresponds to an expected yaction of identificatious above the magnitude Πιτ of 16 optical catalogue of the order of (77 + 11)," Since the number of radio sources is 63, this corresponds to an expected fraction of identifications above the magnitude limit of the optical catalogue of the order of (77 $\pm$ 14)."
+ Ou this sis we adopted Q = 0.5., On this basis we adopted $Q$ = 0.8.
+ This value is also in agreement with the results obtained by Tanuner et al. (, This value is also in agreement with the results obtained by Hammer et al. (
+1995) for the radio sources in the Canada-Frauce Redshift Survey aud w Fomalout et al. (,1995) for the radio sources in the Canada-France Redshift Survey and by Fomalont et al. (
+01997) aud Richards et al. (,1997) and Richards et al. (
+1998) for he radio sources in the Hubble Deep Field (IDF).,1998) for the radio sources in the Hubble Deep Field (HDF).
+ Using I baud CCD data with a magnitude lini of I~ 21.5 nag Iauuner et al., Using I band CCD data with a magnitude limit of I $\sim$ 24.5 mag Hammer et al.
+ found 32 optical ideutificatious for 36 racdio sources with 55.9eu2 16543. while Richards et al.," found 32 optical identifications for 36 radio sources with $_{\rm 5.0~GHz}>16\mu$ Jy, while Richards et al."
+ found 26 optical counterparts down to I-225 mae for 29 radio sources with S45ον> 805v., found 26 optical counterparts down to I=25 mag for 29 radio sources with $_{\rm 8.5~GHz}>9\mu$ Jy.
+ However. to check how this assuuptiou could affect our results. we repeated the likelihood ratio analysis using different values of Q in the ranee 0.51.0.," However, to check how this assumption could affect our results, we repeated the likelihood ratio analysis using different values of $Q$ in the range 0.5–1.0."
+ No substautial, No substantial
+02cm (PO =j,symmetries.
+"Qgoeo« = PorQeo'. 0.2em HPO =Poo'o* =jQQ"". 0.2em Q.2cm With the indices «.5.c... ranging rom 1to3 weget."," However, these representations coincide with the usual ones only when applied to special combinations of quark variables, cubic (spinor) or quadratic (vector) representations of the Lorentz group."
+ eight linearly independent combinations of ," While acting on quark variables, the representations correspond to the $Z_3$ -covering of groups."
+threeundotted indices. and thesamen," In this sense quarks are not like ordinary spinors or fermions, and as such, do not obey the usual Dirac equation."
+"umber of combinations ofdotted Theycanbe arranged asfollows: 0 em BORO"" Dou? Qui! 0.2em QOQQu QQo. qoo 0.2cm HPO."," If the sigma-matrices are to be replaced by the non-hermitian matrices $\pi^{\mu}_{A {\dot{B}}}$, instead of the usual wave-like solutions of Dirac's equation we shall get the exponentials of complex wave vectors, and such solutions cannot propagate."
+ qQ'quu. 0.2cm while finite algebra generaterd by Qos.The invariant 3-Jormmapping these combinations ontos," Nevertheless, as argued in \cite{Kerner5}, certain tri-linear and bi-linear combinations of such solutions behave as usual plane waves, with real wave vectors and frequencies, if there is a convenient coupling of non-propagating solutions in the $k$ -space."
+ome eieht-dimensional, 0.2cm We are greatly indebted to Michel Dubois-Violette for numerous discussions and enlightening remarks.
+"luminosities, and we might expect them to be located in more massive dark matter halos.","luminosities, and we might expect them to be located in more massive dark matter halos."
+" We note, however, that the studies of optically selected QSOs also indicate a weak or no trend of clustering length with the object luminosity (??)."," We note, however, that the studies of optically selected QSOs also indicate a weak or no trend of clustering length with the object luminosity ."
+. Clustering analysis of the SDSS quasars shows mild or no evolution of the real-space correlation length at zx2., Clustering analysis of the SDSS quasars shows mild or no evolution of the real-space correlation length at $z\lesssim2$.
+" Having estimated the mass scale and therefore the space density of AGNs, we can formally compute their duty cycle following the approach of?."," Having estimated the mass scale and therefore the space density of AGNs, we can formally compute their duty cycle following the approach of."
+". Under the simplifying assumption that the halo lifetime is approximately independent of mass, the probability for the AGN to be active is simply where nagn is obtained from the dN/dz fit (eq. 2))"," Under the simplifying assumption that the halo lifetime is approximately independent of mass, the probability for the AGN to be active is simply where $n_{\rm AGN}$ is obtained from the $dN/dz$ fit (eq. \ref{frm:dndz}) )"
+ for the median redshift for the given subsample., for the median redshift for the given subsample.
+" The limiting X-ray luminosity is ill-defined for our sample because the X-ray detections extend to very low limits in terms of the number of detected X-ray photons?),, and in this regime, a wide range of possible intrinsic intensities to the given number of photons(?)."," The limiting X-ray luminosity is ill-defined for our sample because the X-ray detections extend to very low limits in terms of the number of detected X-ray photons, and in this regime, a wide range of possible intrinsic intensities corresponds to the given number of photons."
+". However, as a guide correspondsfor the typical luminosity scale one can use the median logL, reported in Table 1.."," However, as a guide for the typical luminosity scale one can use the median $\log L_{x}$ reported in Table \ref{tab:values}."
+" With these caveats, the probabilities given by eq."," With these caveats, the probabilities given by eq."
+" 19 and reported in column (9) of Table 1 correspond to the probability for the dark matter halos with vg;>Vmin to host an AGN with the instantaneous soft-band X-ray luminosity of order (logL,) or above.", \ref{eq:duty:cycle} and reported in column (9) of Table \ref{tab:values} correspond to the probability for the dark matter halos with $\vmax>\Vmin$ to host an AGN with the instantaneous soft-band X-ray luminosity of order $\langle\log L_{x}\rangle$ or above.
+ These Pactive generally decline with mean luminosity and/or redshift., These $P_{\rm active}$ generally decline with mean luminosity and/or redshift.
+" Interestingly, our results indicate that AGNs are quite common at low redshifts — approximately 10% of dark matter halos with Vmax>280km (or mass Miso>3.7x10?17! Mo) host an AGN with a soft-band X-ray luminosity of ~4x10” or above."," Interestingly, our results indicate that AGNs are quite common at low redshifts — approximately $10\%$ of dark matter halos with $\vmax>280\,$ (or mass $\M200>3.7\times10^{12}\,h^{-1}\,M_{\odot}$ ) host an AGN with a soft-band X-ray luminosity of $\sim 4\times10^{42}\,$ or above."
+ We measured the clustering properties of X-ray selected AGNs using a sample of 1282 sources with spectroscopic redshifts from the 9.3 deg? survey in the Boóttes region — one of the most accurate such measurement to date., We measured the clustering properties of X-ray selected AGNs using a sample of 1282 sources with spectroscopic redshifts from the 9.3 $^2$ survey in the Boöttes region — one of the most accurate such measurement to date.
+" In agreement with previous studies of X-ray and optically selected AGNs, we find that the real-space correlation function can be approximated in the radial range 1—20/7! a power law with a of y=1.84+0.12."," In agreement with previous studies of X-ray and optically selected AGNs, we find that the real-space correlation function can be approximated in the radial range $1-20\,h^{-1}$ comoving Mpc by a power law with a slope of $\gamma=1.84\pm0.12$."
+" The comovingcorrelationMpc bylength is ~6.4/7! slopeMMpc, showing only weak trends with redshift at z=0.17—1.6 (or with X-ray luminosity)."," The correlation length is $\approx
+6.4\,h^{-1}$ Mpc, showing only weak trends with redshift at $z=0.17-1.6$ (or with X-ray luminosity)."
+" Matching the observed of the AGNs to those of dark clusteringmatter propertieshalos in the high- cosmological simulations, we find that the X-ray AGNs reside in halos with the maximum rotational velocity z320kms!,, or with total masses ~4.1x10!7!Mo, also with no detectable redshift trend."," Matching the observed clustering properties of the AGNs to those of dark matter halos in the high-resolution cosmological simulations, we find that the X-ray AGNs reside in halos with the maximum rotational velocity $\approx 320\,$, or with total masses $\sim 4.1\times10^{12}\,h^{-1}\,M_{\odot}$, also with no detectable redshift trend."
+ The lack of a redshift or luminosity dependence of the AGN clustering is inconsistent with the scenarios in which the AGN luminosities in the active state are similar fractions of the Eddington luminosity., The lack of a redshift or luminosity dependence of the AGN clustering is inconsistent with the scenarios in which the AGN luminosities in the active state are similar fractions of the Eddington luminosity.
+" However, it can be explained in the scenario in which the AGN activity is triggered by major mergers of gas-rich galaxies, and the instantaneous luminosity passes through many levels after each trigger(?)."," However, it can be explained in the scenario in which the AGN activity is triggered by major mergers of gas-rich galaxies, and the instantaneous luminosity passes through many levels after each trigger."
+. Our results reveal another interesting aspect of the AGN clustering which was predicted in?., Our results reveal another interesting aspect of the AGN clustering which was predicted in.
+. The redshift measurements in our sample are sufficiently accurate to detect peculiar motions of objects in excess of ~100kms!.," The redshift measurements in our sample are sufficiently accurate to detect peculiar motions of objects in excess of $\sim 100\,$."
+". The comparison of the two-point correlation functions projected on the line of sight and on the sky plane reveals no signatures of the redshift-space distortions, which allows us to put limits on the fraction of AGNs located in the satellite subhalos the host dark matter halos."," The comparison of the two-point correlation functions projected on the line of sight and on the sky plane reveals no signatures of the redshift-space distortions, which allows us to put limits on the fraction of AGNs located in the satellite subhalos the host dark matter halos."
+ We find that the X-ray AGNS are predominantly located in the central galaxies of the host dark matter halos and tend to satellite galaxies., We find that the X-ray AGNs are predominantly located in the central galaxies of the host dark matter halos and tend to satellite galaxies.
+" Quantitatively, we limit the fraction of AGNs in non-central galaxies to be «0.12 at the CL."," Quantitatively, we limit the fraction of AGNs in non-central galaxies to be $<0.12$ at the CL."
+ We also exclude the model in which the probability for a galaxy with Vmax>300 to host an AGN is the same for central galaxies and satellite galaxies in more massive halos ," We also exclude the model in which the probability for a galaxy with $\vmax>300\,$ to host an AGN is the same for central galaxies and satellite galaxies in more massive halos \\ref{sec:res:vmin-f}) )."
+The central locations of the quasar host galaxies are expected in the trigger model because mergers of equally-sized galaxies preferentially occur at the centers of dark matter halos(?)., The central locations of the quasar host galaxies are expected in the trigger model because mergers of equally-sized galaxies preferentially occur at the centers of dark matter halos.
+". Finally, we compared the number densities of the AGNs to that of the dark matter halos with the mass corresponding to the AGN clustering amplitude."," Finally, we compared the number densities of the AGNs to that of the dark matter halos with the mass corresponding to the AGN clustering amplitude."
+" We find that the fraction of halos with active X-ray AGNs decreases with increasing z — and, correspondingly, with L, — probably reflecting a lower probability for an object to have a instantaneous luminosity."," We find that the fraction of halos with active X-ray AGNs decreases with increasing $z$ — and, correspondingly, with $L_{x}$ — probably reflecting a lower probability for an object to have a higher instantaneous luminosity."
+" At the lowest redshifts in our highersample, probes such low luminosity that ray AGNs become quite common."," At the lowest redshifts in our sample, probes such low luminosity that X-ray AGNs become quite common."
+" At z=0.37, the dra--detected sources are located in more than of the dark matter halos with vg>280km or M>3.7x1012/ΓΙ Mo."," At $z=0.37$, the -detected sources are located in more than of the dark matter halos with $\vmax>280\,$ or $M>3.7\times10^{12}\,h^{-1}\,M_{\odot}$ ."
+ This work was supported by the Smithsonian Institution and contracts NAS8-39073 and NAS8-03060 (CXC)., This work was supported by the Smithsonian Institution and through contracts NAS8-39073 and NAS8-03060 (CXC).
+" AK is throughsupported by the NSF grant AST-0708154, by NASA grant NAGS-13274, and by Kavli Institute for Cosmological Physics at the University of Chicago through grant NSF PHY- and an endowment from the Kavli Foundation."," AK is supported by the NSF grant AST-0708154, by NASA grant NAG5-13274, and by Kavli Institute for Cosmological Physics at the University of Chicago through grant NSF PHY-0551142 and an endowment from the Kavli Foundation."
+ We thank O. Gnedin and A. Franceschini for useful discussions and comments., We thank O. Gnedin and A. Franceschini for useful discussions and comments.
+ MMT., MMT.
+ MMT. , MMT.
+ MMT. B, MMT.
+ MMT. Bl, MMT.
+ MMT. Bll, MMT.
+ MMT. Bllu, MMT.
+ MMT. Bllup, MMT.
+When the ecquilibrium state is reached. OJ/Or=0. since (his equation is linear. we can write ils solution as J=J.+J;. where J. is the spectrum of the continuum photons. anc J; is the spectrin of the recombination photons injected near the line center.,"When the equilibrium state is reached, $\partial J/\partial \tau=0. $ Since this equation is linear, we can write its solution as $J= J_c + J_i$ , where $J_c$ is the spectrum of the continuum photons, and $J_i$ is the spectrum of the recombination photons injected near the line center."
+ For the continuum photons. there is no injection and so C'(/)=0.," For the continuum photons, there is no injection and so $C(t)=0$."
+ Equation (13)) is reduced to where lis a constant of integration determined by the boundary. condition: as 7—x. or)—0. J.—Ja. where Ja;=(dx/c)na is the unperturbed. intensity.," Equation \ref{eq:FP}) ) is reduced to where $A$ is a constant of integration determined by the boundary condition: as $x \rightarrow \infty$, $\phi(x)\rightarrow 0$, $J_c \rightarrow J_{c0} $, where $J_{c0}= (4\pi/c)\, n_{c0}$ is the unperturbed intensity."
+ Therefore. 4-255454.," Therefore, $A=2\gamma_S J_0$."
+ The equation is then solved as an initial value problem of an ordinary differential equation with the condition J.(re)4Jig forc ——2x€. required by photon number conservation.," The equation is then solved as an initial value problem of an ordinary differential equation with the condition $J_c(x) \rightarrow J_{c0}$ for $x \rightarrow - \infty$, required by photon number conservation."
+ In Fig. L..," In Fig. \ref{Jcon},"
+ we plot the solution for a (vpical case: T=2.65 and 55=1.4x10.*. corresponding to the physical conditions of unheated gas al z=10 in a flat ACDAM model.," we plot the solution for a typical case: $T=2.65$ K and $\gamma_S =1.44 \times 10^{-6}$, corresponding to the physical conditions of unheated gas at $z=10$ in a flat $\Lambda$ CDM model."
+ An absorption feature is produced near the center of the line., An absorption feature is produced near the center of the line.
+ For the injected photons. we assume a constant source CU)=C.," For the injected photons, we assume a constant source $C(t) = C$."
+ This should generally be an accurate approximation because (he timescale for (he evolution of (he production rate ol photons from sources is likely be of order the Hubble time. which is much longer (han the relaxation time. which is of order the time to redshift across the width of the spectral profile of the cosmic background.," This should generally be an accurate approximation because the timescale for the evolution of the production rate of photons from sources is likely be of order the Hubble time, which is much longer than the relaxation time, which is of order the time to redshift across the width of the spectral profile of the cosmic background."
+ We also inject photons with a single Irequency al (he line center. Le. @Cr)=96e).," We also inject photons with a single frequency at the line center, i.e., $\psi(x)=\delta(x)$."
+ Because photons change randomly in [requency over a thermal width at each scattering. and the number of scatterings each photon undergoes is very large. injecting the photons with a thermal width instead does not make any difference to the final profile [we have verified this by obtaining solutions with cre) given by a Voigt profile. which did not result in any signilicant change to the resulting spectral profile].," Because photons change randomly in frequency over a thermal width at each scattering, and the number of scatterings each photon undergoes is very large, injecting the photons with a thermal width instead does not make any difference to the final profile [we have verified this by obtaining solutions with $\psi(x)$ given by a Voigt profile, which did not result in any significant change to the resulting spectral profile]."
+ The solution is where O(.r) is the Heaviside step Function., The solution is where $\Theta(x)$ is the Heaviside step function.
+ The value of the integration constant D can again be determined [rom boundary conditions: as à—2€. we must now have Z;(r)—0. from which we obtain B=2C.," The value of the integration constant $B$ can again be determined from boundary conditions: as $x \to +\infty$, we must now have $J_i(x) \to 0$, from which we obtain $B = 2 C$."
+ The profile of (he line can again be obtained by solving ((15)). with the initial condition J;(r=—2€)Jj. where Jig is the photon intensity due to the injected photons.," The profile of the line can again be obtained by solving \ref{eq:injdelta}) ), with the initial condition $J_i(x= -\infty) = J_{i0}$, where $J_{i0}$ is the photon intensity due to the injected photons."
+ The result is shown in Fie. 2..," The result is shown in Fig. \ref{Jinj},"
+ for (wo dilferent temperatures: we see (hat some of the photons injected at the line center are scattered to the right (blue)side., for two different temperatures; we see that some of the photons injected at the line center are scattered to the right (blue)side.
+ We nole that (he solution for the spectral profile of continuum and injected photons has a characteristic width in the variable 7=v/Avyp that is substantially larger than unity.," We note that the solution for the spectral profile of continuum and injected photons has a characteristic width in the variable $x = \nu /
+\Delta\nu_D$ that is substantially larger than unity,"
+in NCC L10£ just iuside the edge.,in NGC 1404 just inside the edge.
+ The results of the spectral aualyses are sumauiarized in Table 1.., The results of the spectral analyses are summarized in Table \ref{tab:spectra}.
+ The best fits for the cluster region (W1) aud Dor the galaxy emission inside the edge (E) are shown in Figure L., The best fits for the cluster free-stream region (W1) and for the galaxy emission inside the edge (E) are shown in Figure \ref{fig:edgespec}.
+ We fiud a temperature (1.53um Κον) and abuucdauce (0.12umZ. ) lor gas in the [ree stream region.," We find a temperature $1.53^{+0.10}_{-0.13}$ keV) and abundance $0.42^{+0.2}_{-0.13}\,\Zs$ ) for gas in the free stream region."
+" These values are in excellent agreement with XMM-Newton measurements by Buote (2002) for diffuse gas δ from NGC 1399 (+2"" from NGC 1101). who finds the data at large radius from NGC 1399 well represented by a single temperature APEC model with temperature FP—1.35£0.15 kkeV (errors lo) aud abundauce 4A~0.1—0.6Z.."," These values are in excellent agreement with XMM-Newton measurements by Buote (2002) for diffuse gas $\sim 8'$ from NGC 1399 $\sim 2'$ from NGC 1404), who finds the data at large radius from NGC 1399 well represented by a single temperature APEC model with temperature $kT = 1.35 \pm 0.15$ keV (errors $1\sigma$ ) and abundance $A \sim 0.4 - 0.6\,\Zs$."
+ They are also in reasouable agreement with Jones (1997). who find kT=1.26(i kkeV and A=0.51TnZ. for radii r8’—LO’ [rom NGC 1399 and by O'Sullivan (2003). who find at the same radius 47—1.21.15 kkeV. aud A0.1—0.7Z. (see their Figure 1) using ROSAT PSPC data. and with Scharl (2001) who [iuc AD—1.6 using Chandra ACIS-I data.," They are also in reasonable agreement with Jones (1997), who find $kT = 1.26^{+0.08}_{-0.11}$ keV and $A=0.51^{+0.23}_{-0.16}\,\Zs$ for radii $r \sim 8'-10'$ from NGC 1399 and by O'Sullivan (2003), who find at the same radius $kT \sim 1.2-1.45$ keV and $A \sim 0.4 - 0.7\,\Zs$ (see their Figure 1) using ROSAT PSPC data, and with Scharf (2004) who find $kT \sim 1.6$ using Chandra ACIS-I data."
+ Due to limited statistics we are unable to [ουν vary all the parameters of the fit lor region WO. the dilfuse gas just outside the edge containing the staguation region.," Due to limited statistics we are unable to freely vary all the parameters of the fit for region W0, the diffuse gas just outside the edge containing the stagnation region."
+ We fix the abuudance at the best fit value for region WI and allow the temperature to vary., We fix the abundance at the best fit value for region W1 and allow the temperature to vary.
+ We find a temperature 41.660235 jj5kkeV; consistent with that for region WI uoagiven the large errors., We find a temperature $1.66^{+0.25}_{-0.15}$ keV consistent with that for region W1 given the large errors.
+ The somewhat cooler temperature ( LkkeV) measured by O'Sullivan. (2003) with ROSAT data when using a, The somewhat cooler temperature $\sim 1$ keV) measured by O'Sullivan (2003) with ROSAT data when using a
+‘Table 3 shows the results with this latter scatter (σος) taken as 0. 0.10. or 0.20 mag.,"Table 3 shows the results with this latter scatter $\sigma _{PLC}$ ) taken as 0, 0.10, or 0.20 mag."
+ The solutions correspond to the four solutions of Table 2., The solutions correspond to the four solutions of Table 2.
+ Lt is clear that within reasonable limits the value of σρες has little ellect., It is clear that within reasonable limits the value of $\sigma _{PLC}$ has little effect.
+ We adopt solution 4 with epic=0.10 2.38+0.10).," We adopt solution 4 with $\sigma _{PLC} = 0.10$ $-2.38
+\pm 0.10$ )."
+ Phis is 0.11 mag brighter than than the zero-point used by PN. or an increase in the distance scale of 5 percent.," This is 0.11 mag brighter than than the zero-point used by PMB, or an increase in the distance scale of 5 percent."
+ The Llipparcos proper motions and their errors (see section 2 above) were converted into components in galactie longitude and latitude (a.=qncosb. anc ye) and their errors. using the constants provided in the Llipparcos Catalogue and taking into account the correlation between proper motions in Right Ascension and. Declination.," The Hipparcos proper motions and their errors (see section 2 above) were converted into components in galactic longitude and latitude $ \mu_{\ell *}
+= \mu_{\ell} \cos b$, and $ \mu_{b}$ ) and their errors, using the constants provided in the Hipparcos Catalogue and taking into account the correlation between proper motions in Right Ascension and Declination."
+ Then if yr is the proper motion in mas ancl d is the distance in kpc. the corresponding velocity is αμαkms+ where s=4.74047 (a value also taken from the Llipparcos Catalogue)," Then if $\mu$ is the proper motion in mas and $d$ is the distance in kpc, the corresponding velocity is $\kappa \mu d\ {\rm km\ s^{-1}}$ where $\kappa = 4.74047$ (a value also taken from the Hipparcos Catalogue)."
+ In analysing the proper motions account must be taken in the weighting not only of the standard errors of the proper motion components themselves (σι. σε) but also the scatter about an adopted galactic model due to the random motions of the stars.," In analysing the proper motions account must be taken in the weighting not only of the standard errors of the proper motion components themselves $\sigma_{\ell *}, \sigma_{b}$ ) but also the scatter about an adopted galactic model due to the random motions of the stars."
+ Phe weights were taken as proportional to the reciprocals of the squares of σεν and oj). where. and a. and 5 are the axes of the velocity ellipsoid for which we have adopted the values 13. 9 and 5 kms.7 for Copheids from Delhave (1965)(see also PAIB).," The weights were taken as proportional to the reciprocals of the squares of $\sigma_{\kappa
+\mu \ell *}$ and $\sigma_{\kappa \mu b}$, where, and $\rm \alpha, \beta$ and $\rm \gamma$ are the axes of the velocity ellipsoid for which we have adopted the values 13, 9 and 5 $\rm km\ s^{-1}$ for Cepheids from Delhaye (1965)(see also PMB)."
+ Of the Cepheids for which astrometric cata were available from. Hipparcos. the following stars have been omitted in the analyses for the reasons given below.," Of the Cepheids for which astrometric data were available from Hipparcos, the following stars have been omitted in the analyses for the reasons given below."
+ In most cases the stars were rejected. because of the large peculiar velocities (given. below) derived. in. preliminary analyses using PL distances and the proper motions., In most cases the stars were rejected because of the large peculiar velocities (given below) derived in preliminary analyses using PL distances and the proper motions.
+ L110. 1129 and. 130 indicate an anomalous entry in the relevant field. Ln. of the Lipparcos catalogue (for more details see the explanatory supplement).," H10, H29 and H30 indicate an anomalous entry in the relevant field, Hn, of the Hipparcos catalogue (for more details see the explanatory supplement)."
+ In the following LILO) indicates that the star is Uageed in LILO as being in a binary or multiple system: 129 indicates that 20 percent or more of the data were rejected in the astrometric solution and 1130 indicates that the goodness of [it statistic is greater than 3. possibly implying à poor astrometric DP Vel: HELP46610: Photometry too sparse: see paper AN. Per: HIP22275: Binary: see paper AN Cir: HIP 72773: Binary: see paper 1: 10. UN Per: LIP10332: Peculiar velocity ~500kms|: 1110. 1129. SS CAla: LILP36088: Peculiar velocity ~200kms.|. V Vel: LIP45949: Rejected by ΙΔΗΣ due to discrepant radial velocity.," In the following H10 indicates that the star is flagged in H10 as being in a binary or multiple system; H29 indicates that 20 percent or more of the data were rejected in the astrometric solution and H30 indicates that the goodness of fit statistic is greater than 3, possibly implying a poor astrometric DP Vel; $\rm HIP\, 46610$; Photometry too sparse; see paper AW Per; $\rm HIP\, 22275$; Binary; see paper AX Cir; $\rm HIP\, 72773$ ; Binary; see paper 1; H10, UX Per; $\rm HIP\, 10332$; Peculiar velocity $\rm \sim 500\ km\ s^{-1}$; H10, H29, SS CMa; $\rm HIP\, 36088$; Peculiar velocity $\rm \sim 200\ km\ s^{-1}$ V Vel: $\rm HIP\, 45949$; Rejected by PMB due to discrepant radial velocity."
+ Itejected by Calcwell Coulson (1987) due to discrepant colours., Rejected by Caldwell Coulson (1987) due to discrepant colours.
+ Peculiar velocity ~SOkms.+. SU Cru: LP59996: Rejected by PMD due to cliserepant racial velocity.," Peculiar velocity $\rm \sim 80\ km\ s^{-1}$ SU Cru; $\rm HIP\, 59996$; Rejected by PMB due to discrepant radial velocity."
+ Peculiar velocity ~240kms.+: LILO. SY Nor: HIP77913: Peculiar velocity ~100kms 1110. TW Nor: HIPT8771: Peculiar velocity ~100kms+.," Peculiar velocity $\rm \sim 240\ km\ s^{-1}$; H10, SY Nor; $\rm HIP\, 77913$; Peculiar velocity $\rm \sim 100\ km\ s^{-1}$; H10, TW Nor; $\rm HIP\, 78771$; Peculiar velocity $\rm \sim 100\ km\ s^{-1}$."
+ ‘This is an important star in. or in the direction of. the open cluster Lvnga 6.," This is an important star in, or in the direction of, the open cluster Lynga 6."
+ Caldwell Coulson (LOST) rejected it as a PL calibrator because il gave a cliscrepant zero-point but more recently Lancy Stobie (1994) have included it in a zero-point determination bv the cluster moethod., Caldwell Coulson (1987) rejected it as a PL calibrator because it gave a discrepant zero-point but more recently Laney Stobie (1994) have included it in a zero-point determination by the cluster method.
+" ]t is useful to begin with a simple. first order (""Oort type) solution for galactic rotation from the proper motions."," It is useful to begin with a simple, first order (“Oort"" type) solution for galactic rotation from the proper motions."
+" The relevant equation is generally written in the form: where i5, and c, are the components of the local solar motion (vith respect to voung objects) towards the galactic centre and in the direction of galactic rotation.", The relevant equation is generally written in the form; where $u_{\rm o}$ and $v_{\rm o}$ are the components of the local solar motion (with respect to young objects) towards the galactic centre and in the direction of galactic rotation.
+ Phe Oort constants are defined by and where Ris the distance of the star from the galactic centre anc © the angular velocity (=O/R. O being the circular velocity).," The Oort constants are defined by and where $R$ is the distance of the star from the galactic centre and $\rm
+\Omega$ the angular velocity $\rm \Theta /R$, $\rm \Theta$ being the circular velocity)."
+ Subscript 70 indicates the values of these quantities at the solar position., Subscript “o” indicates the values of these quantities at the solar position.
+ Llowever. the quantity of interest is often not £3 itself but Lt is useful therefore to solve equation 6 in the form: Thus here and in the rest of the paper we work in ternis of A and. €. but explicit values of D. are given. for. our final solutions.," However, the quantity of interest is often not $B$ itself but It is useful therefore to solve equation 6 in the form: Thus here and in the rest of the paper we work in terms of $A$ and $\Omega_{\rm o}$, but explicit values of $B$ are given for our final solutions."
+" Decause of the uniformity of the Llipparcos reference frame and the fact that it is tied to the positions of extragalactic objects. the values of D and©, Owhich are clirecthy allecteck by any rotation of the reference. frame) should be considerably. more reliable than previous values."," Because of the uniformity of the Hipparcos reference frame and the fact that it is tied to the positions of extragalactic objects, the values of $B$ and$ \Omega_{\rm o}$ (which are directly affected by any rotation of the reference frame) should be considerably more reliable than previous values."
+ΗΕ ΕΕ the sole source discovered with üu the core of NGC 6201 (C09). and was classified as a GLAINB based on dts. X-ray. spectrum.,"This object was the sole source discovered with in the core of NGC 6304 (G09), and was classified as a qLMXB based on its X-ray spectrum."
+ It was composed of two componcuts: an ILatmosplere neutron star spectrum at the distance of NGC 6301. aud a hard power-law component which dominates the spectra at high cnereies.," It was composed of two components: an H-atmosphere neutron star spectrum at the distance of NGC 6304, and a hard power-law component which dominates the spectrum at high energies."
+ A areinal positional offset between the thermal source (candidate GLAINB) detected on a soft dmage (1.5 keV) and a harder source (power-law conrponeut) detected ou a hard Πμασο (>1.5keV) of the ddata teutatively sueecsted that the two spectral colponcuts are due to distinct point sources (C09)., A marginal positional offset between the thermal source (candidate qLMXB) detected on a soft image $<1.5\keV$ ) and a harder source (power-law component) detected on a hard image $>1.5\keV$ ) of the data tentatively suggested that the two spectral components are due to distinct point sources (G09).
+ Tt is shown in the following subsections that the ssource ls resolved. spatially and spectrally. iuto two sources detected withC," It is shown in the following subsections that the source is resolved, spatially and spectrally, into two sources detected with."
+handra. Iu Figure L. two sources; aand 15.5.. are detected2027 (0> 3) within the core radius of NGC 6301. separated bv an angular distance of ~137.," In Figure \ref{fig:core}, two sources, and , are detected $\sigma>3$ ) within the core radius of NGC 6304, separated by an angular distance of $\sim13\arcsec$."
+ The brightest one IS.0.. accouuting for 71 counts (iucludiug about two backeround cots) is positionally cousistent with the candidate qLMXDL7... within 26.," The brightest one –, accounting for 74 counts (including about two background counts) -- is positionally consistent with the candidate qLMXB, within $2\sigma$."
+" The relative offset between the aand ppositious of the source isppiull.G"": the error on the aastrometry (statistical aud systematic errors) accounts for (lo).", The relative offset between the and positions of the source is; the error on the astrometry (statistical and systematic errors) accounts for $1\sigma$ ).
+ Since ccannot resolve the two sources in the core. the ceutroid position obtained with is in between the two sources detected with Chundru//ACTS (Figure 11.," Since cannot resolve the two sources in the core, the centroid position obtained with is in between the two sources detected with /ACIS (Figure \ref{fig:core}) )."
+ As ineutioned above. the extracted events are left nubinued and the spectiin is fitted with the assmued model. inchiding a systematic uncertaiutv. and the confidence regions are derived using the Cash-statistic (?)..," As mentioned above, the extracted events are left unbinned and the spectrum is fitted with the assumed model, including a systematic uncertainty, and the confidence regions are derived using the Cash-statistic \citep{cash79}."
+ The ealactic absorption (hydrogen column deusity. ιο) is kept fixed at the value in the direction of NGC 6301: Nyy0.266«1072atomsem2 (Nypoo=0.266 afterward).," The galactic absorption (hydrogen column density, ) is kept fixed at the value in the direction of NGC 6304: $\nh=0.266\tee{22} \unit{atoms\percmsq}$ $\nhtt=0.266$ afterward)."
+ The derived best-fit parameters and confidence regions (Table 3)) are compared to the previously reported values (C09)., The derived best-fit parameters and confidence regions (Table \ref{tab:res}) ) are compared to the previously reported values (G09).
+ The source lis spectrally soft with all events haviug energies below25keV., The source is spectrally soft with all events having energies below.
+. It is assuned to correspond to the qLMXDL7.. aud therefore a NS lydrogcu atmosphere model is used for the spectral fitting.," It is assumed to correspond to the qLMXB, and therefore a NS hydrogen atmosphere model is used for the spectral fitting."
+ For this ssource the only one in this observation — the nuuber of counts allows for spectral binning (~15 counts per bius). while having an acceptably approximate Cassia uncertainty in each biu.," For this source – the only one in this observation – the number of counts allows for spectral binning $\sim15$ counts per bins), while having an acceptably approximate Gaussian uncertainty in each bin."
+ The chosen backerouud is a circular reeion of aaround the core. excluding (ECE= 99%) around cach detected source.," The chosen background is a circular region of around the core, excluding $ECF=99\%$ ) around each detected source."
+ A 5-bin spectra is obtained in the irauge., A 5-bin spectra is obtained in the range.
+ The model used is the tabulated model of NS wdrogen atmosphere (2).. together with fixed Nyy 1.266.," The model used is the tabulated model of NS hydrogen atmosphere \citep{zavlin96}, together with fixed $\nhtt=0.266$ ."
+ The best-fit parameters of this \2--fit ave: Tig=128!bieV aud Ry=7.3!kan., The best-fit parameters of this -fit are: $\kteff=128\ud{40}{30}\eV$ and $\rinfty=7.3\ud{10.5}{3.7}\km$.
+ Although the counting statistic is poor. the best-fit nvalues are in agreement witli he previous results. and the model fitted is statistically acceptable (0.95)).," Although the counting statistic is poor, the best-fit values are in agreement with the previous results, and the model fitted is statistically acceptable )."
+" For colpletcuess. the spectrum is fitted with an absorbed ΝΟΥ law. leading to a best-fit photon iudex a=3,507 with(0.27). which inclicates a soft source. as expected for a typical qLAINB."," For completeness, the spectrum is fitted with an absorbed power law, leading to a best-fit photon index $\alpha=3.5\ud{0.5}{0.4}$ with, which indicates a soft source, as expected for a typical qLMXB."
+ This provides further support to the aremuent that lis the qLAINB detect in17., This provides further support to the argument that is the qLMXB detect in.
+. In Table 3.. we report best-fit values determined using nof the 4? statistic with binned data. but the Cashestatistic using uubimned data. which typically produces smaller error regious.," In Table \ref{tab:res}, we report best-fit values determined using not the $\chi^2$ statistic with binned data, but the Cash-statistic using unbinned data, which typically produces smaller error regions."
+ The results of Monte Carlo simulation are also eiven in this table., The results of Monte Carlo ``goodness-of-fit'' simulation are also given in this table.
+ The best-fit ILatinosphere paralctersfor aare consistent with typical values for quiesceut NSs, The best-fit H-atmosphere parametersfor are consistent with typical values for quiescent NSs
+We have shown the power of a new technique. that of millimetre/submillimetre polarimetry. in. deducing new constraints on models of compact. f[at-spectrum. raclio-jet sources. which have been well-stucicd using other techniques for many vears.,"We have shown the power of a new technique, that of millimetre/submillimetre polarimetry, in deducing new constraints on models of compact, flat-spectrum, radio-jet sources, which have been well-studied using other techniques for many years."
+ One of the biggest limitations on this experiment has been the uncertainty din the jet. direction on the very small (milliparsec) scales sampled: at. these wavelengths., One of the biggest limitations on this experiment has been the uncertainty in the jet direction on the very small (milliparsec) scales sampled at these wavelengths.
+ Polarisation-sensitive VLBI measurements at l mm would be the ultimate tool for this experiemnt. however that is likely to be some vears away.," Polarisation-sensitive VLBI measurements at 1 mm would be the ultimate tool for this experiemnt, however that is likely to be some years away."
+ A combination of the highest. possible frequency VLBI ancl millimetresubmillimetre photopolarimetric monitoring with further theoretical development is likely to be the way forward in the near future., A combination of the highest possible frequency VLBI and millimetre--submillimetre photopolarimetric monitoring with further theoretical development is likely to be the way forward in the near future.
+ We thankJason Stevens of the Joint Astronomy Centre. Llawaii. for useful discussions and for carrving out the 1995 service observing runs.," We thank Jason Stevens of the Joint Astronomy Centre, Hawaii, for useful discussions and for carrying out the 1995 service observing runs."
+ We also thank Alan Marscher for providing unpublished VLBI information on a number of sources and the anonymous referee. for his. comments and suggestions that have helped. to improve the paper., We also thank Alan Marscher for providing unpublished VLBI information on a number of sources and the anonymous referee for his comments and suggestions that have helped to improve the paper.
+ The ΙΟΝΤ is operated by the Joint Astronomy Centre on behalf of the Ulx Particle Physics ancl Astronomy Council (PPARC). the National Research Council. of Canada. the Netherlands Organisation. for the Acvwancement of Pure Research (ZWO) and the University. of Hawaii.," The JCMT is operated by the Joint Astronomy Centre on behalf of the UK Particle Physics and Astronomy Council (PPARC), the National Research Council of Canada, the Netherlands Organisation for the Advancement of Pure Research (ZWO) and the University of Hawaii."
+ IN acknowledges the receipt of an FPL scholarship from the Ministerio de Educacionn v Ciencia of Spain., RN acknowledges the receipt of an FPI scholarship from the Ministerio de Educaciónn y Ciencia of Spain.
+yhotometry This↴⋅∙∙ highlights the importance. of these types of coustraints for validating ccandidates. eiven that blends involving plisically associated stars would generally be spatially uuresolved wou high-resolution iuagiue with adaptive optics or speckle interferometry. and they would typically also be below the seusitivitv liuits of our cenutroid notion analvsis. so that they would not be detected by hose means.,"photometry This highlights the importance of these types of constraints for validating candidates, given that blends involving physically associated stars would generally be spatially unresolved by our high-resolution imaging with adaptive optics or speckle interferometry, and they would typically also be below the sensitivity limits of our centroid motion analysis, so that they would not be detected by those means."
+ Therefore. the oulv. bleuds we need to be concerned about for aare those consisting of stars in the backeround of the arect that are orbited by other stars or by transiting λαο».," Therefore, the only blends we need to be concerned about for are those consisting of stars in the background of the target that are orbited by other stars or by transiting planets."
+ Iu order to estimate the frequeney of the bleud scenarios (1.ce.. background configurations) that remain xossible after applying aand all other observational constraints. we follow a mocedure simular to that described by Torresetal.(2001). for Kepler-Odd. We appeal to the Desaucoon Galactic structure models of Robinetal.(2003) to xediet the wmmber density of background stars of cach spectral type (1iass) and brightuess around I&epler-10. iu wilfinacuitucde bins. aud we make use of estimates ofthe requencies of transiting plaucts and of eclipsing binaries roni recent studies by the TTeam to infer the nuniber density of leds.," In order to estimate the frequency of the blend scenarios (i.e., background configurations) that remain possible after applying and all other observational constraints, we follow a procedure similar to that described by \cite{Torres:11} for d. We appeal to the Besançoon Galactic structure models of \cite{Robin:03} to predict the number density of background stars of each spectral type (mass) and brightness around Kepler-10, in half-magnitude bins, and we make use of estimates of the frequencies of transiting planets and of eclipsing binaries from recent studies by the Team to infer the number density of blends."
+ Using coustrainuts from our liel-resolution iuagine (specifically.thesensitivitycurvespresentedbyBatalhaetal.2011.theirFie.9) we calculate the area around the target within which blends would go undetected. aud with this the expected iuuber of The receut release by Boruckietal.(20115). of a List of 1235 candidate transiting planets (ΟΤΕ) from pprovides a means to estimate planet frequencies needed for our calculations. with significant advantages over the calculations of Torresetal.(2011) for dd. which were based on the earlier list of candidates published by Boruckietal.(2011a).," Using constraints from our high-resolution imaging \citep[specifically, the
+sensitivity curves presented by][their Fig.~9]{Batalha:11} we calculate the area around the target within which blends would go undetected, and with this the expected number of The recent release by \cite{Borucki:11b} of a list of 1235 candidate transiting planets (KOIs) from provides a means to estimate planet frequencies needed for our calculations, with significant advantages over the calculations of \cite{Torres:11} for d, which were based on the earlier list of candidates published by \cite{Borucki:11a}."
+. Not oulv is the sample now much larger. but the knowledee of the rate of false positives for Hs also mich improved. aud that rate is believed to be relatively siuall depending on the level of vetting of the candidate. according to Doruckietal. 2011: less than according to the recent study by Mortou&Joluson 2011)).," Not only is the sample now much larger, but the knowledge of the rate of false positives for is also much improved, and that rate is believed to be relatively small depending on the level of vetting of the candidate, according to \citealt{Borucki:11a}; ; less than according to the recent study by \citealt{Morton:11}) )."
+ Thus. our results will not be significautly affected by the assumption that all of the candidates are plaucts (see also below).," Thus, our results will not be significantly affected by the assumption that all of the candidates are planets (see also below)."
+ An additional assumption we male is that this census is larecly complete., An additional assumption we make is that this census is largely complete.
+" Amone these candidates we count a total of 267 having radi in the range allowed bv ffor the tertiaries of viable bleuds G.c.. between 0.12 aud l.51BRg4,)."," Among these candidates we count a total of 267 having radii in the range allowed by for the tertiaries of viable blends (i.e., between 0.42 and $R_{\rm Jup}$ )."
+ With the total umuber of ttareets beme 156.153 (Doruckictal.2011b). the relevant frequency of transiting plauets for our bleud calculation is =267/156.1530.0017.," With the total number of targets being 156,453 \citep{Borucki:11b}, , the relevant frequency of transiting planets for our blend calculation is $f_{\rm planet} = 267/156,\!453 = 0.0017$."
+ Slawsonetal(2011) fühave recently published a catalog of the 2165RV] eclipsing binaries found im the field. from the first four mouths of observation.," \cite{Slawson:11} have recently published a catalog of the 2165 eclipsing binaries found in the field, from the first four months of observation."
+" Ouly 1e 1225 detached systeiis among these are considered rere, since binaries m the category of semü-detached. over-coutact. or ellipsoidal variables would uot produce ight curves with a shape consistent with a transit."," Only the 1225 detached systems among these are considered here, since binaries in the category of semi-detached, over-contact, or ellipsoidal variables would not produce light curves with a shape consistent with a transit."
+ The requency of eclipsing binarics for our purposes is then fep=1225/156.1530.0078.," The frequency of eclipsing binaries for our purposes is then $f_{\rm EB} = 1225/156,\!453 = 0.0078$."
+ Table 3. preseuts the results of our calculation of he frequency of bleuds. separately for backeround lends with stellar tertiavics (eclipsing binaries) aud with planetary tertiarics.," Table \ref{tab:new_stats} presents the results of our calculation of the frequency of blends, separately for background blends with stellar tertiaries (eclipsing binaries) and with planetary tertiaries."
+ Columus 1 and 2 eive the Ap maguitude range of each biu and the magnitude differcuce Ap relative to the target. calculated at the upper οσο of each bin.," Columns 1 and 2 give the $K\!p$ magnitude range of each bin and the magnitude difference $\Delta K\!p$ relative to the target, calculated at the upper edge of each bin."
+ Cobhuuu 3 reports the mean uunmber density of starsper square degree obtained frou the Besangoou models. for stars in the mass rauge allowed ly aas shown in Figure 2.," Column 3 reports the mean number density of starsper square degree obtained from the Besançoon models, for stars in the mass range allowed by as shown in Figure \ref{fig:backstar}."
+ In columu | we lst the niaxiumui augular separation Όρος at which stars in the corresponding maguitude bin would go undetected iu our imaging observations. taken from the information in the work of Batalhaetal.(2011).," In column 4 we list the maximum angular separation $\rho_{\rm max}$ at which stars in the corresponding magnitude bin would go undetected in our imaging observations, taken from the information in the work of \cite{Batalha:11}."
+. The product of he area iaplied by this radius aud the stellar deusities in the previous cobhuun sive the umuber of stars in he appropriate nass range. listed iu column 6 in mits of 109.," The product of the area implied by this radius and the stellar densities in the previous column give the number of stars in the appropriate mass range, listed in column 6 in units of $10^{-6}$."
+" Multiplying these fieures bv the frequency of eclipsing binarics frp, then eives the umuber of vackeround star| blends iu columu 7.", Multiplying these figures by the frequency of eclipsing binaries $f_{\rm EB}$ then gives the number of background star+star blends in column 7.
+ À similar calculation for the backeround star|planet blends. nakiug use of filaners IS presented in cohuuus 710.," A similar calculation for the background star+planet blends, making use of $f_{\rm planet}$, is presented in columns 7–10."
+ We suni up the contributions from cach maguitude bin at he bottom of cohuuus 6 aud 10., We sum up the contributions from each magnitude bin at the bottom of columns 6 and 10.
+ The total wuuber of lends we expect (blend frequency) is given iu he last liue of the table by adding these two values oecther. aud is BF=1.62«10.*.," The total number of blends we expect (blend frequency) is given in the last line of the table by adding these two values together, and is ${\rm BF} = 1.62
+\times 10^{-8}$."
+ The calculations slow hat backerouud bleuds consisting of star|planet pairs contribute to this frequency about three times more thu ckeround eclipsing binarics., The calculations show that background blends consisting of star+planet pairs contribute to this frequency about three times more than background eclipsing binaries.
+ While we have assumed up to now that any companions to wwithin ANp=2 mae of the tarect would have been secu spectroscopicallv. we note that relaxing this condition to aduuch more conservative AJp—1 has uo effect at all on the contribution from eclipsing binaries. aud a uceligible effect on the contribution of star|planet scenarios.," While we have assumed up to now that any companions to within $\Delta K\!p = 2$ mag of the target would have been seen spectroscopically, we note that relaxing this condition to a much more conservative $\Delta K\!p = 1$ has no effect at all on the contribution from eclipsing binaries, and a negligible effect on the contribution of star+planet scenarios."
+" To obtain a Bayesian estimate of the probability that lis indeed a planet as opposed to a false positive (or equivalently. the ""false alarm rate”. FAR) we follow the eeneral mcthodoloey of Torresetal.(2011) ancl compare the priori likelihoods of blends aud of planets: FAR = DF/PF."," To obtain a Bayesian estimate of the probability that is indeed a planet as opposed to a false positive (or equivalently, the “false alarm rate”, FAR) we follow the general methodology of \cite{Torres:11} and compare the likelihoods of blends and of planets: FAR = BF/PF."
+ If the bleudfrequency issuticicutly, If the blendfrequency issufficiently
+We begin with the proof of Theorem 1.1. that strongly relies on the results obtained in [5]..,We begin with the proof of Theorem \ref{theo1} that strongly relies on the results obtained in \cite{Ibrahim09}.
+ Let ΝεN be any arbitrary integer., Let $N\in \N$ be any arbitrary integer.
+" Using (2.2)). we estimate ||f||;« in the following way: where Step 2 of the proof of [5.Theorem1.1]. asserts that: while ο,Lemuina3.1]. gives: Iu order to estimate Ay. we proceed iu the following way: hence (see Delinition 2.2)) Using the well known result (see for instance [1])) we finally obtain Inequalities (3.1)). (3.3)). (3.3)) aud (3.1)) imply: Optimizing the above inequality with respect to the variable Ν΄ (see Step 2 of the proof of 0. 3.2])). we directly arrive into the result."," Using \ref{dyad1}) ), we estimate $\|f\|_{L^{\infty}}$ in the following way: where Step 2 of the proof of \cite[Theorem~1.1]{Ibrahim09} asserts that: while \cite[Lemma 3.1]{Ibrahim09} gives: In order to estimate $A_{3}$, we proceed in the following way: hence (see Definition \ref{def2}) ) Using the well known result (see for instance \cite{Farkas}) ) we finally obtain Inequalities \ref{fer0}) ), \ref{fer1}) ), \ref{fer2}) ) and \ref{fer3}) ) imply: Optimizing the above inequality with respect to the variable $N$ (see Step 2 of the proof of \cite[Lemma~3.2]{Ibrahim09}) ), we directly arrive into the result."
+ Oo , $\hfill{\Box}$ 
+We searched for pulsations in Vega as a near pole-on configuration is particularly favorable for the seismology of rapidly rotating stars.,We searched for pulsations in Vega as a near pole-on configuration is particularly favorable for the seismology of rapidly rotating stars.
+" Equivalent photospheric line profiles were calculated, using the multiplex information of all photospheric spectral lines in the observed wavelength range."," Equivalent photospheric line profiles were calculated, using the multiplex information of all photospheric spectral lines in the observed wavelength range."
+" In this instrumental configuration, the highest attainable radial velocity calibration was obtained by taking the telluric water vapor line profiles present in the spectrum as absolute wavelength references."," In this instrumental configuration, the highest attainable radial velocity calibration was obtained by taking the telluric water vapor line profiles present in the spectrum as absolute wavelength references."
+" A subsequent radial velocity analysis with Period04 was performed and low amplitude (A « 8ms~!)) multi-periodic variations in the residual radial velocity curves were detected, with frequencies ranging roughly from 5 -cd,, depending on the run."," A subsequent radial velocity analysis with Period04 was performed and low amplitude (A $<$ ) multi-periodic variations in the residual radial velocity curves were detected, with frequencies ranging roughly from 5 -, depending on the run."
+" The frequencies detected in this work, if confirmed, are consistent with values occasionally reported in the past, as described in the review by Fernie (1981).."," The frequencies detected in this work, if confirmed, are consistent with values occasionally reported in the past, as described in the review by \citet{fernie}."
+" The observed amplitudes of the variations are very small with respect to the intrinsic calibration precision, but taking the telluric water vapor lines of the spectrum as a reference enabled us to reach a precision on the individual radial velocity measurement of mms~!."," The observed amplitudes of the variations are very small with respect to the intrinsic calibration precision, but taking the telluric water vapor lines of the spectrum as a reference enabled us to reach a precision on the individual radial velocity measurement of $^{-1}$."
+" This auto-calibration of the spectrum should rule out, in principle, any instrumental"," This auto-calibration of the spectrum should rule out, in principle, any instrumental"
+280 respectively.,280 respectively.
+ The limits on fay can also be checked using the fisher discriminant function (Barreiro Hobson 2001. Cavónn et al.," The limits on $f_{NL}$ can also be checked using the fisher discriminant function (Barreiro Hobson 2001, Cayónn et al."
+ 2003)., 2003).
+" An optimal linear function of the measured. variables (here. skewness of wavelet coefficients on 15 different scales) is Ileve is the 15 element vector that contains the skewness values at the different scales considered here. and /(x) is the fisher discriminant fonction that optimally puts together inlormation contained in in (he sense of maximizing the difference between (he expected mean values of/ from Gaussian and non-Gaussian models. ancl minimizing (heir dispersions. W=V,4+Vy."," An optimal linear function of the measured variables (here, skewness of wavelet coefficients on 15 different scales) is Here is the 15 element vector that contains the skewness values at the different scales considered here, and $t({\bf{x}})$ is the fisher discriminant function that optimally puts together information contained in in the sense of maximizing the difference between the expected mean values of $t$ from Gaussian and non-Gaussian models, and minimizing their dispersions. ${\bf{W}}={\bf{V}}_0+{\bf{V}}_1$,"
+ the sum of the covariance matrices of the test statistic in the Gaussian (subscript 0) and non-Gaussian (subscript 1) cases., the sum of the covariance matrices of the test statistic in the Gaussian (subscript 0) and non-Gaussian (subscript 1) cases.
+ pg and poy are vectors containing the mean values of the test statistic in the Gaussian and non-Gaussian cases respectively., $_0$ and $_1$ are vectors containing the mean values of the test statistic in the Gaussian and non-Gaussian cases respectively.
+ Thus for different. fay. the fisher discriminant Function is found for each of the Gaussian realizations. for each of (he non-Gaussian realizations. and for the data.," Thus for different $f_{NL}$, the fisher discriminant function is found for each of the Gaussian realizations, for each of the non-Gaussian realizations, and for the data."
+ The probability (hat the data are drawn Irom one or the other hypothesis can then be estimated., The probability that the data are drawn from one or the other hypothesis can then be estimated.
+ Looking at the fraction of non-Gaussian simulations that have values larger (for positive fy) and smaller (for negative fy) we deduce similar limits on [νι as derived above using the 4? test., Looking at the fraction of non-Gaussian simulations that have values larger (for positive $f_{NL}$ ) and smaller (for negative $f_{NL}$ ) we deduce similar limits on $f_{NL}$ as derived above using the $\chi^2$ test.
+ Darreiro Hobson (2001) found that the fisher discriminant can do better than X7 s at distinguishing between Gaussian and non-Gaussian hypotheses., Barreiro Hobson (2001) found that the fisher discriminant can do better than $\chi^2$ 's at distinguishing between Gaussian and non-Gaussian hypotheses.
+ However since we have only 300 non-Gaussian simulations to obtain the fractions from. the accuracy of this method is not better but just comparable for our case here.," However since we have only 300 non-Gaussian simulations to obtain the fractions from, the accuracy of this method is not better but just comparable for our case here."
+" Figure 10 shows the histograms of the fisher discriminants of 1000 Gaussian realizations. of 300 non-Gaussian realizations and of the cata. for fy, values of 120 and 250: these values are close to the the lo and 2o limits derived using the 4? above."," Figure 10 shows the histograms of the fisher discriminants of 1000 Gaussian realizations, of 300 non-Gaussian realizations and of the data, for $f_{NL}$ values of 120 and 250; these values are close to the the $\sigma$ and $\sigma$ limits derived using the $\chi^2$ above."
+ For these values of fp. 0.74 and 0.95 of the non-Gaussian simulations. respectively. have larger values of the fisher discriminant (han the data.," For these values of $f_{NL}$, 0.74 and 0.95 of the non-Gaussian simulations, respectively, have larger values of the fisher discriminant than the data."
+ We see (hat results obtained using fisher discriminants are consistent with those obtained using the \7test’., We see that results obtained using fisher discriminants are consistent with those obtained using the $\chi^2$.
+ We note that if the skewness sienal in the data showed deviations [rom Caussianily on parücular scales then (his method involving wavelet (ransforms could be used to obtain scale dependent constraints on νι., We note that if the skewness signal in the data showed deviations from Gaussianity on particular scales then this method involving wavelet transforms could be used to obtain scale dependent constraints on $f_{NL}$.
+ However the skewness signal in (he data is well consistent with Gaussianitv., However the skewness signal in the data is well consistent with Gaussianity.
+ The skewness spectrum of Fie l(a) really flattens out with the use of another, The skewness spectrum of Fig 1(a) really flattens out with the use of another
+field).,field).
+" The value of temperature anisotropy is not constant over the whole computational domain. but depends on the spatial position. and 7,/T ranges from 0.98 to 4.9."," The value of temperature anisotropy is not constant over the whole computational domain, but depends on the spatial position, and $T_\parallel/T_\perp$ ranges from 0.98 to 4.9."
+ This suggest that the heating is connected with the formation of magnetic topological structures., This suggest that the heating is connected with the formation of magnetic topological structures.
+ An interesting feature that emerges bv examining the electron distribution function is the presence of non-thermal features., An interesting feature that emerges by examining the electron distribution function is the presence of non-thermal features.
+ However. the deviation from a bi-Maxwellian distribution is more accentuated when (he particles aud (the mean magnetic field are considered in a small portion of the box.," However, the deviation from a bi-Maxwellian distribution is more accentuated when the particles and the mean magnetic field are considered in a small portion of the box."
+ We show in Figure (23)) the electron distribution function in the Cr.jy) plane for four dillerent nested boxes.," We show in Figure \ref{pdf_xy}) ) the electron distribution function in the $(x,y)$ plane for four different nested boxes."
+ Velocilies are normalized to the speed of light., Velocities are normalized to the speed of light.
+" The solid line indicates the direction of the mean magnetic field within each box. and {, is the length of (he boxes. in ion gvroradius units."," The solid line indicates the direction of the mean magnetic field within each box, and $L$ is the length of the boxes, in ion gyroradius units."
+ Ht turns out that as the magnetic field ancl the particles are averaged over larger and larger boxes. the distribution function tends to become more and more bi-Maxwellian aud anv non-thermal feature is lost.," It turns out that as the magnetic field and the particles are averaged over larger and larger boxes, the distribution function tends to become more and more bi-Maxwellian and any non-thermal feature is lost."
+ A similar result holds in the (7.z) plane. while the distributions in (he direction orthogonal to (he mean magnetic field remains roughly gvrotropic.," A similar result holds in the $(x,z)$ plane, while the distributions in the direction orthogonal to the mean magnetic field remains roughly gyrotropic."
+ This result has important implications for satellite observations., This result has important implications for satellite observations.
+ Idee. il suggests that Che (ime window over which the particles and the magnetic field are averaged allects the distribution functions and therelore the anisotropy. measured.," Indeed, it suggests that the time window over which the particles and the magnetic field are averaged affects the distribution functions and therefore the anisotropy measured."
+ By assuming a steady. wind flowing al 250 km/s. one can transform the box lengths into lrequencies al which data are collected. aud estimate what is the level of anisotropy measured at clilferent frequencies.," By assuming a steady wind flowing at 250 km/s, one can transform the box lengths into frequencies at which data are collected, and estimate what is the level of anisotropy measured at different frequencies."
+ The upper bound of such anisotropies are shown. lor electrons. in Figure (24)).," The upper bound of such anisotropies are shown, for electrons, in Figure \ref{anisotropy}) )."
+ Given (he small length of the computational box used in this study. (he range of frequencies plotted in Figure (24)) is relatively high.," Given the small length of the computational box used in this study, the range of frequencies plotted in Figure \ref{anisotropy}) ) is relatively high."
+ However. future missions and instruments max be able to collect data at such frequencies. aid therefore will be able to verily (liis prediction.," However, future missions and instruments may be able to collect data at such frequencies, and therefore will be able to verify this prediction."
+ We have performed two-dimensional fullvy-kinetic PIC simulations of decaving magnetic fluctuations al electron scales., We have performed two-dimensional fully-kinetic PIC simulations of decaying magnetic fluctuations at electron scales.
+ The scope of this work is to address some open questions regarding the dissipation of turbulent. [Inctuations al small scales. and the plausibilitw of treating such [Iuctuations in the framework of linear We have shown the results of the first simulations that use realistic parameters in a lairly large computational domain that extends to electron scales.," The scope of this work is to address some open questions regarding the dissipation of turbulent fluctuations at small scales, and the plausibility of treating such fluctuations in the framework of linear We have shown the results of the first simulations that use realistic parameters in a fairly large computational domain that extends to electron scales."
+ This represents an improvenient in the context of the previous 2D The results of our simulations suggest that a turbulent cascade in the form ofa power-law could proceed very far below the electron gvroradius., This represents an improvement in the context of the previous 2D The results of our simulations suggest that a turbulent cascade in the form of a power-law could proceed very far below the electron gyroradius.
+ It is interesting that al such scales the, It is interesting that at such scales the
+For the sources identified with galaxies we convert the counts in the 0.5—2.0 keV band to the flux assuming a 1 keV thermal plasma of 0.3 solar metallicity. yielding the effective area of the telescope of 1054 em? and a mean energy of the photon of 0.94 keV. In estimating the net flux from the detected source. we also subtract the background due to the diffuse emission of the Coma cluster.,"For the sources identified with galaxies we convert the counts in the $0.5-2.0$ keV band to the flux assuming a 1 keV thermal plasma of 0.3 solar metallicity, yielding the effective area of the telescope of 1054 $^2$ and a mean energy of the photon of 0.94 keV. In estimating the net flux from the detected source, we also subtract the background due to the diffuse emission of the Coma cluster."
+ For this we use a larger (80” in radius) circle. centered on the source. excluding the central 20”. containing most of the source flux (see above for details) and we scale the flux according to the relative area-exposure—solid angle product (cm s aremin*).," For this we use a larger $80^{\prime\prime}$ in radius) circle, centered on the source, excluding the central $20^{\prime\prime}$, containing most of the source flux (see above for details) and we scale the flux according to the relative area--exposure--solid angle product $^2$ s $^2$ )."
+ This implicitly assumes a flat distribution of the diffuse cluster X-ray emission on the 80” scale. which is appropriate even for substructure in the Coma diffuse emission.," This implicitly assumes a flat distribution of the diffuse cluster X-ray emission on the $80^{\prime\prime}$ scale, which is appropriate even for substructure in the Coma diffuse emission."
+ An underestimate of the source flux due to subtraction of the scattered flux is less than in this procedure., An underestimate of the source flux due to subtraction of the scattered flux is less than in this procedure.
+" However. when halos of galaxies are filled with hotter medium. like in the case of NGC4889 (GMP 2921). the use of the 80"" aperture underestimates the “background’."," However, when halos of galaxies are filled with hotter medium, like in the case of NGC4889 (GMP 2921), the use of the $80^{\prime\prime}$ aperture underestimates the 'background'."
+ We also tested the background subtraction in 40 aperture., We also tested the background subtraction in $40^{\prime\prime}$ aperture.
+ Significant differences were found for galaxies GMP 2921. 3152. 3816.," Significant differences were found for galaxies GMP 2921, 3152, 3816."
+ For these galaxies a smaller aperture was adopted (40” in radius) to estimate the Coma diffuse background with, For these galaxies a smaller aperture was adopted $40^{\prime\prime}$ in radius) to estimate the Coma diffuse background with
+as distance estimators since they are less affected by obscuration.,as distance estimators since they are less affected by obscuration.
+ Further uncertainties include the mass of the NS and the gravitational redshift., Further uncertainties include the mass of the NS and the gravitational redshift.
+ Up to now. three photospheric radius expansion bursts have been reported from RNTE on Mav 2001 aud June and August 2005 (?7)..," Up to now, three photospheric radius expansion bursts have been reported from RXTE on May 2004 and June and August 2005 \citep{0748:wolff05apj, dist:galloway08apjs}."
+ The burst on Mav 2001 shows the largest peak flux., The burst on May 2004 shows the largest peak flux.
+ Based on spectral analysis of this burst. 7. estimated a distance of 7.7 Ispe (5.9 kpe) for a helmu(livdrogen)-dominated burst photosphere.," Based on spectral analysis of this burst, \citet{0748:wolff05apj} estimated a distance of 7.7 kpc (5.9 kpc) for a helium(hydrogen)-dominated burst photosphere."
+ ? used the three bursts to infer a distance of διXN) pe. where X is 0 for solar abundances ik 0.7 for hydrogen-rich material.," \citet{dist:galloway08apjs} used the three bursts to infer a distance of 8.3/(1+X) kpc, where X is 0 for solar abundances and 0.7 for hydrogen-rich material."
+ Further. they constrained the distance to T.l kpe areuing that the bursts shor rise times and low duration sugeest ieuitiou in a Ποσο cuviromuent.," Further, they constrained the distance to 7.4 kpc arguing that the bursts short rise times and low duration suggest ignition in a He-rich environment."
+ We note that the ? obtain a peak fiux for the brightest biunst of «110 5 cres 7 1.9! lower than the flux obtained bv ? due to the ower effective area of LUEASOFT v5.3. used in the latter uper (?)..," We note that the \citet{dist:galloway08apjs} obtain a peak flux for the brightest burst of $\times$ $^{-8}$ ergs $^{-2}$ $^{-1}$, lower than the flux obtained by \citet{0748:wolff05apj} due to the lower effective area of LHEASOFT v5.3, used in the latter paper \citep{dist:galloway08apjs}. ."
+ We estimate that Compton scatteriug in a highly ionised plasma (log ©— 1) for a hieli-incliuationOo source such as οσα reduce the flux by 1 to for a column deusity between 1072 and 10°? ?., We estimate that Compton scattering in a highly ionised plasma (log $\xi\sim$ 4) for a high-inclination source such as can reduce the flux by 1 to for a column density between $^{22}$ and $^{23}$ $^{-2}$.
+ If the asma las a lower degree of ionisation. the flux is reduced by a larger factor.," If the plasma has a lower degree of ionisation, the flux is reduced by a larger factor."
+ Therefore. if the flux of the burst on May. 2001 is nuderestimated by. this translates iuto a distance of 6.86.5 kpc for a helinm-domunatcc burst photosphere and of 5.25.0 kpc for a lydrogeu-cominated burst photosphere.," Therefore, if the flux of the burst on May 2004 is underestimated by, this translates into a distance of 6.8–6.5 kpc for a helium-dominated burst photosphere and of 5.2–5.0 kpc for a hydrogen-dominated burst photosphere."
+ These values are consistent with au upper Huit to the distance of & 7 kpc. required by plivsical EoSs (see Fig. 1).," These values are consistent with an upper limit to the distance of $\approxlt$ 7 kpc, required by physical EoSs (see Fig. \ref{fig:contours}) )."
+ We analysed four obscrvations of sspanning 19 mouths after a sienificant decrease of accretion ou the NS was detected., We analysed four observations of spanning 19 months after a significant decrease of accretion on the NS was detected.
+ We couclucle that: - the emission above 3 keV decreased siguificautlv within six months of the cud of the accretion phase - the NS crust has a hieh thermal couductivity aud low nupuriv material - the NS crust is very likely close to reach thermal equilibrium with the core he NS has a low-mecitim mass aud is cooling by standard mechauisuis - a consistent solution between spectral fitting to the thermalspectra aud the heating curves requires a distauce of $6 kpe to GT6.., We conclude that: - the emission above 3 keV decreased significantly within six months of the end of the accretion phase - the NS crust has a high thermal conductivity and low impurity material - the NS crust is very likely close to reach thermal equilibrium with the core - the NS has a low-medium mass and is cooling by standard mechanisms - a consistent solution between spectral fitting to the thermalspectra and the heating curves requires a distance of $\approxlt$ 6 kpc to .
+optical mages in the same fields (Almaini et al 1997).,optical images in the same fields (Almaini et al 1997).
+ A constraint ou the bias parameter of X-ray sources from the CCF results can be derived by taking iuto account that a fraction fz2/3 of the CCF signal arises from sources clustered with the catalogued galaxies., A constraint on the bias parameter of X-ray sources from the CCF results can be derived by taking into account that a fraction $f\approx 2/3$ of the CCF signal arises from sources clustered with the catalogued galaxies.
+ As that contribution scales linearly with by. the local volue cinissivity scales XTom: ," As that contribution scales linearly with $b_X$, the local volume emissivity scales $\propto {3\over
+1+2b_X}$."
+Since the ACN-onlv. local elissivity is also 10°egcn2sἘν we can safely derive that ὃν«2 as otherwise the total volume eumissivitv will be significantly less than the ACN cuiissivity.," Since the AGN-only local emissivity is also $\sim 10^{39}\, {\rm
+erg}\, {\rm cm}^{-2}\, {\rm s}^{-1}$, we can safely derive that $b_X<2$ as otherwise the total volume emissivity will be significantly less than the AGN emissivity."
+ Iu the recent vears large. complete samples of X-ray selected. AGN have been built.," In the recent years large, complete samples of X-ray selected AGN have been built."
+ This has allowed the direct 1icasurement ofthe 3D spatial correlation fiction £(e) or these objects and its comparison with the spatial correlation function of galaxies selected at other wavebauds., This has allowed the direct measurement of the 3D spatial correlation function $\xi(r)$ for these objects and its comparison with the spatial correlation function of galaxies selected at other wavebands.
+ Carrera et al (1998) used two complete samples of X-ray selected AGN in pencil CALL survey regions. spanunius a wide redshift range (0<2 2) to search for clustering signals aud deriving its amplitude aud redshift evolution.," Carrera et al (1998) used two complete samples of X-ray selected AGN in pencil beam survey regions, spanning a wide redshift range $0<z<2$ ) to search for clustering signals and deriving its amplitude and redshift evolution."
+ Clustering is ond to be significaut at z<1., Clustering is found to be significant at $z<1$.
+ When the spatial correlation fiction is fitted oa standard power-law form €(r})=(11:LS (foy conioviug rj. it is seen hat comoving or slower clusterius evolution is LI)excluded. and that even for stable or linear erowth the values of ry permitted by the data are of the same order as the ones derived from clustering of ZRAS galaxies.," When the spatial correlation function is fitted to a standard power-law form $\xi(r)=(1+z)^p({r\over r_0})^{-1.8}$ (for comoving $r$ ), it is seen that comoving or slower clustering evolution is excluded, and that even for stable or linear growth the values of $r_0$ permitted by the data are of the same order as the ones derived from clustering of $IRAS$ galaxies."
+ Carrera ct al (1998) conclude that N-ray selected AGN are not significantly biased 0.7<by«2., Carrera et al (1998) conclude that X-ray selected AGN are not significantly biased $0.7<b_X<2$.
+ Akvlas. Pliouis Georgantopoulos (1999) have used the ROSAT All Sky Survey sources to derive a local angular correlation fiction from which they estimate a somewhat higher correlation leugth (ry7—9 Mpc). consistent with optically selected QSO clustering (La Franca et al 1998) and comoving clustering evolution.," Akylas, Plionis Georgantopoulos (1999) have used the $ROSAT$ All Sky Survey sources to derive a local angular correlation function from which they estimate a somewhat higher correlation length $r_0\sim 7-9\, {\rm Mpc}$ ), consistent with optically selected QSO clustering (La Franca et al 1998) and comoving clustering evolution."
+ The obvious weakness of this method is that it is not based on 3D but 2D data., The obvious weakness of this method is that it is not based on 3D but 2D data.
+ The method of auto-correlating the NRB intensity at various separations has becu, The method of auto-correlating the XRB intensity at various separations has been
+likely helium- and iron-rich white dwarfs with periods of ~ seconds.,likely helium- and iron-rich white dwarfs with periods of $\sim$ seconds.
+ However. dilfusion theory suggests that in stars of such high gravity. iron. would settle out of the lavers where it would be required to excite pulsations.," However, diffusion theory suggests that in stars of such high gravity, iron would settle out of the layers where it would be required to excite pulsations."
+ To understand the phenomenon we recall that the instability blue edge occurs when the thermal timescale becomes too small compared with the pulsation period. where 7 is the temperature in the excitation laver. ὃν is the specific heat per unit mass at constant volume. and Arn is the mass lving above the excitation laver (seeCox1980.Ch.10for details)..," To understand the phenomenon we recall that the instability blue edge occurs when the thermal timescale becomes too small compared with the pulsation period, where $T$ is the temperature in the excitation layer, $c_{\rm v}$ is the specific heat per unit mass at constant volume, and $\Delta m$ is the mass lying above the excitation layer \citep[see][Ch. 10 for details]{Cox80}."
+ Let us assume an ideal gas and that the opacity in the lavers above the excitation zone is expressed as with constants ay. a(O<a1). and 6.," Let us assume an ideal gas and that the opacity in the layers above the excitation zone is expressed as with constants $\kappa_0$, $a (0\le a \le 1)$, and $b$."
+ Then. we can integrate. approximately. [from the stellar surface to. the excitation zone. the hydrostatie equation ancl the equation of cdilfusive Dow of radiation to obtain where ge is the mean molecular weight and A is the stellar radius.," Then, we can integrate approximately, from the stellar surface to the excitation zone, the hydrostatic equation and the equation of diffusive flow of radiation to obtain where $\mu$ is the mean molecular weight and $R$ is the stellar radius."
+ Substituting equation (3)) into equation (1)). and talking into account ὃνx1 [we have The strong dependence of my on Zige is responsible for the blue and the red. edges of instability range.," Substituting equation \ref{eq_deltam}) ) into equation \ref{eq_tauthdef}) ), and taking into account $c_{\rm v} \propto 1/\mu$, we have The strong dependence of $\tau_{\rm th}$ on $T_{\rm eff}$ is responsible for the blue and the red edges of instability range."
+ Since the s-mechanism is optimal when the period. of pulsation is comparable to my. the instability range of a short (long) period mode is hotter (cooler) as can be seen in Figs.," Since the $\kappa$ -mechanism is optimal when the period of pulsation is comparable to $\tau_{\rm th}$, the instability range of a short (long) period mode is hotter (cooler) as can be seen in Figs."
+ 2 and 1.., \ref{Z-stable} and \ref{Fe-stable}.
+ Fora given L/M ratio. the blue-edge can be shifted ward by increasing my. by reducing so and the mean molecular weight 5.," For a given L/M ratio, the blue-edge can be shifted blue-ward by increasing $\tau_{\rm th}$, by reducing $\kappa_0$ and the mean molecular weight $\mu$."
+ Phe former can be achieved by reducing Z without significantly reducing the iron abundance., The former can be achieved by reducing $Z$ without significantly reducing the iron abundance.
+ The mean molecular weight is reduced. by increasing .X., The mean molecular weight is reduced by increasing $X$.
+ The ellect is somewhat complicated: because an increase in .V tends to increase #&y., The effect is somewhat complicated because an increase in $X$ tends to increase $\kappa_0$.
+ Figs., Figs.
+ 2 and 4. seem to indicate that the ellect of reducing the mean molecular weight exceeds the ellect of increasing opacity., \ref{Z-stable} and \ref{Fe-stable} seem to indicate that the effect of reducing the mean molecular weight exceeds the effect of increasing opacity.
+ Predieted pulsation periods for unstable fundamental racial niodes. and for first. second. ancl third. racial harmonics we shown in Lig.," Predicted pulsation periods for unstable fundamental radial modes, and for first, second and third radial harmonics are shown in Fig."
+ 6 for a range of compositions and [or logL/AÀA4.=1.2 and 3.," \ref{Periods} for a range of compositions and for $\log L/M=1,
+2$ and 3."
+ The central column represents mocels with a standard. metallicity and reduced hydrogen abundance. while the left and. right columns represent a uniform metal enhancement and a selective iron enhancement. respectively.," The central column represents models with a standard metallicity and reduced hydrogen abundance, while the left and right columns represent a uniform metal enhancement and a selective iron enhancement respectively."
+ Note again that the lower L/Al pulsators for a selective iron enhancement are bluer than for a uniform enhancement. and. also that higher-order modes are preferentially excited in bluer pulsators.," Note again that the lower $L/M$ pulsators for a selective iron enhancement are bluer than for a uniform enhancement, and also that higher-order modes are preferentially excited in bluer pulsators."
+ For comparison. the top centre panel shows the pulsation periods and period ranges for several LC14026 and extreme heliumestar pulsators.," For comparison, the top centre panel shows the pulsation periods and period ranges for several EC14026 and extreme helium-star pulsators."
+ A consequence of the models is that bluer EC14026 stars should: show higher harmonies and. conversely. there should be fewer hot sdBVS pulsating in the racial fundamental mode.," A consequence of the models is that bluer EC14026 stars should show higher harmonics and, conversely, there should be fewer hot sdBVS pulsating in the radial fundamental mode."
+ This is clearly rellected by the meanslope of the logP(logTi) relation for EC14026 stars (excluding 11605|072 and 11930|2752) being steeper than that for a single mode.," This is clearly reflected by the meanslope of the $\log P(\log
+T_{\rm eff})$ relation for EC14026 stars (excluding 1605+072 and 1930+2752) being steeper than that for a single mode."
+ Fig., Fig.
+ 7. compares the observations for. LEC14026. variables with pulsation periods for unstable radial modes with Al=0(0.50M.. NX=0.75.Z0.02 and enhanced iron (f=10): logLfAL=LS was chosen to be representative of the EC14026 variables (see bie. 1)).," \ref{Blowup} compares the observations for EC14026 variables with pulsation periods for unstable radial modes with $M=0.50\Msolar$ , $X=0.75, Z=0.02$ and enhanced iron $f=10$ ); $\log L/M = 1.8$ was chosen to be representative of the EC14026 variables (see Fig. \ref{sdbv_obs}) )."
+ Ht is evident that only the coolest stars are likely to pulsate in the radial funcamental mode while many of the hottest can only excite the second or third harmonic., It is evident that only the coolest stars are likely to pulsate in the radial fundamental mode while many of the hottest can only excite the second or third harmonic.
+ Some EC14026 stars lic bevond the edee for all of these models: a higher value off might have been more appropriate. but the overall result is unallected.," Some EC14026 stars lie beyond the blue-edge for all of these models; a higher value of$f$ might have been more appropriate, but the overall result is unaffected."
+ Two exceptions are 11605|072 and 110502752., Two exceptions are 1605+072 and 1930+2752.
+ Phe first. has evolved awav from the extreme horizontal-branch and has a higher L/M ratio (lig. 1)., The first has evolved away from the extreme horizontal-branch and has a higher $L/M$ ratio (Fig. \ref{sdbv_obs}) ).
+ The second. is tidallv distorted. by à very close binary companion (Billeresetal.2000) ancl our classical racial models may be inadequate., The second is tidally distorted by a very close binary companion \citep{Bil00} and our classical radial models may be inadequate.
+ An indication of the overall instability is given in Fig., An indication of the overall instability is given in Fig.
+ δ where the exponential growth times are shown as contour plots for compositions representative of extremehelium and 1C14026. variables.," \ref{Growth}
+ where the exponential growth times are shown as contour plots for compositions representative of extremehelium and EC14026 variables."
+ The growth times are represented: by the dimensionless quantity —v f/w;., The growth times are represented by the dimensionless quantity $-\omega_r/\omega_i$ .
+ Divide by 2x to find the growth time in pulsation eveles. ancl then multiply by pulsation period. P? to obtain the physical erowth time Ty.," Divide by $2\pi$ to find the growth time in pulsation cycles, and then multiply by pulsation period $P$ to obtain the physical growth time $\tau_g$."
+ Thus. for the EC14026 stars with Af=0.50M. and logL/AM~L3. Fig.," Thus, for the EC14026 stars with $M=0.50\Msolar$ and $\log L/M\sim1.7$, Fig."
+ ο indicates logP/s~2 and Fig.," \ref{Periods} indicates $\log
+P/{\rm s}\sim 2$ and Fig."
+" S indicates log(ωνω)4. vielding log7,/s~5."," \ref{Growth} indicates $\log
+(-\omega_r/\omega_i)\sim4$, yielding $\log \tau_g/{\rm s}\sim5$."
+ For the extreme helium stars. the same exercise also gives log 5.," For the extreme helium stars, the same exercise also gives $\log \tau_g/{\rm
+ s}\sim5$ ."
+ With growth. times  one day. these modes are highly. unstable.," With growth times $\sim$ one day, these modes are highly unstable."
+ The non-racdial modes most usually observed. in stars may be divided into pressure or p-moces and gravity or gamocdes., The non-radial modes most usually observed in stars may be divided into pressure or p-modes and gravity or g-modes.
+ Pressure modes are trapped in the outer stellar envelope and may be thought of as a three dimensional generalization of the racial modes. ciseussecl in the previous section., Pressure modes are trapped in the outer stellar envelope and may be thought of as a three dimensional generalization of the radial modes discussed in the previous section.
+ Therefore their. stability has essentially been discussed already. although the precise periods of p-moces with spherical degree /2O will differ [rom the corresponding racial = 0) values.," Therefore their stability has essentially been discussed already, although the precise periods of p-modes with spherical degree $l>0$ will differ from the corresponding radial $l=0$ ) values."
+ Gravity modes propagate most strongly in the deep interior of the star., Gravity modes propagate most strongly in the deep interior of the star.
+ Consequently. complete stellar models including both core ancl envelope. are required., Consequently complete stellar models including both core and envelope are required.
+" Following previous work (Fontainectal.2003) we have assumed that the POGITIG6 variables are extreme horizontal branch stars and have thus constructed. series of ""zero-age"" horizontal- models having a helium-burning core with a mass of AJ. and with a hvdrogen-rich. envelope with mass A. ranging from 0.0003 to 0.0034 M.", Following previous work \citep{Fon03} we have assumed that the PG1716 variables are extreme horizontal branch stars and have thus constructed series of “zero-age” horizontal-branch models having a helium-burning core with a mass of $M_{\rm c}$ and with a hydrogen-rich envelope with mass $M_{\rm e}$ ranging from 0.0003 to 0.0034 .
+.. We have considered two cases of A: 0.476 and0.486 M... The surface. lavers," We have considered two cases of $M_{\rm c}$ ; $0.476$ and$0.486\Msolar$ , The surface layers"
+" (vy,2103 3) (Meier&Turner2"," $n_{H_{2}}> 10^{4}$ $^{-3}$ \citep[e.g.,][]{GS04,WEGSSV05}. \citep[][]{MT05}, \citep[e.g.,][]{I90,DRGGGM92}."
+005j..," \citep[e.g.,][]{GMFN00,GMFN01,GMFUN02,UGFMR04,UGMFN06}
+ \citep[e.g.,][]{MT05,P07,Bet08,NCSDHW08,KNGMCGE08,GGPFU08}."
+ is optically thick iu most starburst environments (6.8.al. 2007).. has a widely spaced. rotational ladder with iJ transitious in the subnun and can be scusitive to chemical effects (PDReffectsandIRpuimpine:e.g.Fuenteetal.1993:Aaltoct 2002).," is optically thick in most starburst environments \citep[e.g.][]{DRGGGM92,MT04, KWWBRM07}, has a widely spaced rotational ladder with its $>$ 3 transitions in the submm and can be sensitive to chemical effects \citep[PDR effects and IR pumping;
+e.g.][]{FMCB93,APHC02}."
+. Tere we present a study of the physical conditions of the very deuse gas componcut of a nearby starburst uucleus using IIC4N. With its large electric dipole ionmient (jj = 3.72 debywe versus 3.0 debye for ICN). low opacity and closely spaced rotational ladder accessible to powerful ceutimeter and mullimeterwave interferometers. IIC;N iswell-suited to high resolution inuagiueg of the structure and excitation of the densest conrponeut of the molecular eas (ee.Morrisetal.1976:VandenBoutetal.1983:Aladro 2011).," Here we present a study of the physical conditions of the very dense gas component of a nearby starburst nucleus using $_{3}$ N. With its large electric dipole moment $\mu$ = 3.72 debye versus 3.0 debye for HCN), low opacity and closely spaced rotational ladder accessible to powerful centimeter and millimeter-wave interferometers, $_{3}$ N iswell-suited to high resolution imaging of the structure and excitation of the densest component of the molecular gas \citep[eg.][]{MTPZ76,VLSW83,AMMMB11}."
+" TIC3N(5 has a critical density of 5«105 ? aud au upper level euergy of E/k = 6.6 IK: the 1615 line has a critical deusity of 5«10° cua? aud ai upper level enerev of E/k = 59k. We have observed the J = 5 Land J = 1615 hues of TIC3N at <2” rresolution in IC 312 with the Very Large (VLAÀ) and the Plateau de Bure Interferometer (PdD). respectively,"," $_3$ N(5--4) has a critical density of $5\times 
+10^{4}$ $^{-3}$ and an upper level energy of E/k = 6.6 K; the 16–15 line has a critical density of $5\times 10^{5}$ $^{-3}$ and an upper level energy of E/k = 59 K. We have observed the J = 5–4 and J = 16–15 lines of $_{3}$ N at $<$ resolution in IC 342 with the Very Large (VLA) and the Plateau de Bure Interferometer (PdB), respectively."
+ IC 312 is one of the closest (D~3Mpc.or 2005).. large spirals with active nuclear star formation(Becklin 1983)..," IC 342 is one of the closest \citep[D$\sim$3 Mpc, or 2\arcsec\ = 29 pc;
+eg.][]{SCH02,K05}, , large spirals with active nuclear star formation\citep[][]{BGMNSWW80,TH83}. ."
+ It is nearly face-on. with bright molecular lineciuission both iu deuse clouds aud a diffuse medii," It is nearly face-on, with bright molecular lineemission both in dense clouds and a diffuse medium"
+2011).,.
+. Iu particular. Zheng. Coil Zehavi (2007: hereafter ZCZOT) develop a plenomenological approach to extract information about ealaxy evolution from 2.~1 to ~0 bv performing TOD modeling of two-point correlation fictions of DEEP2 and SDSS ealaxies.," In particular, Zheng, Coil Zehavi (2007; hereafter ZCZ07) develop a phenomenological approach to extract information about galaxy evolution from $z\sim 1$ to $z\sim 0$ by performing HOD modeling of two-point correlation functions of DEEP2 and SDSS galaxies."
+ With the iuferred ealaxy-halo connection at two redshifts. they establish an evolutionary link usine the typical erowth of dark matter halos obtained from munerical simulations.," With the inferred galaxy-halo connection at two redshifts, they establish an evolutionary link using the typical growth of dark matter halos obtained from numerical simulations."
+ Even with the progress made iu establishing the evolutionary link between galaxies and halos. our uuderstaudius of the specifies of stellar mass growth within the dark matter halos is still far from complete.," Even with the progress made in establishing the evolutionary link between galaxies and halos, our understanding of the specifics of stellar mass growth within the dark matter halos is still far from complete."
+ Calaxies can grow their stellar mass by star formation. accretion of sinaller satellite ealaxies or major mereiug.," Galaxies can grow their stellar mass by star formation, accretion of smaller satellite galaxies or major merging."
+ It is nuportaut to quantity the contribution of all these Xocesses dn order to have a complete picture of the assenmablv listory of salaxies within their lost dark uatter halos., It is important to quantify the contribution of all these processes in order to have a complete picture of the assembly history of galaxies within their host dark matter halos.
+ ZCZOT derive the mean stellar masses of central galaxies at +~1 aud :~O as a function of je present-day host halo mass., ZCZ07 derive the mean stellar masses of central galaxies at $z\sim1$ and $z\sim0$ as a function of the present-day host halo mass.
+ After roughly accounting or the contribution of mereius of ceutral aud satellite 43valanics. they inter the star-formation contribution to the ellar mass assembly.," After roughly accounting for the contribution of merging of central and satellite galaxies, they infer the star-formation contribution to the stellar mass assembly."
+ They &ud that in ceutral galaxies cated iu relatively low-1ass halos (~5<Lolth HAE.) ιο bulk of the stellar mass results frou star formation ποποσα Do Land 2~0. while only a simall fraction of stars formed since ;~1l du central galaxies of halos as nassive as ~5&10025.TALL.," They find that in central galaxies located in relatively low-mass halos $\sim 5\times 10^{11} \Msunh$ ) the bulk of the stellar mass results from star formation between $z\sim1$ and $z\sim 0$, while only a small fraction of stars formed since $z\sim1$ in central galaxies of halos as massive as $\sim 5\times 10^{12} \Msunh$."
+ For these massive lalos. ucreiue becomes more dnportant aud constitutes the dominant contribution to the stellar mass erowth (see Fie.," For these massive halos, merging becomes more important and constitutes the dominant contribution to the stellar mass growth (see Fig."
+ 9 in ZCZOT for details}., 9 in ZCZ07 for details).
+ The results reflect the so-called “downsizing” star formation pattern in which 1ο sites of active star formation shift from highiauass ealaxies at carly times to lower-iass svstenis at later epochs (e.e.. Cowieetal. 1996)). mmanitested iu ferius of jalo lass.," The results reflect the so-called “downsizing” star formation pattern in which the sites of active star formation shift from high-mass galaxies at early times to lower-mass systems at later epochs (e.g., \citealt{Cowie96}) ), manifested in terms of halo mass."
+ Iu this paper. we study the theoretical predictions for stellar mass evolution as a function of dark matter halo nass using SAX catalogs.," In this paper, we study the theoretical predictions for stellar mass evolution as a function of dark matter halo mass using SAM catalogs."
+ One of the main advantages of the SANs is that we can trace the full evolution of 1ο individual galaxies within their dark matter halos. allowing au explicit study of the different processes that contribute to the build up of the stellar couteut of ealaxies.," One of the main advantages of the SAMs is that we can trace the full evolution of the individual galaxies within their dark matter halos, allowing an explicit study of the different processes that contribute to the build up of the stellar content of galaxies."
+ We compare and contrast these predictions with the results of ZCZüT., We compare and contrast these predictions with the results of ZCZ07.
+ We gauge the potential of the phenomenological methods to coustrain galaxy formation models. as well as test some of the assumptions of such methods.," We gauge the potential of the phenomenological methods to constrain galaxy formation models, as well as test some of the assumptions of such methods."
+ Iu particular. we check the validity of the simple evolutionary. approach preseuted by 6201.," In particular, we check the validity of the simple evolutionary approach presented by ZCZ07."
+ For these purposes. we use two SAAD catalogs. the “MIPA™ (Crotonetal.2006:DeLucia& Blaizot 2007)) aud “Durham” (Bowerctal.2006:Font 2008)) catalogs produced from the Mallenninm Run costuological N-body simulation (Spriugeletal. 2005)).," For these purposes, we use two SAM catalogs, the “MPA” \citealt{Croton06, DeLucia07}) ) and “Durham” \citealt{Bower06, font08}) ) catalogs produced from the Millennium Run cosmological -body simulation \citealt{Springel05}) )."
+" We mainly focus on the stellar mass growth as a function of halo mass since 2~1 for case of comparison with 26701,", We mainly focus on the stellar mass growth as a function of halo mass since $z\sim 1$ for ease of comparison with ZCZ07.
+ Nonetheless. we briefly study also the stellar mass erowth since 2~ in the MPA catalog to explore the processes involved iu galaxy formation at higher redshifts.," Nonetheless, we briefly study also the stellar mass growth since $z\sim 2$ in the MPA catalog to explore the processes involved in galaxy formation at higher redshifts."
+ These SAM catalogs have been previously used to study the stellar mass evolution in galaxies (6.9.. Guo&White2008:Strinecrctal.Foutanot 2009)).," These SAM catalogs have been previously used to study the stellar mass evolution in galaxies (e.g., \citealt{Guo08,Stringer09,Fontanot09}) )."
+ ILoxvever. those studies focus on the iutegrated stellar mass to make a direct comparison to observations and did not investigate iu detail its evolution as a functionof halo mass.," However, those studies focus on the integrated stellar mass to make a direct comparison to observations and did not investigate in detail its evolution as a function of halo mass."
+ Most plivsical processes involved im ealaxy formation models depend explicitly ou the mass of the dark matter halo., Most physical processes involved in galaxy formation models depend explicitly on the mass of the dark matter halo.
+ Additionally. the modeling of other observables such as the ealaxy-galaxy merecr rate strouely rely ou the precise knowledge of the ealaxy-halo connection (see e.g. Topkinsetal.2010)).," Additionally, the modeling of other observables such as the galaxy-galaxy merger rate strongly rely on the precise knowledge of the galaxy-halo connection (see e.g., \citealt{Hopkins10}) )."
+ Thus it is physically 1ueanimgful aud informative to study galaxy evolution as a function of halo mass., Thus it is physically meaningful and informative to study galaxy evolution as a function of halo mass.
+ The paper is organized as follows: 822 describe the SAMI catalogs that we use auc the sample selection., The paper is organized as follows: 2 describe the SAM catalogs that we use and the sample selection.
+ Iu 833 we present and discuss our results for the erowth of stellar mass as a function of halo mass., In 3 we present and discuss our results for the growth of stellar mass as a function of halo mass.
+ Iu bl we compare our results to those fouud by ZCZUÜT aud we couclude i 855., In 4 we compare our results to those found by ZCZ07 and we conclude in 5.
+) Tn this work. we use the publicly available mocks ealaxv catalogs produced with the ~MPA™ (Crotouet 2007)) aud “Durham (Bowerctal. 20061) soimi-analvtie niodels of ealaxy formation. both based upon dark matter halo evolution oe1i the Milleuniuuni Run simulation (Spriugeletal.2005)).," In this work, we use the publicly available mock galaxy catalogs produced with the “MPA” \citealt{Croton06, DeLucia07}) ) and “Durham” \citealt{Bower06}) ) semi-analytic models of galaxy formation, both based upon dark matter halo evolution in the Millennium Run simulation \citealt{Springel05}) )."
+" The Millennium Run followed the evolution of ~101"" ark inafter particles in a Universe.", The Millennium Run followed the evolution of $\sim 10^{10}$ dark matter particles in a Universe.
+ The VAiunlation uses a periodic box of 500%tAIpe ou a side with mass resolution per particle of 8.6«1055.TAL..., The simulation uses a periodic box of $500 \mpch$ on a side with mass resolution per particle of $8.6 \times 10^8 \Msunh$.
+ The oeutial couditious of the simulation were ecucrated with 'osnological paraiucters obtained from the combined alysis of 24FCRS aud WALTAPL CMB data., The initial conditions of the simulation were generated with cosmological parameters obtained from the combined analysis of 2dFGRS and WMAP1 CMB data.
+ The halos oei the simulation were identified m cach time step using a friend-offricuds algorithim with a linking leneth of 0.2 ines the mean particle separation., The halos in the simulation were identified in each time step using a friend-of-friends algorithm with a linking length of 0.2 times the mean particle separation.
+ More details can be ound in Spriugeletal.2005.., More details can be found in \citealt{Springel05}.
+ Iu the SAMs. galaxies are assumed to form at the ceuter of dark iatter halos.," In the SAMs, galaxies are assumed to form at the center of dark matter halos."
+ The evolution of the xuvonie component of ealaxies is modeled using simple mt plysically motivated analytic prescriptions., The evolution of the baryonic component of galaxies is modeled using simple but physically motivated analytic prescriptions.
+" These include radiative cooling of lot gas. star formation in he cold disk. supernova feedback. black hole growth aud ACN feedback through the ""quasi and vracdio” epochs of ACN evolution. metal euricluneut of the inter-ealactic and iutra-cluster meditu. as well as ealaxy morphology shaped through merecrs and mereer-induced starbursts."," These include radiative cooling of hot gas, star formation in the cold disk, supernova feedback, black hole growth and AGN feedback through the “quasar” and “radio” epochs of AGN evolution, metal enrichment of the inter-galactic and intra-cluster medium, as well as galaxy morphology shaped through mergers and merger-induced starbursts."
+ As incutioned above. these models are aimed at reproducing inteerated ealaxy observables.," As mentioned above, these models are aimed at reproducing integrated galaxy observables."
+" To that effect. the SAMS recover reasonably well the ealaxy huuinosity function in different bands (οιο,, Crotonctal.2006:Bowerct 2006)) as well as the stellar mass functions for a range of redshifts (c.g.. BowerotFoutanotetal. 2009: but see also Li&White2009:Marchesinietal. 20093)."," To that effect, the SAMs recover reasonably well the galaxy luminosity function in different bands (e.g., \citealt{Croton06, Bower06}) ) as well as the stellar mass functions for a range of redshifts (e.g., \citealt{Bower06, Fontanot09}; ; but see also \citealt{li09,Marchesini09}) )."
+ The predictions for the fraction of blue aud red ceutral galaxies. however. are not fullv correct (see Baldryetal. 2006)). and modifications to the treatineut of gas cooling (e... Violaetal. 20083) and the AGN feedback (6.2... Baldryetal. 2006)) have been proposed to alleviate these discrepancies.," The predictions for the fraction of blue and red central galaxies, however, are not fully correct (see \citealt{Baldry06}) ), and modifications to the treatment of gas cooling (e.g., \citealt{Viola08}) ) and the AGN feedback (e.g., \citealt{Baldry06}) ) have been proposed to alleviate these discrepancies."
+ The SAMs also appear to overestimate the fraction of red satellites (Weinmannetal.2006:Coil2008:ane2009)).," The SAMs also appear to overestimate the fraction of red satellites \citealt{Weinmann06,Coil08, Kang08, DeLucia09, SeekKim09}))."
+ Note that satellite galaxies iu these modols are treated differently than central galaxies., Note that satellite galaxies in these models are treated differently than central galaxies.
+ Ouly ceutral, Only central
+A plot of the integrated radio continuum spectrum for the SNR IC 443 ts shown in Fig. 7..,A plot of the integrated radio continuum spectrum for the SNR IC 443 is shown in Fig. \ref{global-spectrum}.
+ Our new integrated flux measurements at 74 and 330 MHz are indicated by filled circle symbols., Our new integrated flux measurements at 74 and 330 MHz are indicated by filled circle symbols.
+unstructured halo population.,unstructured halo population.
+ A simple Kolmogorov-Smirnov test provides an estimate of the significance of this result: the probability that the two distributions were drawn from the same parent distribution is only +1.3%., A simple Kolmogorov-Smirnov test provides an estimate of the significance of this result: the probability that the two distributions were drawn from the same parent distribution is only $\approx1.3\%$.
+ This calculation provides a clear quantitative straightforwarddemonstration that the global spatial coincidence between globular clusters and tidal debris streams visible in Fig., This straightforward calculation provides a clear quantitative demonstration that the global spatial coincidence between globular clusters and tidal debris streams visible in Fig.
+ | is almost certainly not due to chance alignment., 1 is almost certainly not due to chance alignment.
+ Further inspection of Fig., Further inspection of Fig.
+" 1 suggests a more complicated some streams possess more clusters for their picture:size than do others, while apparentlyseveral are devoid of clusters altogether."," 1 suggests a more complicated picture: some streams apparently possess more clusters for their size than do others, while several are devoid of clusters altogether."
+ To investigate this in more detail we considered each major halo overdensity individually., To investigate this in more detail we considered each major halo overdensity individually.
+" For now, we restrict ourselves to those stellar substructures previously published in the literature."," For now, we restrict ourselves to only those stellar substructures previously published in the literature."
+only These are highlighted in Fig., These are highlighted in Fig.
+" 1: the giant stellar stream (Ibataetal.2001);; the northeast structure (Zuckeretal.2004;Ibata2005);; the four minor-axis tangent streams A-D (Ibataetal.2007),, the latter of which extends to a coherent arc to the east of M31; and the southwest cloud and northwest stream (McConnachieetal."," 1: the giant stellar stream \citep{ibata:01}; the northeast structure \citep{zucker:04,ibata:05}; the four minor-axis tangent streams A-D \citep{ibata:07}, the latter of which extends to a coherent arc to the east of M31; and the southwest cloud and northwest stream \citep{mcconnachie:09}."
+" 2009).. Ibataetal.(2007) also describe the “major- diffuse structure"", an ill-defined feature extending to the southwest of M31 (region 1 in Fig."," \citet{ibata:07} also describe the ``major-axis diffuse structure”, an ill-defined feature extending to the southwest of M31 (region 1 in Fig."
+ 1)., 1).
+" Although there is evidently an overdensity along the major axis here, unlike the other streams it is not identifiable as a structure; indeed there are additional easilyfaint overdensities singleextending east and north of this feature."," Although there is evidently an overdensity along the major axis here, unlike the other streams it is not easily identifiable as a single structure; indeed there are additional faint overdensities extending east and north of this feature."
+" We defer analysis of this region until it has been properly characterized, but note that numerous clusters (2:8— 10) project onto it."," We defer analysis of this region until it has been properly characterized, but note that numerous globular clusters $\approx8-10$ ) clearly project onto it."
+" We globularutilized FITS versions of clearlyboth our metal-poor stellar density map, and an additional map including stars with photometric metallicities -1.4< [Fe/H]S-0.7, to delineate the edges of each substructure."," We utilized FITS versions of both our metal-poor stellar density map, and an additional map including stars with photometric metallicities $-1.4\la[$ $/$ $]\la-0.7$, to delineate the edges of each substructure."
+" These more metal-rich stars alter the appearance of several of the overdensities (e.g.,Ibataetal.2007;McConnachie 2009):: most notably the giant stream, which increases in radial extent and fans westward; the southwest cloud, which becomes more prominent; and stream C, which is considerably broadened — indeed, this feature is known to consist of two distinct overlapping (Chapmanetal.2008)."," These more metal-rich stars alter the appearance of several of the overdensities \citep[e.g.,][]{ibata:07,mcconnachie:09}: most notably the giant stream, which increases in radial extent and fans westward; the southwest cloud, which becomes more prominent; and stream C, which is considerably broadened – indeed, this feature is known to consist of two distinct overlapping components \citep{chapman:08}."
+". For each componentssubstructure we measured the mean and standard deviation of pixels in numerous nearby regions, and defined the edge of the structure by following a contour level above the local "," For each substructure we measured the mean and standard deviation of pixels in numerous nearby regions, and defined the edge of the structure by following a contour level above the local background."
+No globular clusters were overplotted during thisbackground. process., No globular clusters were overplotted during this process.
+ Our results are shown in Fig., Our results are shown in Fig.
+ 3., 3.
+" Although the edges of these substructures are by nature difficult to define, they are all sufficiently unambiguous not to alter the conclusions we draw below."," Although the edges of these substructures are by nature difficult to define, they are all sufficiently unambiguous not to alter the conclusions we draw below."
+" We next counted the fraction of mock systems in which or more globular clusters overlap spatially with a given Ναςsubstructure, where Νας is the number of clusters observed to project onto that feature in the real M31 halo."," We next counted the fraction of mock systems in which $N_{{\rm gc}}$ or more globular clusters overlap spatially with a given substructure, where $N_{{\rm gc}}$ is the number of clusters observed to project onto that feature in the real M31 halo."
+" We also considered the system globally, grouping together all the identified substructures."," We also considered the system globally, grouping together all the identified substructures."
+ Our results are summarized in Table 1., Our results are summarized in Table 1.
+" Taking all substructures together, we find it very unlikely that the observed spatial overlap of 27 clusters with these features can be explained by random alignment: the probability sits at just ~0.25%."," Taking all substructures together, we find it very unlikely that the observed spatial overlap of $27$ clusters with these features can be explained by random alignment: the probability sits at just $\sim0.25\%$."
+ This strongly reinforces the result of our previous calculation involving the average local stellar densities., This strongly reinforces the result of our previous calculation involving the average local stellar densities.
+" Individually, the northwest stream and eastern arc are particularly well endowed with clusters."," Individually, the northwest stream and eastern arc are particularly well endowed with clusters."
+ The fraction of mock in which >6 clusters fall within the northwest stream is systems~ 2.7%; this falls to below 1% if the additional boundary-straddling cluster to the south is included., The fraction of mock systems in which $\ge6$ clusters fall within the northwest stream is $\sim2.7\%$ ; this falls to below $1\%$ if the additional boundary-straddling cluster to the south is included.
+" Similarly, the frequency with which at least the 11 observed clusters fall within the eastern arc is ~0.5% in the mock systems."," Similarly, the frequency with which at least the $11$ observed clusters fall within the eastern arc is $\sim0.5\%$ in the mock systems."
+" For the southwest cloud our calculated probability sits at ~2.5%, while for stream C it is less significant at ~ 7.8%."," For the southwest cloud our calculated probability sits at $\sim2.5\%$, while for stream C it is less significant at $\sim7.8\%$ ."
+" Notably however, stream C is the only case where velocity information exists for both the field substructure and an associated cluster (EC4), unambiguously linking the two (Collinsetal.2009)."," Notably however, stream C is the only case where velocity information exists for both the field substructure and an associated cluster (EC4), unambiguously linking the two \citep{collins:09}."
+. The stream is notable as the structure where the observed giantnumber of globular clusters onlyapproximately matchesthe number expected in a smoothly-distributed system., The giant stream is notable as the only structure where the observed number of globular clusters approximately matchesthe number expected in a smoothly-distributed system.
+ Since this stream is by far the most luminous stellar substructure, Since this stream is by far the most luminous stellar substructure
+"The remaining energy dependant terms are now A(E) and 7(£), with A(E) now usefully being the “y-axis” flux scale factor and T(E) being the othogonal“x-axis” temporal scale factor.","The remaining energy dependant terms are now $A(E)$ and $\tau(E)$ , with $A(E)$ now usefully being the “y-axis"" flux scale factor and $\tau(E)$ being the othogonal“x-axis"" temporal scale factor."
+" A yet further conjecture ventured here, the Pulse Shape Conjecture, posits that a group of GRB pulses all have the same shape."," A yet further conjecture ventured here, the Pulse Shape Conjecture, posits that a group of GRB pulses all have the same shape."
+" In this paradigm, € is not only constant in a pulse across energies, an entire group of pulses all are hypothesized to have the same shape: €=1."," In this paradigm, $\xi$ is not only constant in a pulse across energies, an entire group of pulses all are hypothesized to have the same shape: $\xi = 1$."
+" Given that the Pulse Shape Conjecture holds, then Eq. (9) "," Given that the Pulse Shape Conjecture holds, then Eq. \ref{NorrisFull}) )"
+can be further simplified to It is suspected that this functional form is not unique in that other mathematical functions can also adequately fit all of the pulse conjectures and describe the same group of GRB pulses., can be further simplified to It is suspected that this functional form is not unique in that other mathematical functions can also adequately fit all of the pulse conjectures and describe the same group of GRB pulses.
+" In the above cross-energy pulse descriptions, the meaning of A(£) is interesting as it appears to be a type of spectrum for a pulse that involves the flux from the entire pulse."," In the above cross-energy pulse descriptions, the meaning of $A(E)$ is interesting as it appears to be a type of spectrum for a pulse that involves the flux from the entire pulse."
+" Note that A(E) will be different than a spectrum taken at a specific time for of the pulse, say at the time of peak energy of one of the energy channels."," Note that $A(E)$ will be different than a spectrum taken at a specific time for of the pulse, say at the time of peak energy of one of the energy channels."
+" Interestingly, A(E) can be fit from incomplete pulse information, for example computable when only a fractionof a single pulse is recoverable across energy bands."," Interestingly, $A(E)$ can be fit from incomplete pulse information, for example computable when only a fractionof a single pulse is recoverable across energy bands."
+" Previously, A(E) has only been determined for a single pulse in a single GRB: the main pulse in BATSE trigger 2193 (Nemiroff|2000)."," Previously, $A(E)$ has only been determined for a single pulse in a single GRB: the main pulse in BATSE trigger 2193 \citep{Nem00}."
+". There, A(E) was fit to 16 BATSE MER channels and was seen to peak at an energy of about 150 keV. Note that A(£) is intrinsically a type of “photon count spectrum”, as P(t) is a measurement of photon counts rather than energy flux."," There, $A(E$ ) was fit to 16 BATSE MER channels and was seen to peak at an energy of about 150 keV. Note that $A(E)$ is intrinsically a type of “photon count spectrum"", as $P(t)$ is a measurement of photon counts rather than energy flux."
+ The meaning of 7T(E) appears to be the pulse duration as a function of energy., The meaning of $\tau(E)$ appears to be the pulse duration as a function of energy.
+" Note that T(E) has (also) only been determined for single pulse in a single GRB, in fact the same pulse (BATSE trigger 2193) as with A(E) ∥"," Note that $\tau(E)$ has (also) only been determined for single pulse in a single GRB, in fact the same pulse (BATSE trigger 2193) as with $A(E)$ \citep{Nem00}."
+" There, 7T(E) appears approximated as a power law where T(E)ο.E97."," There, $\tau(E)$ appears approximated as a power law where $\tau(E) \propto E^{-0.7}$."
+ Informal observation of numerous GRB pulses indicates that 7(E) usually decreases monotonically with increasing energy., Informal observation of numerous GRB pulses indicates that $\tau(E)$ usually decreases monotonically with increasing energy.
+" Note that as with A(E), T(E) may also be fit from incomplete pulse information, for example computable when only a faction of a single pulse is recoverable across energy bands."," Note that as with $A(E)$, $\tau(E)$ may also be fit from incomplete pulse information, for example computable when only a faction of a single pulse is recoverable across energy bands."
+ Note that the total fluence of a Norris-model pulse is directly proportional to A(E) times 7(E)., Note that the total fluence of a Norris-model pulse is directly proportional to $A(E)$ times $\tau(E)$.
+" Since A(E) has a clear peak, pulses might be seen as envelopes effectively spanning only certain times and energies."," Since $A(E)$ has a clear peak, pulses might be seen as envelopes effectively spanning only certain times and energies."
+" In this section, an attempt is made to determine whether the Norris pulse model can be constrained by the Pulse Start, Pulse Scale, and Pulse Shape conjectures and still adequately fit isolated GRB pulses adequately across BATSE energy bands."," In this section, an attempt is made to determine whether the Norris pulse model can be constrained by the Pulse Start, Pulse Scale, and Pulse Shape conjectures and still adequately fit isolated GRB pulses adequately across BATSE energy bands."
+" Note that invoking these conjectures provides many additional constraints, so that many fewer free parameters will be available to fit the pulses across the BATSE energy bands."," Note that invoking these conjectures provides many additional constraints, so that many fewer free parameters will be available to fit the pulses across the BATSE energy bands."
+" As a consequence, one might expect fits to be less good than fits without these constraints."," As a consequence, one might expect fits to be less good than fits without these constraints."
+" Therefore, what is being tested is not whether such constraints provide better fits — but whether these more tightly constrained models can create acceptable fits at all."," Therefore, what is being tested is not whether such constraints provide better fits – but whether these more tightly constrained models can create acceptable fits at all."
+ BATSE data was used for several reasons., BATSE data was used for several reasons.
+" First, BATSE incorporated the largest set of GRB detectors ever flown, and so recorded relatively high count rates for GRBs."," First, BATSE incorporated the largest set of GRB detectors ever flown, and so recorded relatively high count rates for GRBs."
+" Next, BATSE, so far, was one of the longest running GRB detectors ever flown, and so over its decade of operation had time for wait for relatively bright GRBs to occur."," Next, BATSE, so far, was one of the longest running GRB detectors ever flown, and so over its decade of operation had time for wait for relatively bright GRBs to occur."
+" Next, BATSE data has now been around for over two decades and so is relatively well understood."," Next, BATSE data has now been around for over two decades and so is relatively well understood."
+" Next, BATSE data is in the public domain and easily available over the web in simple formats, for example ASCII."," Next, BATSE data is in the public domain and easily available over the web in simple formats, for example ASCII."
+" Last, created their bright isolated pulse catalog from BATSE data, which is a catalog with useful data from which this study can draw and compare."," Last, \citet{Nor05} created their bright isolated pulse catalog from BATSE data, which is a catalog with useful data from which this study can draw and compare."
+" In general, four channel BATSE LAD data were used."," In general, four channel BATSE LAD data were used."
+" The four channels used were channel 1 (20 - 50 KeV), channel 2 (50 - 100 KeV), channel 3 (100 - 300 KeV), and channel 4 (300 - 1000 MeV)."," The four channels used were channel 1 (20 - 50 KeV), channel 2 (50 - 100 KeV), channel 3 (100 - 300 KeV), and channel 4 (300 - 1000 MeV)."
+" In many GRBs, the counts in channel 4 are so low that the burst may not be easily discernible,"," In many GRBs, the counts in channel 4 are so low that the burst may not be easily discernible,."
+" Therefore, only the fits for the first three BATSE LAD energy channels are reported in Table 1."," Therefore, only the fits for the first three BATSE LAD energy channels are reported in Table 1."
+" The BATSE LAD data type with 64-ms time resolution was used, as it had the highest temporal resolution that covered greatest temporal range in and around pulses."," The BATSE LAD data type with 64-ms time resolution was used, as it had the highest temporal resolution that covered greatest temporal range in and around pulses."
+ The procedure for determining which GRB pulses were selected for fitting to the scalabale Norris pulse function was as follows., The procedure for determining which GRB pulses were selected for fitting to the scalabale Norris pulse function was as follows.
+ All of the GRBs in the [Norrisetall(2005) catalog were considered initially., All of the GRBs in the \citet{Nor05} catalog were considered initially.
+" These pulses were considered initially by because they were isolated, of relatively long duration, were relatively fluent, and had relatively long lag times between cross correlations between BATSE energy channels."," These pulses were considered initially by \citet{Nor05} because they were isolated, of relatively long duration, were relatively fluent, and had relatively long lag times between cross correlations between BATSE energy channels."
+ A visual inspection was then done excluding some GRBs that were subjectively deemed too convolved with other pulses to yield a meaningful result., A visual inspection was then done excluding some GRBs that were subjectively deemed too convolved with other pulses to yield a meaningful result.
+" Admittedly, this procedure might bias the results so that only pulses that might fit are selected, but this was considered acceptable since a key premise being tested is that GRB pulses can be described by the energy-extended Norris pulse form, not all of them."," Admittedly, this procedure might bias the results so that only pulses that might fit are selected, but this was considered acceptable since a key premise being tested is that GRB pulses can be described by the energy-extended Norris pulse form, not all of them."
+" Once a GRB and its dominant pulse were selected, a range of times inside the GRB when the pulse being considered was clearly uncontaminated by surrounding pulses were recorded."," Once a GRB and its dominant pulse were selected, a range of times inside the GRB when the pulse being considered was clearly uncontaminated by surrounding pulses were recorded."
+ The pulse light curve was then rebinned in time so that the pulse rise time occurred between about 10 to 20 time bins., The pulse light curve was then rebinned in time so that the pulse rise time occurred between about 10 to 20 time bins.
+ This allowed the pulse to be temporally resolved while retaining a relatively large signal to noise per time bin., This allowed the pulse to be temporally resolved while retaining a relatively large signal to noise per time bin.
+" Two background intervals, typically lasting 100 seconds, were selected well before the pulse and after the pulse and fit with a single third degree polynomial."," Two background intervals, typically lasting 100 seconds, were selected well before the pulse and after the pulse and fit with a single third degree polynomial."
+" A possible range of reasonable start times, t;, were estimated based on visual inspection of the pulse light curves in the four BATSE LAD energy channels."," A possible range of reasonable start times, $t_o$, were estimated based on visual inspection of the pulse light curves in the four BATSE LAD energy channels."
+" To test for goodness of fit, candidate t, values in this range were assumed, in turn."," To test for goodness of fit, candidate $t_o$ values in this range were assumed, in turn."
+" For each {ο assumed, the best fit r and A values were recorded that minimized thex? per degree of freedom (v, the number of time bins) for each energy channel."," For each $t_o$ assumed, the best fit $\tau$ and $A$ values were recorded that minimized the$\chi^2$ per degree of freedom $\nu$ , the number of time bins) for each energy channel."
+" The value of t, that allowed the best fit A and 7 values were recorded, as well as error estimates for these values."," The value of $t_o$ that allowed the best fit $A$ and $\tau$ values were recorded, as well as error estimates for these values."
+0.5 truecm Strong magnetic fields of neutron stars. tvpicallv of order ~107 G (Bhattacharva. Srinivasan 1995) ave considered crucial to the radio pulsar activity but ecuilibrium arguments concerning NS magnetic field pressure and ram pressure balance lead to magnetic field strength estimates [or heavily accreting NS sources (for example X-ray. bursters) of <10? G. In (hese sources no persistent pulsations have been found despite careful searches (Wood el al.,"0.5 truecm Strong magnetic fields of neutron stars, typically of order $\sim 10^{12}$ G (Bhattacharya, Srinivasan 1995) are considered crucial to the radio pulsar activity but equilibrium arguments concerning NS magnetic field pressure and ram pressure balance lead to magnetic field strength estimates for heavily accreting NS sources (for example X-ray bursters) of $\ltorder10^9$ G. In these sources no persistent pulsations have been found despite careful searches (Wood et al."
+ 1991: Vaughan et al., 1991; Vaughan et al.
+ 1994)., 1994).
+ Different. scenarios of magnetic field evolution lead to different estimates of magnetic field strength and decay., Different scenarios of magnetic field evolution lead to different estimates of magnetic field strength and decay.
+ Thus. it is desirable to find a direct wav of measuring or estimating the neutron stars magnetic field in these binary svstens.," Thus, it is desirable to find a direct way of measuring or estimating the neutron star's magnetic field in these binary systems."
+ Anv model used to make such estimates should have minimum assumptions about observed data and model relationships to produce accurate results., Any model used to make such estimates should have minimum assumptions about observed data and model relationships to produce accurate results.
+ lt develops a new probe of the magnetic field strength and geometry. based on radial magnetoacoustie oscillations.," It develops a new probe of the magnetic field strength and geometry, based on radial magnetoacoustic oscillations."
+ Titarchuk. Lapidus Muslimov. (1993). hereafter TLMO98. considered. (he possibility of dvnanmical adjustments of a Ixeplerian disk (to (he innermost sub-Ixeplerian. boundary," Titarchuk, Lapidus Muslimov (1998), hereafter TLM98, considered the possibility of dynamical adjustments of a Keplerian disk to the innermost sub-Keplerian boundary"
+Hu Section L. aud for the reddening m Section 5..,"in Section \ref{RotationSection}, and for the reddening in Section \ref{ReddeningSection}."
+ Finally iu Section 6 we consider how the range of distances proposed affect estimates of the outburst aud quiescenut AX-aayv Inuuinuosities aud the nature of the companion star. and in Section 7 we sumunarize our conclusions.," Finally in Section \ref{DiscussionSection} we consider how the range of distances proposed affect estimates of the outburst and quiescent X-ray luminosities and the nature of the companion star, and in Section \ref{ConclusionSection} we summarize our conclusions."
+ Tieh-+vesolution observations of wwere obtained with the UV-Visual Echelle Spectrograph (UVES) on UT2 (RKueyou) at the VLT ou 2002 Aueust 915., High-resolution observations of were obtained with the UV-Visual Echelle Spectrograph (UVES) on UT2 (Kueyen) at the VLT on 2002 August 9–15.
+ These spanned a waveleneth range of aand included a total exposure time of bly vieldiug a superb quality average spectrun (signal-to-noise ratio near the Ines greater than 100 without binning)., These spanned a wavelength range of and included a total exposure time of hr yielding a superb quality average spectrum (signal-to-noise ratio near the lines greater than 100 without binning).
+" A 17 slit was used eiviug a nomiual resolution near the ]liues of z £1000. corresponding to a kincmatic resolution of 5,"," A 1” slit was used giving a nominal resolution near the lines of $\simeq44\,000$, corresponding to a kinematic resolution of $^{-1}$."
+ The starlight will be concentrated to the ceuter of the slit. however. vieldiug a somewhat ligher resolution.," The starlight will be concentrated to the center of the slit, however, yielding a somewhat higher resolution."
+ From the spatial profiles near wwe estinate a ican ποσο of about 0.57. muplviug that the resolution could be as good as | with a 1 st.," From the spatial profiles near we estimate a mean seeing of about 0.5”, implying that the resolution could be as good as $^{-1}$ with a 1” slit."
+ In practice. this will not be realized and we expect a value between these extremes.," In practice, this will not be realized and we expect a value between these extremes."
+ We therefore leave the instrumental resolution as a free parameter m subsequeut fits., We therefore leave the instrumental resolution as a free parameter in subsequent fits.
+ For the rest of this work we will focus oulv upon the D11 aud D2 lines aad irespectivelv)., For the rest of this work we will focus only upon the 1 and D2 lines and respectively).
+ Pipeline optimal extractions were supplied. aud were of good quality in this region: a manual reduction of selected. spectra viclded uceligible inrprovenmoenut., Pipeline optimal extractions were supplied and were of good quality in this region; a manual reduction of selected spectra yielded negligible improvement.
+ We supplement these data with lower-esolution observations of the ΠΠ aud Is lines aand rrespectively)., We supplement these data with lower-resolution observations of the H and K lines and respectively).
+ These were obtained with the RGO spectrograph of the Anglo-Australian Telescope over 2002 June 61l. with the R1200D erating vielding a wavelength coverage ofAA.," These were obtained with the RGO spectrograph of the Anglo-Australian Telescope over 2002 June 6–11, with the R1200B grating yielding a wavelength coverage of."
+. The nuages were de-biased and flat-fielded. aud the spectrasubsequently extracted using conveutioual optimal extraction techniques (Eorne1986)..," The images were de-biased and flat-fielded, and the spectrasubsequently extracted using conventional optimal extraction techniques \citep{Horne:1986a}."
+ Wavelenetl calibrations were interpolated between Cur comparison lamp nuages. obtained every 30uunius.," Wavelength calibrations were interpolated between CuAr comparison lamp images, obtained every mins."
+ The kinematic resolution is estimated to be lkkiunss|., The kinematic resolution is estimated to be $^{-1}$.
+ We show Lue profiles of the two —ines in Fie., We show line profiles of the two lines in Fig.
+ d. relative to the local standard of rest (LSR).," \ref{NaDFig}
+ relative to the local standard of rest (LSR)."
+ Thev are clearly complex with a larec velocity Ispersion. sugecsting a relatively lavee distance.," They are clearly complex with a large velocity dispersion, suggesting a relatively large distance."
+ This inunediatelv rules out suggestions that lis a local object at — I.3kkpe (Mauche&Gorenstein 1986: Predehletal. 1991): the ine couclusion was rawn bv άλαςκαὶetal.(1998). for different reasons., This immediately rules out suggestions that is a local object at $\sim 1.3$ kpc \citealt{Mauche:1986a}; \citealt{Predehl:1991a}) ); the same conclusion was drawn by \citet{Zdziarski:1998a} for different reasons.
+ We cau obtain anu approximate cstimate of the istribution of wwith velocity by converting the line profiles iuto apparent optical depths (Fig. 2))," We can obtain an approximate estimate of the distribution of with velocity by converting the line profiles into apparent optical depths (Fig. \ref{NaDTauFig}) ),"
+ as discussed by Sembach.Danks.&Savage (1993).., as discussed by \citet{Sembach:1993a}. .
+ This is not the true, This is not the true
+smaller.,smaller.
+ There are also species that are not easy to interpret. such as NH and NH;.," There are also species that are not easy to interpret, such as $^+$ and $_2^+$."
+ The trends show a minimum or a maximum. when varying the CR rates.," The trends show a minimum or a maximum, when varying the CR rates."
+ À number of species we considered have a strong dependence on temperature and ionization rate., A number of species we considered have a strong dependence on temperature and ionization rate.
+ Especially CS exhibits are very complex interplay. and does not show a particular trend.," Especially CS exhibits are very complex interplay, and does not show a particular trend."
+ Molecular hydrogen. H». is a key species in the chemistry of the ISM.," Molecular hydrogen, $_2$, is a key species in the chemistry of the ISM."
+ The formation of H». which mainly occurs on grain surfaces at these densities and metallicity. depend on the specific properties and temperatures of the grains. and is subject to a considerable uncertainty (cf..Cazaux&Tielens 2004).," The formation of $_2$, which mainly occurs on grain surfaces at these densities and metallicity, depend on the specific properties and temperatures of the grains, and is subject to a considerable uncertainty \citep[cf.,][]{Cazaux2004}."
+. As a result the H» abundance is uncertat as well. especially when the gas is not predominantly molecular. 1.e.. in the radical region.," As a result the $_2$ abundance is uncertain as well, especially when the gas is not predominantly molecular, i.e., in the radical region."
+ Therefore. the column densities of key species such as O-. OH. OH. and H»O' might be affected. especially when the integrated column density is dominated by the abundances in the radical region.," Therefore, the column densities of key species such as $^+$, OH, $^+$, and $_2$ $^+$ might be affected, especially when the integrated column density is dominated by the abundances in the radical region."
+ In order to investigate this. we lowered the H» formation rate by a factor 3. and here we summarize the results for some of the species. considering Model set la::: The abundance of E» is lower by a factor of ~34 in the unshielded region of the cloud. and the H to H» transition occurs now more gradual.," In order to investigate this, we lowered the $_2$ formation rate by a factor 3, and here we summarize the results for some of the species, considering Model set 1a: The abundance of $H_2$ is lower by a factor of $\sim 3$ in the unshielded region of the cloud, and the H to $_2$ transition occurs now more gradual."
+" In the molecular region. there is only a noticeable effect when there is no full transition to H+. which ts for the highest considered CR rate. é=35-1074 κ,"," In the molecular region, there is only a noticeable effect when there is no full transition to $_2$, which is for the highest considered CR rate, $\zeta = 5\cdot 10^{-13}$ $^{-1}$."
+ The column density of atomic hydrogen is ~10 percent higher for the lowest up to ~80 percent higher for the highest CR rate., The column density of atomic hydrogen is $\sim10$ percent higher for the lowest up to $\sim80$ percent higher for the highest CR rate.
+ Molecular hydrogen column densities only change significantly for the highest two CR rates. ~20 and ~40 percent lower. respectively.," Molecular hydrogen column densities only change significantly for the highest two CR rates, $\sim20$ and $\sim40$ percent lower, respectively."
+ The H column density is lower by a factor ~3 for the lowest and ~5 to ~15 percent higher for the highest two CR rates., The $^+$ column density is lower by a factor $\sim3$ for the lowest and $\sim 5$ to $\sim 15$ percent higher for the highest two CR rates.
+ The H column density is dominated by Hy + H > H™ + H> in the radical region at £=5-10 s! and thus lower. while at higher CR rates the column density is increased because of the only partial transition to H»H;O:: The trend seen for the O column density ts very close to what is seen for H™. i.e. a factor of ~3 lower for z=5+I0 s7!. while slightly higher for the highest CR rates.," The $^+$ column density is dominated by $_2^+$ + H $\rightarrow$ $^+$ + $_2$ in the radical region at $\zeta=5\cdot10^{-17}$ $^{-1}$ and thus lower, while at higher CR rates the column density is increased because of the only partial transition to $_2$: The trend seen for the $^+$ column density is very close to what is seen for $^+$, i.e, a factor of $\sim3$ lower for $\zeta=5 \cdot10^{-17}$ $^{-1}$ , while slightly higher for the highest CR rates."
+ Similar trends are seen for OH and H»O'., Similar trends are seen for $^+$ and $_2$ $^+$.
+ HO is only significantly affected (lowered by a factor ~ 2) at the highest CR rate., $_2$ O is only significantly affected (lowered by a factor $\sim 2$ ) at the highest CR rate.
+ OH column density is lower by a factor -3 for the lowest CR rates. and dominated by the abundance in the unshielded region of the cloud.," OH column density is lower by a factor $\sim3$ for the lowest CR rates, and dominated by the abundance in the unshielded region of the cloud."
+ This difference is much less (5-30 percent) for the higher CR rates., This difference is much less (5-30 percent) for the higher CR rates.
+ H:O and H307 ,$_2$ O and $_3$ $^+$ 
+Submillimetre Galaxies (SALGs) contribute significantly to the rapid buildup of stellar mass in the Universe at 2~ 2.,"Submillimetre Galaxies (SMGs) contribute significantly to the rapid buildup of stellar mass in the Universe at $z\,\sim\,$ 2."
+ Llowever. their selection at iis inherently biased towards: colcder-cdust sources. (alesetal.2000:Blain 2004)..," However, their selection at is inherently biased towards colder-dust sources \citep{eales00a,blain04a}. ."
+ Recent work (e.g.Chap-manetal.2004:Casey200960). has demonstrated that S5Ofin--aint. high-redshift Ultraluminous Infrared Calaxies (ULIRGs) exist. and mav contribute significantly to the cosmic star formation rate density at its peak.," Recent work \citep[e.g.][]{chapman04a,casey09b} has demonstrated that -faint, high-redshift Ultraluminous Infrared Galaxies (ULIRGs) exist and may contribute significantly to the cosmic star formation rate density at its peak."
+ (2009c) describe a population of Iliuminous galaxies at 1.5 whose infrared. luminosities exceed ~ 1077 bbut are S5Ofun--faint due to hotter. characteristic dust temperatures., \citet{casey09b} describe a population of luminous galaxies at $z\sim$ 1.5 whose infrared luminosities exceed $\sim$ $^{12}$ but are -faint due to hotter characteristic dust temperatures.
+ Sparse infrared. data. particularly in the wwaveleneth. range. along with poor volume density constraints have limited the interpretation of these subnim-faint ULIRGs.," Sparse infrared data, particularly in the wavelength range, along with poor volume density constraints have limited the interpretation of these submm-faint ULIRGs."
+ Similar studies of other infrarecd-DIuminous galaxy populations. selected at.24/nm.. oor l.2mm for example. present even more evidence for diverse populations of luminous. dusty starbursts at 11 which clo notnecessarily intersect (seeDev 2009)..," Similar studies of other infrared-luminous galaxy populations, selected at, or 1.2mm for example, present even more evidence for diverse populations of luminous, dusty starbursts at 1 which do notnecessarily intersect \citep[see][]{dey08a,bussmann09a,younger09a}. ."
+"remains lower than in non-BCGs, especially at 5<reg20 kpc.","remains lower than in non-BCGs, especially at $5<\rm r_{eff}<20$ kpc."
+" To better understand the non-monotonic variation of unm—1 with reg(r), and remove possible selection effects (from the flux limit of the sample) we re-examine this relation with the E/S0s divided by M; (Figure 12)."," To better understand the non-monotonic variation of ${{{\rm r}_{eff}(g)}\over{{\rm r}_{eff}(r)}}-1$ with $\rm r_{eff}(r)$, and remove possible selection effects (from the flux limit of the sample) we re-examine this relation with the E/S0s divided by $M_r$ (Figure 12)."
+ For moderate luminosities there is simply a steep increase in colour gradient with radius., For moderate luminosities there is simply a steep increase in colour gradient with radius.
+" The turnover at large radii is produced by M.«—23 galaxies, in which colour gradient is ‘suppressed’ and the correlation with radius is flattened."," The turnover at large radii is produced by $M_r<-23$ galaxies, in which colour gradient is `suppressed' and the correlation with radius is flattened."
+" 'This effect is also seen, even more strongly, in the BCGs."," This effect is also seen, even more strongly, in the BCGs."
+" In Paper I, by comparing model and aperture colours, we found evidence that stronger colour gradients were associated with E/SO galaxies with high luminosity relative to c."," In Paper I, by comparing model and aperture colours, we found evidence that stronger colour gradients were associated with E/S0 galaxies with high luminosity relative to $\sigma$."
+" Figure 13 plots o—1 against 10logo+M; (approximately the c residual relative to the c—M,. relation; Forbes and Ponman 1999).", Figure 13 plots ${{{\rm r}_{eff}(g)}\over{{\rm r}_{eff}(r)}}-1$ against $10~{\rm log}~\sigma+M_r$ (approximately the $\sigma$ residual relative to the $\sigma$ $M_r$ relation; Forbes and Ponman 1999).
+ In the full E/SO sample the colour gradient increases by factor of two from the highest to the lowest ratios of o to aluminosity., In the full E/S0 sample the colour gradient increases by a factor of two from the highest to the lowest ratios of $\sigma$ to luminosity.
+" In the M,«—22.5 non-BCG sample this increase is an even larger factor of three.", In the $M_r<-22.5$ non-BCG sample this increase is an even larger factor of three.
+" In marked contrast, the colour gradient in BCGs shows little or no dependence on the c residual and remains at ~0.08."," In marked contrast, the colour gradient in BCGs shows little or no dependence on the $\sigma$ residual and remains at $\simeq 0.08$."
+" Figure 14 shows color gradient against o7/r2g, a more physical quantity as it is a direct measure of mass density (subject to a mild dependence on long-axis orientation)."," Figure 14 shows color gradient against $\sigma^2/\rm r_{eff}^2$, a more physical quantity as it is a direct measure of mass density (subject to a mild dependence on long-axis orientation)."
+ We express this in the form log (o?/r2q) for c in km s! and reg in kpc., We express this in the form log $\sigma^2/\rm r_{eff}^2$ ) for $\sigma$ in km $\rm s^{-1}$ and $\rm r_{eff}$ in kpc.
+ In these units the mean density for the E/SOs is 3.38 with a scatter 0.42., In these units the mean density for the E/S0s is 3.38 with a scatter 0.42.
+ This can also be converted to units of solar masses Mo., This can also be converted to units of solar masses $M_{\odot}$ .
+" From Boylan-Kolchin, Ma and Quataert (2006) the mass within r<reg is approximately 2.90reMz!GE, where the gravitational constant G—A30x103 pc km s!)?."," From Boylan-Kolchin, Ma and Quataert (2006) the mass within $\rm r<r_{eff}$ is approximately $2.90 {{\sigma^2 \rm r_{eff}}\over{G}}$, where the gravitational constant $G=4.30\times10^{-3}$ pc $\rm M_{\odot}^{-1}$ km $\rm s^{-1})^2$."
+" Within Το the volumeis 25125, and the mean mass density, log(Mo/kpc?)=(c2,2reg)+5.21."," Within $r_{eff}$ the volume is ${4\over 3}\pi \rm r_{eff}^3$ , and the mean mass density, ${\rm log}~(\rm M_{\odot}/kpc^3)={\rm log}~(\sigma^2/\rm r_{eff}^2)+5.21$."
+" The full E/S0 sample and the M,<—22.5 non-BCGs show a similar decrease in colour gradient with increasing mass density, whereas for BCGs πω—1~0.08 over the whole density range."," The full E/S0 sample and the $M_r<-22.5$ non-BCGs show a similar decrease in colour gradient with increasing mass density, whereas for BCGs ${{{\rm r}_{eff}(g)}\over{{\rm r}_{eff}(r)}}-1\simeq 0.08$ over the whole density range."
+ Figure 15 showscolour gradient against density, Figure 15 showscolour gradient against density
+The ο spectrum (Jones et 22004) shows a bright featureless continuum ancl no redshift’ could be measured.,The 6dFGS spectrum (Jones et 2004) shows a bright featureless continuum and no redshift could be measured.
+ The optical LD is a 2dEGIUS (Colless et 22001) galaxy at 220.154 whose spectrum shows strong. narrow emission lines.," The optical ID is a 2dFGRS (Colless et 2001) galaxy at $z$ =0.154 whose spectrum shows strong, narrow emission lines."
+ This is one of the fe AT20€1 SOUPCOS in our sample which was uncdetected at GCGllz., This is one of the few AT20G sources in our sample which was undetected at GHz.
+ The racio spectrum rises steeply [rom 0.8 to 20€ClIz. and so must peak somewhere above (αν.," The radio spectrum rises steeply from 0.8 to GHz, and so must peak somewhere above GHz."
+ The MMIIZ SUMSS images shows extended emission with a deconvolved largest angular size of aarcsec at PA 367., The MHZ SUMSS images shows extended emission with a deconvolved largest angular size of arcsec at PA $^\circ$.
+ This appears to be the same object as 4044. though the cataloguec PALIN) position is significantly offset. from the more accurate AT20€ raclio xosition.," This appears to be the same object as $-$ 4044, though the catalogued PMN position is significantly offset from the more accurate AT20G radio position."
+" ""There is a faint (stellar) optical LD at the AT2OC »osition.", There is a faint (stellar) optical ID at the AT20G position.
+ This source appears extended in he A\IAUI2 SUAISS image (ceconvolvecl major axis aaresee at. PA 257)., This source appears extended in the MHz SUMSS image (deconvolved major axis arcsec at PA $^\circ$ ).
+ Phe 6dECS spectrum shows strong. broad Balmer emission lines tvpical of a Sevfert 1 galaxy or QSO.," The 6dFGS spectrum shows strong, broad Balmer emission lines typical of a Seyfert 1 galaxy or QSO."
+ This source is so close to a bright star hat no optical is LD possible., This source is so close to a bright star that no optical is ID possible.
+ Dlank field on the blue DSS image., Blank field on the blue DSS image.
+ No known optical LD., No known optical ID.
+ One of the few sources in our sample which shows an inverted radio spectrum at all frequencies rom OS to 95CGCGllz., One of the few sources in our sample which shows an inverted radio spectrum at all frequencies from 0.8 to GHz.
+ The DSS optical ID is a 19th magnitude galaxy whose redshift is currently unknown., The DSS optical ID is a 19th magnitude galaxy whose redshift is currently unknown.
+ This source is well below the limit. of he SUMSS catalogue at MMlIE. but examination of the original SUMSSimage shows a marginal (3m) ," This source is well below the limit of the SUMSS catalogue at MHz, but examination of the original SUMSS image shows a marginal $\sigma$ "
+scattered bright regions iu the optical.,scattered bright regions in the optical.
+ In these eas-rich galaxies show a complicated fractal pattern full of shells aud clumps., In these gas-rich galaxies show a complicated fractal pattern full of shells and clumps.
+ In the more massive dlirs the distribution is usually much more extended {μα even the oldest stellar populatious., In the more massive dIrrs the distribution is usually much more extended than even the oldest stellar populations.
+ In. some cases evidence [or gas accretion lias been observed (e.g.. Wilcots filler 1998).," In some cases evidence for gas accretion has been observed (e.g., Wilcots Miller 1998)."
+ In less maassive cllrrs the may be displaced with respect to the stellar liglit. distribujon. either off-centered or showing a ring-like structure with a central depression (e.g.. Young> Lo ↽19974).," In less massive dIrrs the may be displaced with respect to the stellar light distribution, either off-centered or showing a ring-like structure with a central depression (e.g., Young Lo 1997a)."
+"- DIrrs are characterized⋅ by gaSEo.23 mag 27. Alyy109AZ... and Adi,<101AL.."," DIrrs are characterized by $\mu_V \la 23$ mag $^{-2}$, $M_{\rm HI} \la 10^9 M_{\odot}$, and $M_{\rm tot} \la 10^{10} M_{\odot}$."
+ The more massive dLrrs 5ow Increasing chemical enrichment with decreasing age., The more massive dIrrs show increasing chemical enrichment with decreasing age.
+ Iu low-mass «]rrs gas aud stars may show distinct spatial distributious aud kinematics., In low-mass dIrrs gas and stars may show distinct spatial distributions and kinematics.
+ Solid-body rotation is comaunon aniong the more massive ος. while low-1nass αιW. occasionally do not show measturale rotation.," Solid-body rotation is common among the more massive dIrrs, while low-mass dIrrs occasionally do not show measurable rotation."
+ DIrrs are founliu clusters. groups. aid in the field.," DIrrs are found in clusters, groups, and in the field."
+ They exhibit little coucentration toward massive galaxies., They exhibit little concentration toward massive galaxies.
+ Massive dlrrs may continue to [orm stars over a Hubble time (Hunter 1997)., Massive dIrrs may continue to form stars over a Hubble time (Hunter 1997).
+ Very low-mass. gas-poor dirrs. on the other hand. may bT evolving towards dwarf spheroidals. (," Very low-mass, gas-poor dIrrs, on the other hand, may be evolving towards dwarf spheroidals. ("
+dEs) are spherical or elliptical in appearance. compact. have hie[un central stellar densities and are typically fainter than A:=-l7Tn jag.,"dEs) are spherical or elliptical in appearance, compact, have high central stellar densities and are typically fainter than $M_V = -17$ mag."
+ Other characteristics iuclude pyS+dyPl mae 27. MgSTas105AL. and MjXa“:109AL..," Other characteristics include $\mu_V \la 21$ mag $^{-2}$, $M_{\rm HI}
+\la 10^8 M_{\odot}$, and $M_{tot} \la 10^9 M_{\odot}$."
+ They teud to be found close to massive ealaxies. usually have little or uo detectable eas. aud are ofteu not rotationally supported.," They tend to be found close to massive galaxies, usually have little or no detectable gas, and are often not rotationally supported."
+ If. eas is detected it may slow au asyuiuetric distribution. be less extended thau the underlying stellar elit. aud may be kinematically distinct (Young Lo 1997b: Sage. Welch. Mitchell 1998: Welch. Sage.e Mitchell 1998).," If gas is detected it may show an asymmetric distribution, be less extended than the underlying stellar light, and may be kinematically distinct (Young Lo 1997b; Sage, Welch, Mitchell 1998; Welch, Sage, Mitchell 1998)."
+ DEs may show pronounced nuclei that can contribute up to of the galaxy's light., DEs may show pronounced nuclei that can contribute up to of the galaxy's light.
+ The fraction of these nucleated dEs (dECN)) increases among more Iuminous dEs., The fraction of these nucleated dEs (dE(N)) increases among more luminous dEs.
+ The surface density profiles of dEs aud dE(N)s are best described by Sérrsic’s (1968) generalization of a cle Vaucouleurs rt law and exponential profiles (Jerjen et 2000). (, The surface density profiles of dEs and dE(N)s are best described by Sérrsic's (1968) generalization of a de Vaucouleurs $r^{1/4}$ law and exponential profiles (Jerjen et 2000). (
+dSplis) are cdilluse. gas-deficieut. low-surlace-briglituess dwarls with very lite ceutral concentration.,"dSphs) are diffuse, gas-deficient, low-surface-brightness dwarfs with very little central concentration."
+" They are characterized by Adyο—11 mag. pyz22 mag >7. AL;""EEAl. and Mu~10*A..."," They are characterized by $M_V \ga -14$ mag, $\mu_V \ga 22$ mag $^{-2}$, $M_{\rm HI} \la 10^5 M_{\odot}$, and $M_{tot} \sim 10^7 M_{\odot}$."
+ DSphis are the faintest. least massive galaxies known.," DSphs are the faintest, least massive galaxies known."
+ The upper liits fo| their are usually even below the amouuts expected from uormal stellar tnass loss., The upper limits for their are usually even below the amounts expected from normal stellar mass loss.
+ DSplis aο usually louucd iu close proximity of massive galaxies aud do not appear to be supporte by roation., DSphs are usually found in close proximity of massive galaxies and do not appear to be supported by rotation.
+ Their velocity dispersious indicate the presence of a significant dark colnponelnt wjen vlrial equilibrium is assure., Their velocity dispersions indicate the presence of a significant dark component when virial equilibrium is assumed.
+ The cdille'ence 1 the spatial distribution of dEs aud dSphs on the one hand and ο» on the other hand is also kuown as or as the morphology—-deusity relation for cdwar galaxies aud way be a signature of euvironmeutal effects oi galaxy evolution., The difference in the spatial distribution of dEs and dSphs on the one hand and dIrrs on the other hand is also known as or as the morphology–density relation for dwarf galaxies and may be a signature of environmental effects on galaxy evolution.
+ form [rom the debris torn out of lore Massive galaxies curing iuteractious auc inergers., form from the debris torn out of more massive galaxies during interactions and mergers.
+ Iu conu‘ast to other cwarf galaxies they do 1kot contal1 cark matter and may lave hiel metallicities or their luminosity depeucing on the evolutlonary stage of the parent galaxy (Duc Mirabel 1995)., In contrast to other dwarf galaxies they do not contain dark matter and may have high metallicities for their luminosity depending on the evolutionary stage of the parent galaxy (Duc Mirabel 1998).
+ Masses. sizes. gas content. etc.," Masses, sizes, gas content, etc."
+ of tidal dwarls deyencl critically on the couditious duriug the ineractiou aud on the parent galaxy., of tidal dwarfs depend critically on the conditions during the interaction and on the parent galaxy.
+ It is possible that certain cllrrs aud dSplis are tida cdwarfs that [oriued. relatively early., It is possible that certain dIrrs and dSphs are tidal dwarfs that formed relatively early.
+ For instance. various auhors suggested that the Milky Way dwarf compaious may be part of two or more major debris streatis in polar orbits (e.g.. Majewski 199L: Lvuden-Bell Lyudeu-Bell 1995).," For instance, various authors suggested that the Milky Way dwarf companions may be part of two or more major debris streams in polar orbits (e.g., Majewski 1994; Lynden-Bell Lynden-Bell 1995)."
+ Potential caucliclates for 1rearby tidal dwarf galaxies base on their intriiie properties are discussed by Hunter et ((2000)., Potential candidates for nearby tidal dwarf galaxies based on their intrinsic properties are discussed by Hunter et (2000).
+ The dwarl galaxies studied most exteusively to date are cdwart galaxies within the Local Group., The dwarf galaxies studied most extensively to date are dwarf galaxies within the Local Group.
+ Of the above mentioned types. the Local Croup contains dlirs. «Es. aud dSplis as well as galaxies that may be in transition from cdlIrrs to dSphs.," Of the above mentioned types, the Local Group contains dIrrs, dEs, and dSphs as well as galaxies that may be in transition from dIrrs to dSphs."
+are: P!X1x1077 dyne em. P~6.5x107? dyne em. and prs~2.5%107'? dyne em.,"are: $P^{t}\la 1\times10^{-9}$ dyne $^{-2}$ , $P^{r}\sim 6.5\times10^{-10}$ dyne $^{-2}$, and $P^{ms}\sim 2.5\times10^{-10}$ dyne $^{-2}$."
+ The thermal conduction between the cloud and the intercloud medium drives a thin hot halo around the cloud., The thermal conduction between the cloud and the intercloud medium drives a thin hot halo around the cloud.
+ As the primary shock progressively engulfs the cloud. it generates lateral shocks propagating into the cloud.," As the primary shock progressively engulfs the cloud, it generates lateral shocks propagating into the cloud."
+ After ~3000 vr. in the regions where the transmitted shock interacts with the lateral shocks. there are dense (7~5 em?) cold (P.~1—2xIO? Κ). and elongated (L~5xI0 em) structures.," After $\sim3000$ yr, in the regions where the transmitted shock interacts with the lateral shocks, there are dense $n\sim 5$ $^{-3}$ ), cold $T\sim1-2\times 10^{5}$ K), and elongated $L\sim 5\times 10^{17}$ cm) structures."
+ These structures are larger than the critical. Field length and are therefore thermally instable., These structures are larger than the critical Field length and are therefore thermally instable.
+ They cool off to temperatures of a few 10 K. reaching density 2~10 em™ in a few 10° yr (and are visible in Fig. 4)).," They cool off to temperatures of a few $10^{4}$ K, reaching density $n\sim 10$ $^{-3}$ in a few $10^{2}$ yr (and are visible in Fig. \ref{fig:nTAB}) )."
+ These values are consistent with the characteristic temperature and density of optical filaments in SNRs and with those observed in the FilD region (?))., These values are consistent with the characteristic temperature and density of optical filaments in SNRs and with those observed in the FilD region \citealt{bms00}) ).
+ We caution the reader that our simulations nodel a plasma where hydrogen is fully ionized and that for very low values of temperature our radiative losses function Is not calculated accurately., We caution the reader that our simulations model a plasma where hydrogen is fully ionized and that for very low values of temperature our radiative losses function is not calculated accurately.
+ Therefore the values of temperature and density derived when 7 approaches the H ionization temperature may not be completely correct., Therefore the values of temperature and density derived when $T$ approaches the H ionization temperature may not be completely correct.
+ However. our nodel clearly indicates that thermal instabilities can. produce the FilD optical emission.," However, our model clearly indicates that thermal instabilities can produce the FilD optical emission."
+ The upper panels in Fig., The upper panels in Fig.
+ 5. show the synthesized X-ray naps. for setup Sphl. at the same epoch of Fig. 4..," \ref{fig:mappeXAB} show the synthesized X-ray maps, for setup Sph1, at the same epoch of Fig. \ref{fig:nTAB},"
+ derived in the three different energy bands. 0.3—0.5 keV. 0.5—] keV. and 0.3—2 keV. by following the procedure described in Sect. 2.4..," derived in the three different energy bands, $0.3-0.5$ keV, $0.5-1$ keV, and $0.3-2$ keV, by following the procedure described in Sect. \ref{Xsynth}."
+ The synthesized count rate per space bin is in agreemert with the observed one. considering its strong dependence oa the source distance D (the synthesized count rate scales as D7. because the flux seales as D and the angular extension of the bin as D!x D!) and that the distance of the Vela SNR ts known with ~30% uncertainty.," The synthesized count rate per space bin is in agreement with the observed one, considering its strong dependence on the source distance $D$ (the synthesized count rate scales as $D^{-4}$ , because the flux scales as $D^{-2}$ and the angular extension of the bin as $D^{-1}\times D^{-1}$ ) and that the distance of the Vela SNR is known with $\sim30\%$ uncertainty."
+ Although the overall morphology of the synthesized emission in the 0.3—2 keV band for setup Sphl appears similar to the observed one. there are significant differences.," Although the overall morphology of the synthesized emission in the $0.3-2$ keV band for setup Sph1 appears similar to the observed one, there are significant differences."
+ Great part of the synthesized emission above 0.5 keV. shown in the central upper panel of Fig. 5..," Great part of the synthesized emission above 0.5 keV, shown in the central upper panel of Fig. \ref{fig:mappeXAB},"
+ is clearly associated to the plasma heated by the reflected shock (see Fig. 4)).," is clearly associated to the plasma heated by the reflected shock (see Fig. \ref{fig:nTAB}) ),"
+ while the bulk of the soft X-ray emission in the 0.3—0.5 keV band originates in the cloud plasma heated by the transmitted shock., while the bulk of the soft X-ray emission in the $0.3-0.5$ keV band originates in the cloud plasma heated by the transmitted shock.
+ The morphology of the synthesized emission below 0.5 keV is then very different from the synthesized emission in the 0.5—| keV band., The morphology of the synthesized emission below $0.5$ keV is then very different from the synthesized emission in the $0.5-1$ keV band.
+ In setup Sphl. therefore. the morphology of the X-ray emitting knot in the broad band 0.3—2 keV is the result of the superposition of the emission associated with two different physical conditions of the plasma: the cloud material heated by the transmitted shock at North. which contributes with soft emission. and the intercloud material heated by the reflected shock in the central and southern parts of the knot. with harder emission.," In setup Sph1, therefore, the morphology of the X-ray emitting knot in the broad band $0.3-2$ keV is the result of the superposition of the emission associated with two different physical conditions of the plasma: the cloud material heated by the transmitted shock at North, which contributes with soft emission, and the intercloud material heated by the reflected shock in the central and southern parts of the knot, with harder emission."
+ The extension of the X-ray emitting knot in the 0.3—2 keV band is significantly larger than thatin the 0.3—0.5 keV band., The extension of the X-ray emitting knot in the $0.3-2$ keV band is significantly larger than thatin the $0.3-0.5$ keV band.
+ This result does not match the observation of the ΕΠΟ knot. which shows that the X-ray emission in the 0.3—0.5 keV band has almost the same extension as the ones in the 0.5—| keV and 0.3—2 keV bands (Fig. 1)).," This result does not match the observation of the FilD knot, which shows that the X-ray emission in the $0.3-0.5$ keV band has almost the same extension as the ones in the $0.5-1$ keV and $0.3-2$ keV bands (Fig. \ref{fig:XMM}) )."
+ The upper panel of Fig., The upper panel of Fig.
+ 6 shows. for setup Sphl. the emission measure of the plasma with 7«10? K and n>3 em projected in the plane of the sky. which Is a proxy of the optical emission.," \ref{fig:X-HalphaS} shows, for setup Sph1, the emission measure of the plasma with $T<10^{5}$ K and $n>3$ $^{-3}$ projected in the plane of the sky, which is a proxy of the optical emission."
+ We also superimpose. in green. thecorresponding X-ray contour levels at a count rate 2.4 times higher than the minimum. derived from the synthesized X-ray imagesin the 0.3—2 keV band shown in Fig. 5..," We also superimpose, in green, thecorresponding X-ray contour levels at a count rate 2.4 times higher than the minimum, derived from the synthesized X-ray imagesin the $0.3-2$ keV band shown in Fig. \ref{fig:mappeXAB}. ."
+ The impact of the Vela shock front with the spherical cloud of setup Sphl, The impact of the Vela shock front with the spherical cloud of setup Sph1
+where f(s) is the new function that we are proposing in order to improve the fit to the external parts of DM haloes.,where $f(s)$ is the new function that we are proposing in order to improve the fit to the external parts of DM haloes.
+ We choose the following rational function: 'There are several reasons for this choice., We choose the following rational function: There are several reasons for this choice.
+" On the one hand, we aim for a simple formula which allows an easy calculation for the tangential shear."," On the one hand, we aim for a simple formula which allows an easy calculation for the tangential shear."
+" On the other hand, we want to minimize the influence of the new added parameters (b; and 0ο) on the internal profile, where the NFW profile is steeper, except for the very inner regions in which NFW and f(s) are both proportional to r-!."," On the other hand, we want to minimize the influence of the new added parameters $b_1$ and $b_2$ ) on the internal profile, where the NFW profile is steeper, except for the very inner regions in which NFW and $f(s)$ are both proportional to $r^{-1}$."
+" We thus propose in this work the following model as an approximation of the density profile from the inner regions up to large distances well beyond Ri: It is worth noting that the enclosed mass and the characteristic density are now given by these relations: Applying the condition (4)) to the expression (13)) leads to Eq. (14)),"," We thus propose in this work the following model as an approximation of the density profile from the inner regions up to large distances well beyond $R_{\rmn{vir}}$: It is worth noting that the enclosed mass and the characteristic density are now given by these relations: Applying the condition \ref{eq:mv1}) ) to the expression \ref{eq:masa}) ) leads to Eq. \ref{eq:rhos}) ),"
+ which quantifies the influence of external parameters bi and 69 on the value of ps., which quantifies the influence of external parameters $b_1$ and $b_2$ on the value of $\rho_s$.
+" In the case of b;=bz 0, which is equivalent to f(s)=0, we recover the modified NFW characteristic density p™°*."," In the case of $b_1=b_2=0$ , which is equivalent to $f(s)=0$, we recover the modified NFW characteristic density $\rho_s^{\rmn{mod}}$."
+ An example of this new fitting formula for Milky Way size haloes is shown in Figure 1., An example of this new fitting formula for Milky Way size haloes is shown in Figure 1.
+ Here we compare different approximations to fit the numerical density profile., Here we compare different approximations to fit the numerical density profile.
+" In order to reduce the range of variation along the y-axis, we plot the function (p/p)-(r/Ry)? instead of just p/p."," In order to reduce the range of variation along the $y$ -axis, we plot the function $ \left(\rho/\bar{\rho}\right) \cdot (r/R_{\rmn{vir}})^{2} $ instead of just $ \rho/\bar{\rho} $."
+ The relative errors for each one of them are also displayed., The relative errors for each one of them are also displayed.
+" The numerical profile given by cosmological simulations (dotted line) is compared to theNFW model (Eq. 2)),"," The numerical profile given by cosmological simulations (dotted line) is compared to theNFW model (Eq. \ref{eq:nfw}) ),"
+" the modified NFW model (NFW plus the mean matter density of the Universe, Eq. 8))"," the modified NFW model (NFW plus the mean matter density of the Universe, Eq. \ref{eq:nfwmod}) )"
+ and our model (Eq. 12))., and our model (Eq. \ref{eq:ourmodel}) ).
+ The fit in the inner part is quite similar for all the models as expected., The fit in the inner part is quite similar for all the models as expected.
+" There is an obvious discrepancy beyond the virial radius between the NFW fit, the modified NFW approach, and the data from numerical simulations."," There is an obvious discrepancy beyond the virial radius between the NFW fit, the modified NFW approach, and the data from numerical simulations."
+" Our model helps to alleviate this disagreement, providing a reasonable fit going from the inner regions to well beyond the virial radius up to 10Rvir."," Our model helps to alleviate this disagreement, providing a reasonable fit going from the inner regions to well beyond the virial radius up to $10R_{\rmn{vir}}$."
+ 'The description of the DM distribution from small to large distances from the centre of a halo presents an added relevance as it is related to the halo-dark matter correlation function €galo-pbM=(ptr-p (e.g. ?))., The description of the DM distribution from small to large distances from the centre of a halo presents an added relevance as it is related to the halo-dark matter correlation function $\xi_{\rmn {Halo-DM}}= \frac{\left\langle\rho(r)\right\rangle -\bar{\rho}}{\bar{\rho}} $ (e.g. \citealt{Ha08}) ).
+" Therefore, our approximation for the density profiles in the outer regions of DM haloes is also showing information about this correlation function."," Therefore, our approximation for the density profiles in the outer regions of DM haloes is also showing information about this correlation function."
+ In Figure 2 we display €Haio_pm for different bins in the mass range covered in this paper (see Section 3)., In Figure 2 we display $\xi_{\rmn {Halo-DM}}$ for different bins in the mass range covered in this paper (see Section 3).
+" In this plot we can clearly distinguish between the one-halo term, which corresponds to the halo density profile, and the two-halo term, which corresponds to the distribution of matter in its outskirts."," In this plot we can clearly distinguish between the one-halo term, which corresponds to the halo density profile, and the two-halo term, which corresponds to the distribution of matter in its outskirts."
+ Our approximation provides a reasonable description in both regions., Our approximation provides a reasonable description in both regions.
+ To carry out our analysis we have used several high-resolution cosmological simulations., To carry out our analysis we have used several high-resolution cosmological simulations.
+" The initial conditions of these simulations have been set up from the power spectrum corresponding to the ACDM model with the following values of the cosmological parameters: Qm= 0.3, Qa= 0.7, h = 0.7 and og= 0.9."," The initial conditions of these simulations have been set up from the power spectrum corresponding to the $\Lambda$ CDM model with the following values of the cosmological parameters: $\Omega _{m} $ = 0.3, $\Omega _{\Lambda} $ = 0.7, $h$ = 0.7 and $\sigma _{8} $ = 0.9."
+" We used three different boxes whose size is 80, 120 and 250 h! Mpc."," We used three different boxes whose size is 80, 120 and 250 $h^{-1}$ Mpc."
+" The mass resolution in these boxes is 3.14x105, 1.06x10? and 9.59x10?7! respectively, and the space resolution is 1.2, 1.8 and 7.6 A!Mc kpc, respectively."," The mass resolution in these boxes is $3.14 \times 10^8$, $1.06 \times 10^9$ and $9.59 \times 10^9 h^{-1} M_{\odot} $ respectively, and the space resolution is 1.2, 1.8 and 7.6 $h^{-1}$ kpc, respectively."
+" The final output of the simulation is obtained by tracking the evolution of 512? particles in each box to z—0, using the N-body code ART (?).."," The final output of the simulation is obtained by tracking the evolution of $512^3$ particles in each box to $z = 0$, using the $N$ -body code ART \citep{Kr97}."
+ The collapsed DM haloes are detected using Bound Density Maxima halofinder (BDM)., The collapsed DM haloes are detected using Bound Density Maxima halofinder (BDM).
+ We use mean halo density profiles averaged over seven different virial mass bins ranging from 5.35x10!A!Mc to 4.47x1045-7!Mc., We use mean halo density profiles averaged over seven different virial mass bins ranging from $5.35\times10^{11} h^{-1}M_{\odot} $ to $4.47\times10^{14} h^{-1} M_{\odot} $.
+" There is a total of 17,220 haloes in all simulation boxes but we constrain this sample with the following selection criterion in order to select distinct haloes: the centre of all selected haloes must be further than the virial radius of every halo with higher mass."," There is a total of 17,220 haloes in all simulation boxes but we constrain this sample with the following selection criterion in order to select distinct haloes: the centre of all selected haloes must be further than the virial radius of every halo with higher mass."
+" This criterion reduces the number of haloes to 16,679 (see ? for details on this halo sample)."," This criterion reduces the number of haloes to 16,679 (see \citealt{Cu07} for details on this halo sample)."
+ Figure 3 shows the best fit provided by our approximation to some average densityprofiles from simulations., Figure 3 shows the best fit provided by our approximation to some average densityprofiles from simulations.
+ In this plot we can see that there are deviations of the order of at the position smax of the local maximum of p(r)- r?., In this plot we can see that there are deviations of the order of at the position $s_{\rmn{max}}$ of the local maximum of $\rho(r)\cdot r^2$ .
+ This deviation is inherited from the, This deviation is inherited from the
+2001: Molencdi 2002).,2001; Molendi 2002).
+ Phe detailed. properties of these armis are discussed in Paper lL. Outside the arms. the ambient cluster gas is strikinely isothermal at each radius.," The detailed properties of these arms are discussed in Paper I. Outside the arms, the ambient cluster gas is strikingly isothermal at each radius."
+ This suggests that either the gas remains relatively uncdisturbed by AGN uplift or that. conduction is ellicient. alone the azimuthal direction. as expected under action of the heat-flux buovaney instability (LIBI).," This suggests that either the gas remains relatively undisturbed by AGN uplift or that conduction is efficient along the azimuthal direction, as expected under action of the heat-flux buoyancy instability (HBI)."
+ Many cavities are seen in the inner radio cocoon (40 aresec or 3 kpc) including one located at 7 aresec (~0.5 kpe) from the central CIN in the counter-jet. direction., Many cavities are seen in the inner radio cocoon $\sim40$ arcsec or $\sim3$ kpc) including one located at $\sim7$ arcsec $\sim0.5$ kpc) from the central AGN in the counter-jet direction.
+ These cavities are associated with decreased: surface brightness. high temperatures. low thermal pressure. ancl tvpically appear filled: with raclio plasma in projection.," These cavities are associated with decreased surface brightness, high temperatures, low thermal pressure, and typically appear filled with radio plasma in projection."
+ Our pressure map confirms spectrally previous results on the presence of a circular shock front at à1850 aresee (14 kpe)., Our pressure map confirms spectrally previous results on the presence of a circular shock front at $r\sim180$ arcsec (14 kpc).
+ Based on the surface brightness jump. we measure an average Mach. number of Al=1.25 for the shock at roc150 arcsec. in good agreement with the analysis of Forman (2007).," Based on the surface brightness jump, we measure an average Mach number of $M=1.25$ for the shock at $r\sim180$ arcsec, in good agreement with the analysis of Forman (2007)."
+ Under the assumption of a linearly rising temperature profile. we also measure a temperature jump within the ~40 aresee thick shock. consistent with the expectations for a shock with Mach number AJ=1.25.," Under the assumption of a linearly rising temperature profile, we also measure a temperature jump within the $\sim40$ arcsec thick shock, consistent with the expectations for a shock with Mach number $M=1.25$."
+ The thickness and overall weakness of these shocks will make it clillicu to detect. temperature jumps. unambiguously: in hotter clusters like Perseus Fabian 2006: Graham 2008)., The thickness and overall weakness of these shocks will make it difficult to detect temperature jumps unambiguously in hotter clusters like Perseus Fabian 2006; Graham 2008).
+ Another shock. associated with the central racio cocoon. is observed at 40 aresec (~3 kpc).," Another shock, associated with the central radio cocoon, is observed at $\sim40$ arcsec $\sim3$ kpc)."
+ Because of the complexity of this region. we can only estimate the Alach number of this shock. using the observed. projected temperature jump across the front.," Because of the complexity of this region, we can only estimate the Mach number of this shock, using the observed projected temperature jump across the front."
+ This suggests a Mach number of Al1.2., This suggests a Mach number of $M\approxgt1.2$.
+ The relative distances ancl Mach numbers of these shocks suggest that AGN outbursts are common and occur every ~LO Myrs., The relative distances and Mach numbers of these shocks suggest that AGN outbursts are common and occur every $\sim10$ Myrs.
+ Our results show that a significant. mass of metals is lifted up to laree radius in the wake of the radio bubbles., Our results show that a significant mass of metals is lifted up to large radius in the wake of the radio bubbles.
+ At least LO5ΟΝΕ. of Fe is uplifted to a radius of 350400 aresec (2731 kpc) to the north of the AGN., At least $1.0\times10^6 {\rm M_\odot}$ of Fe is uplifted to a radius of $350-400$ arcsec (27–31 kpc) to the north of the AGN.
+ Given that a similar amount of Fe (1.5«10?M.) is present in the X-ray. bright arms (Simionescu 2008). à single generation of bubbles can. in principle. explain the observed Ie excess.," Given that a similar amount of Fe $1.5\times10^6 {\rm M_\odot}$ ) is present in the X-ray bright arms (Simionescu 2008), a single generation of bubbles can, in principle, explain the observed Fe excess."
+ IH is still uncertain how much metal falls towarels the central regions after this process occurs., It is still uncertain how much metal falls towards the central regions after this process occurs.
+ Fe uplift of this magnitude may explain the observed metallicity ridges observed. around. some cool core clusters., Fe uplift of this magnitude may explain the observed metallicity ridges observed around some cool core clusters.
+ Studies of such abundance features can further enhance our understanding of transport processes., Studies of such abundance features can further enhance our understanding of transport processes.
+ This will be the subject of. future work., This will be the subject of future work.
+ We thank WLR. Forman. €. Jones. and LE. Churazov for helpful comments and RAG. Morris for computational support.," We thank W.R. Forman, C. Jones, and E. Churazov for helpful comments and R.G. Morris for computational support."
+ N. Werner and A. Simionescu were supported. by the National Aeronautics and Space Administration through Chandra/Iinstein Postdoctoral Fellowship Aware) Number PES-90056 and PE9-O0070 issued by the Chandra: X-ray Observatory Center. which is operated by the Smithsonian Astrophysical Observatory for and on behalf of the National Aeronautics and. Space Xdministration uncer contract NASS-03060.," N. Werner and A. Simionescu were supported by the National Aeronautics and Space Administration through Chandra/Einstein Postdoctoral Fellowship Award Number PF8-90056 and PF9-00070 issued by the Chandra X-ray Observatory Center, which is operated by the Smithsonian Astrophysical Observatory for and on behalf of the National Aeronautics and Space Administration under contract NAS8-03060."
+" ""his work was supported in part by the US Department of Energy. under contract number DIS-AC'02-Π", This work was supported in part by the US Department of Energy under contract number DE-AC02-76SF00515.
+ All computational analysis was carried out using the KIPAC NOC compute cluster at Stanford University and the Stanford. Linear Accelerator Center (SLAC)., All computational analysis was carried out using the KIPAC XOC compute cluster at Stanford University and the Stanford Linear Accelerator Center (SLAC).
+ )isone.,is one.
+"T hetotalf ractionof gridcellswithimaginaryeigenvaluesi sabio tiguousο Jandzz0,, and at)."," The total fraction of grid cells with imaginary eigenvalues is about at and, and at."
+". Withdecreasingheight, thehistogramismorepeakedandthe dicowirbin peakmevesiaiderddatyabovethes (sraabtaty Op"," With decreasing height, the histogram is more peaked and the peak moves toward larger swirling strengths (smaller $\taus$ )."
+tic thertdeer plotsof Fi ," In all the other plots of Fig. \ref{fig:disdens},"
+"taus« )(i.e., tothelef tof theverticalline)aretakenintoaccount. ("," only strong swirls with < (i.e., to the left of the vertical line) are taken into account. ("
+"ο) The histograms of the inclination angle ι indicate that the swirls are preferentially (i.e., more than expected for an isotropic distribution) horizontal at and preferentially horizontal or vertical at),,whereasatzz )thedistributionisconsistentwithanisotropicdistribution(indicatedbytheg dottedline). (","c) The histograms of the inclination angle $\iota$ indicate that the swirls are preferentially (i.e., more than expected for an isotropic distribution) horizontal at and preferentially horizontal or vertical at, whereas at the distribution is consistent with an isotropic distribution (indicated by the gray dash-dotted line). ("
+"d) The gas pressure p in swirling regions is, on average, reduced to of the horizontal mean )inalldownflowregions(vv, 0) at the same height, consistent with a contribution by a dynamic pressure due to centrifugal forces.","d) The gas pressure $p$ in swirling regions is, on average, reduced to of the horizontal mean in all downflow regions ) at the same height, consistent with a contribution by a dynamic pressure due to centrifugal forces."
+ The gas pressure deficit does not depend much on height and inclination angle. (, The gas pressure deficit does not depend much on height and inclination angle. (
+"e) The swirling regions are significantly rarefied, the average value of )being84%.","e) The swirling regions are significantly rarefied, the average value of being."
+.T herare factiondoe snotde pendmuchonheightandinclination¢ (, The rarefaction does not depend much on height and inclination angle. (
+f) The mean of the relative temperature )is11.0 at all three considered heights.,f) The mean of the relative temperature is at all three considered heights.
+" With increasing depth, the distribution becomes broader and the maximum moves towards lower temperatures."," With increasing depth, the distribution becomes broader and the maximum moves towards lower temperatures."
+" Most of the vertically oriented swirls (dashed lines) are slightly cooler than the horizontal ones (dotted lines), consistent with their predominant location inside cool intergranular lanes. ("," Most of the vertically oriented swirls (dashed lines) are slightly cooler than the horizontal ones (dotted lines), consistent with their predominant location inside cool intergranular lanes. ("
+"g) In horizontal swirls (dotted lines), the median and the mean of the vertical velocity are near zero, in agreement with them being found at the edges of the granules.","g) In horizontal swirls (dotted lines), the median and the mean of the vertical velocity are near zero, in agreement with them being found at the edges of the granules."
+" In vertical swirls (dashed lines), the mean is roughly )atallheights. ("," In vertical swirls (dashed lines), the mean is roughly at all heights. ("
+"h) The mean of the horizontal velocity Jisapproximately33),, irrespectiveo f theheight.","h) The mean of the horizontal velocity is approximately, irrespective of the height."
+Notethatthisvelocityincludesthehorizontalpre , Note that this velocity includes the horizontal proper motion of the whole swirl.
+," Vertical swirls tend to have slightly smaller horizontal velocities, presumably because most of them (being at the center of intergranular lanes) are not subject to the bulk motion of the expanding granules."
+"In Fig. 5,,"," In Fig. \ref{fig:twodens},"
+" we present 2D histograms of the swirling strength and selected observables, taking into account all downflow regions near the optical surface."," we present 2D histograms of the swirling strength and selected observables, taking into account all downflow regions near the optical surface."
+ This restricts the analysis to the intergranular lanes., This restricts the analysis to the intergranular lanes.
+" The top panel shows that strong magnetic fields become increasingly rare with increasing swirling strength, there being no apparent correlation between the two."," The top panel shows that strong magnetic fields become increasingly rare with increasing swirling strength, there being no apparent correlation between the two."
+ This is also true for both polarities of the vertical magnetic field alone., This is also true for both polarities of the vertical magnetic field alone.
+" As is to be expected for the downflow lanes, the emerging intensity in vortex regions is preferentially smaller than the overall mean (bottom panel of Fig. 5))."," As is to be expected for the downflow lanes, the emerging intensity in vortex regions is preferentially smaller than the overall mean (bottom panel of Fig. \ref{fig:twodens}) )."
+ There is no clear trend towards higher intensity with increasing swirling strength., There is no clear trend towards higher intensity with increasing swirling strength.
+" This indicates that, although vortices are locally bright, the contrast is small compared to the brightness variations within intergranular lanes."," This indicates that, although vortices are locally bright, the contrast is small compared to the brightness variations within intergranular lanes."
+ We define vortex features to be contiguous regions of grid cells with a large swirling strength., We define vortex features to be contiguous regions of grid cells with a large swirling strength.
+" To achieve a clear separation, and take a lower cutoff for the swirling period than that used in the preceding section: zz > and <"," To achieve a clear separation, we consider only features above the (warped) optical surface and take a lower cutoff for the swirling period than that used in the preceding section: z > and <."
+ )..n6independent snapshotsof Run Cwe f ind3806 feature sthatsatis, In 6 independent snapshots of Run C we find 3806 features that satisfy these criteria (this number includes isolated grid cells).
+f ythe," The number of features decreases rapidly with size, see panel (a) of Fig. \ref{fig:contig}."
+, It is roughly proportional to the inverse square of the occupied volume (blue line).
+," In the following, we consider only the 622 features with at least 100 grid cells, corresponding to a minimum volume of"
+approximation.,approximation.
+ The model presented below is an attempt to include the most important features of this concept while still keeping the number of parameters low., The model presented below is an attempt to include the most important features of this concept while still keeping the number of parameters low.
+" The pure ice desorption energies are derived from fitting the three-phase model to the experimental pure ice desorption curves with the results: Ey,9= 4400K, Eco,=2440K and Eco=1010K."," The pure ice desorption energies are derived from fitting the three-phase model to the experimental pure ice desorption curves with the results: $E_{\rm H_2O} = 4400 \rm \, K$ , $E_{\rm CO_2} = 2440 \rm \, K$ and $E_{\rm CO} = 1010 \rm \, K$."
+ The H5O value is lower than the one found in ?.., The $_2$ O value is lower than the one found in \citet{Fraser_01}.
+" This discrepancy is probably due to a combination of that we use a single experiment, do not include the pumping speed and fix the pre-exponential factor to v=10s! for every species here."," This discrepancy is probably due to a combination of that we use a single experiment, do not include the pumping speed and fix the pre-exponential factor to $\nu = 10^{12} s^{-1}$ for every species here."
+ The value in ? should thus still be used when modeling the absolute desorption temperature., The value in \citet{Fraser_01} should thus still be used when modeling the absolute desorption temperature.
+ For the purpose of parameterizing the desorption fractions we prefer our value for the sake of consistency., For the purpose of parameterizing the desorption fractions we prefer our value for the sake of consistency.
+" The remainingman model parameters €,swa; and cj are obtained. separately for CO? and CO.", The remaining model parameters $\epsilon_{\rm H_2O- \textit{i}}^{\rm swap}$ and $c_i$ are obtained separately for $_2$ and CO.
+" This is done through a X? analysis, where trapping fractions from the model are compared to those from the binary experiments for a grid: of €5,0.swap and c; values."," This is done through a $\chi^2$ analysis, where trapping fractions from the model are compared to those from the binary experiments for a grid of $\epsilon_{\rm H_2O- \textit{i}}^{\rm swap}$ and $c_i$ values."
+" The minimumLo. x2 value for CO)» is obtained for €CO,-H,0 72500 K, cco, =20.5 ML, but the swapping energy for H?O:CO», which is linked to the ability of CO» to swap with H5O within the solid phase, is not well-constrained between 1600-3200 K (Fig. 5))."," The minimum $\chi^2$ value for $_2$ is obtained for $\epsilon_{\rm CO_2-H_2O}^{\rm swap}$ =2500 K, $c_{\rm CO_2}$ =20.5 ML, but the swapping energy for $_2$ $_2$, which is linked to the ability of $_2$ to swap with $_2$ O within the solid phase, is not well-constrained between 1600-3200 K (Fig. \ref{chisquare}) )."
+" In contrast, the parameter cco, related to available amount of CO, from the ice mantle that can migrate to the surface, is well constrained which suggests that even at laboratory time scales mixture desorption is mainly governed by how large a part of the mantle is eligible for swapping with the surface, rather than the swapping barriers."," In contrast, the parameter $c_{\rm CO_2}$ related to available amount of $_2$ from the ice mantle that can migrate to the surface, is well constrained which suggests that even at laboratory time scales mixture desorption is mainly governed by how large a part of the mantle is eligible for swapping with the surface, rather than the swapping barriers."
+" The CO experiments can be fitted with €CO-H,07 960 K and cco= 80 ML, but these are on only a few experiments and the inequalities €H,0-CObased<€,0-CO, and cco>cco, are alone well constrained."," The CO experiments can be fitted with $\epsilon_{\rm CO-H_2O}^{\rm swap}$ = 960 K and $c_{\rm CO}$ = 80 ML, but these are based on only a few experiments and the inequalities $ \epsilon_{\rm H_2O- CO}^{\rm swap}<\epsilon_{\rm H_2O- CO_2}^{\rm swap} $ and $c_{\rm CO}>c_{\rm CO_2}$ are alone well constrained."
+ Figure 6 shows the simulations of the binary mixture desorption using the optimized model parameters from the previous section., Figure \ref{tpd_sim} shows the simulations of the binary mixture desorption using the optimized model parameters from the previous section.
+ Generally the qualitative agreement is good and the model captures the trends that were observed experimentally., Generally the qualitative agreement is good and the model captures the trends that were observed experimentally.
+ Theexact shapes of the modeled and experimental desorption curves differ for several reasons., Theexact shapes of the modeled and experimental desorption curves differ for several reasons.
+" First, the model does not take into account the range of environments from which"," First, the model does not take into account the range of environments from which"
+of error bars for the time scales. the estimate of A7 is based on the goodness of fit. so that its normalization is not reliable.,"of error bars for the time scales, the estimate of $\chi^2$ is based on the goodness of fit, so that its normalization is not reliable."
+ The time scale model for the G cloud is overplotted in the left panel together with (he best [ive parameter fit V;=—1.7bonu -—34dkmstL.," The time scale model for the G cloud is overplotted in the left panel together with the best five parameter fit – $V_{iss,\alpha}=-1.7$, $V_{iss,\delta}=-3.4$."
+ The ISS parameters for this fit are A=3.5.04105”.," The ISS parameters for this fit are $A=3.5, \theta_{A}= 168^{\circ}$."
+ We note that these velocities are distinct [rom that of the Local Standard of Rest used by Jauncevetal.(2003) in their analvsis of the same time scale data., We note that these velocities are distinct from that of the Local Standard of Rest used by \citet{Jauncey2003} in their analysis of the same time scale data.
+ We conclude that the scintillation screen lies near the edge of the G cloud. since the Leo cloud lies about 20° [rom the source (Figure 1)).," We conclude that the scintillation screen lies near the edge of the G cloud, since the Leo cloud lies about $20^{\circ}$ from the source (Figure \ref{fig:fig1}) )."
+ The inner edee of the G cloud lies closer than 1.3 pe from the Sun and may extend several parsecs away from the inner edge., The inner edge of the G cloud lies closer than 1.3 pc from the Sun and may extend several parsecs away from the inner edge.
+ The orientation of the G cloud edge in celestial coordinates is about 110 compared to 94=108”., The orientation of the G cloud edge in celestial coordinates is about $110^{\circ}$ compared to $\theta_A = 168^{\circ}$.
+ Figure 5. (left) shows the excellent 5 Gllz ISS data overplotted from the vears 1999 for the quasar J13194-385 observed by Deunett-Thorpe&deBruvn (2003).., Figure \ref{fig:1819} (left) shows the excellent 5 GHz ISS data overplotted from the years 1999--2001 for the quasar J1819+385 observed by \citet{Dennett2003}. .
+ This source shows the largest intraday variability presently known at radio wavelengths &deBruyn 2001)., This source shows the largest intraday variability presently known at radio wavelengths \citep{Dennett2001}.
+". The best five parameter fit to these data is for Vj,=—1 aand V;,5=#20ἂν, which is within the uncertainty of the Mie cloud. velocity."," The best five parameter fit to these data is for $V_{iss,\alpha}=-1$ and $V_{iss,\delta}=+20$, which is within the uncertainty of the Mic cloud velocity."
+ The right panel shows the narrow band in (transverse velocity space corresponding to fits to the ISS data with reduced 4?<1.2., The right panel shows the narrow band in transverse velocity space corresponding to fits to the ISS data with reduced $\chi^2 <1.2$.
+ The transverse velocities of the Mic and LIC clouds are (he best fits to the ISS data. although the G and Oph clouds provide fits with reduced mLA.," The transverse velocities of the Mic and LIC clouds are the best fits to the ISS data, although the G and Oph clouds provide fits with reduced $\chi^2 \approx 1.8$."
+ Again. four of the six clouds have transverse velocities that give good-to-reasonable lits to the ISS data.," Again, four of the six clouds have transverse velocities that give good-to-reasonable fits to the ISS data."
+ The transverse velocity of the local standard of rest is an unacceptable fit to the data. as already reported by Dennett-Thorpe&deBruvun(2003).," The transverse velocity of the local standard of rest is an unacceptable fit to the data, as already reported by \citet{Dennett2003}."
+. Our analvsis. which follows that of Dennet(-Thorpe&deBruvn (2003).. agrees with their result Chat the ISS data lie in a narrow region of velocity space. as shown in Figure 5 (right).," Our analysis, which follows that of \citet{Dennett2003}, , agrees with their result that the ISS data lie in a narrow region of velocity space, as shown in Figure \ref{fig:1819} (right)."
+ By combining their time scale data with time celav data on an intercontinental baseline. (hev arrived at a best fit [or Ένα=—33.5 aand Vi;13.5m relative to the local standard of rest.," By combining their time scale data with time delay data on an intercontinental baseline, they arrived at a best fit for $V_{iss,\alpha}=-33.5$ and $V_{iss,\delta}=+13.5$ relative to the local standard of rest."
+ This is plotted with the diamond svmbol in the heliocentric reference svstem that we use in this paper., This is plotted with the diamond symbol in the heliocentric reference system that we use in this paper.
+ Their result is well removed [rom the values for the local clouds. but as stated in a footnote in their paper. Dennett- used an algorithm for interpreting the time delay data that was not corrected for anisotropic scintillation.," Their result is well removed from the values for the local clouds, but as stated in a footnote in their paper, \citet{Dennett2003} used an algorithm for interpreting the time delay data that was not corrected for anisotropic scintillation."
+ We thereforemake no further use of that valueplotted by the, We thereforemake no further use of that valueplotted by the
+As one of the nearest voung (125Myr:?). open clusters. the Pleiades have long been recognized as an important astrophvsical laboratory for studying stellar evolution and the dvnanmiucs of stellar associations.,"As one of the nearest young \citep[125~Myr;][]{stauffer1998} open clusters, the Pleiades have long been recognized as an important astrophysical laboratory for studying stellar evolution and the dynamics of stellar associations."
+ Multüiplicityv studies of the Pleiades have focused both on stellar (e.g..?27) and substellar (e.g..??) members of the cluster.," Multiplicity studies of the Pleiades have focused both on stellar \citep[e.g.,][]{stauffer1984, mermilliod1992, bouvier1997} and substellar \citep[e.g.,][]{martin2000, bouy2006} members of the cluster."
+ ILowever. mixed dwarf svstems have not been. identified.," However, mixed star--brown dwarf systems have not been identified."
+ One of the most extensive studies is the adaptive optics imaging survey of ?.., One of the most extensive studies is the adaptive optics imaging survey of \citet{bouvier1997}.
+ Conducted in the near-IR. it covered. 144 G and Ix stars to a relatively shallow depth.," Conducted in the near-IR, it covered 144 G and K stars to a relatively shallow depth."
+ As a result. its sensitivitv encompassed only stellar and massive substellar companions.," As a result, its sensitivity encompassed only stellar and massive substellar companions."
+ A svstematie high-contrast imaging survey of the Pleiades on a similar scale but at a higher sensiGvily has not been perlormed since., A systematic high-contrast imaging survey of the Pleiades on a similar scale but at a higher sensitivity has not been performed since.
+ Consequently. no substellar companions are known to 20.2 sslars in the Pleiades.," Consequently, no substellar companions are known to $>$ 0.2 stars in the Pleiades."
+" With the frequency of wide substellar companions to field-aged stars now estimated at 22 (0.012.0.072AL,browndwarlsin28.1590AUorbits:?).. the lrequeney of brown dwarf secondares around Sun-like stars in (hie Pleiades is expected to be comparable."," With the frequency of wide substellar companions to field-aged Sun-like stars now estimated at $\approx$ \citep[0.012\,--\,0.072~\Msun\ brown dwarfs in 28\,--\,1590\,AU orbits;][]{metchev2009}, the frequency of brown dwarf secondaries around Sun-like stars in the Pleiades is expected to be comparable."
+ In the present paper. we announce the identification of a low mass companion to the Pleiacl HII 1348. a VV. double-lined spectroscopic binary (7.hereafterreferedtoasprimaryorHII 1348A)..," In the present paper, we announce the identification of a low mass companion to the Pleiad HII 1348, a V double-lined spectroscopic binary \citep[hereafter refered to as the 
+primary or HII~1348A]{queloz1998}."
+ The faint companion. HII 1343D. was already detected. by 2..," The faint companion, HII 1348B, was already detected by \citet{bouvier1997}."
+ llowever. without follow-up astrometric observations and due to a non-negligible probability of background star contamination. ? conservatively assumed that the candidate companion was an unrelated background star.," However, without follow-up astrometric observations and due to a non-negligible probability of background star contamination, \citet{bouvier1997} conservatively assumed that the candidate companion was an unrelated background star."
+ The astrometric measurements confirm the proper motion association of the pair. ancl AO spectra obtained at Palomar and Neck reveal that the companion has a spectral type of M8.," The astrometric measurements confirm the proper motion association of the pair, and AO spectra obtained at Palomar and Keck reveal that the companion has a spectral type of M8."
+empirical approach. ancl represent the uncertainty in terms of two parameters (Macau et al.,"empirical approach, and represent the uncertainty in terms of two parameters (Madau et al."
+ 1998: Dahlen Fransson 1999)., 1998; Dahlen Fransson 1999).
+ The first is a delav time 7. between the binary system. formation. and SNe explosion epochs. which defines a (ime-independent) explosion. probability per white dwarf.," The first is a delay time $\tau$, between the binary system formation and SNe explosion epochs, which defines a (time-independent) explosion probability per white dwarf."
+ The second is an explosion elliciency.. jg. which accounts for the fraction of binary svstems that never result in a SNe.," The second is an explosion efficiency, $\eta$, which accounts for the fraction of binary systems that never result in a SNe."
+ This cllicicney a is constrained using a preliminary determination of the lla rate at z=0.55 of 1.66+0426°10tAMIpeter by Pain et al.," This efficiency $\eta$ is constrained using a preliminary determination of the Ia rate at $z=0.55$ of $1.66\pm0.42\,h^{3}10^{-4}\rmn{Mpc}^{-3}\rmn{yr}^{-1}$ by Pain et al."
+ (2000a) using a sample of 38 La from the SCP (see also Pain et al., \shortcite{pain00a} using a sample of 38 Ia from the SCP (see also Pain et al.
+ 2000b)., 2000b).
+ We consider two SNella evolutions. shown in Fig. 2..," We consider two Ia evolutions, shown in Fig. \ref{snerates}."
+ ‘The first has à small value of 7=0.3€ivr (corresponding to a shallower decline at high-:2.) whilst the second is a larger value of 723.0Gyr. which produces a steeper cdrop-oll.," The first has a small value of $\tau=0.3\,\rmn{Gyr}$ (corresponding to a shallower decline at $z$ ,) whilst the second is a larger value of $\tau=3.0\,\rmn{Gyr}$, which produces a steeper drop-off."
+ For comparison. the z0.55 rate of Pain et al.," For comparison, the $z=0.55$ rate of Pain et al."
+ (2000a) is also shown in Fig. 2..," \shortcite{pain00a} is also shown in Fig. \ref{snerates},"
+ together with the dillerence that arises from using the two SEL. and how these histories tend to over or under-prediet the local SNe rates.," together with the difference that arises from using the two SFHs, and how these histories tend to over or under-predict the local SNe rates."
+ While much. published. data exists. for SNella. by contrast there is à dearth of data for SNelll and a great varicty in their spectra. peak =nagnitudes and light curves.," While much published data exists for Ia, by contrast there is a dearth of data for II and a great variety in their spectra, peak magnitudes and light curves."
+ For the purposes of this paper. we adopt the moclels of Gilliland. Nugent Phillips (1999) (hereafter GNP) which incorporate recent light curve templates. A-corrections and luminosity functions lor lla and LH.," For the purposes of this paper, we adopt the models of Gilliland, Nugent Phillips \shortcite{gilliland99}
+ (hereafter GNP) which incorporate recent light curve templates, $k$ -corrections and luminosity functions for Ia and II."
+ These mocels include. observed. dispersions in. the peak magnitudes: of SNella anc Hl. which are included in our simulations of our SNe counts (see GNP for a full cliseussion of the models and simulations).," These models include observed dispersions in the peak magnitudes of Ia and II, which are included in our simulations of our SNe counts (see GNP for a full discussion of the models and simulations)."
+ The major advantage to using INPS mocels Lies in the fact that the Spectral Energy. Distributions (SEDs) are more realistic than those emploved. previously for rate caleulations and follow the true temporal evolution of these objects., The major advantage to using GNP's models lies in the fact that the Spectral Energy Distributions (SEDs) are more realistic than those employed previously for rate calculations and follow the true temporal evolution of these objects.
+ The SNella template was constructed using data from and citeiuesne.kirO2a in the UV and a large number of erouncl- observations in the optical.," The Ia template was constructed using data from and \\cite{iuesne,kir92a} in the UV and a large number of ground-based observations in the optical."
+ Spectrum: synthesis calculations were carried out bv PN for the SNellls using the non-LTE code PILOENIN 9.0 (Hauschildt.ctal.1997:HIauschildt.Baron.&Allard. 1997)..," Spectrum synthesis calculations were carried out by PN for the II's using the non-LTE code PHOENIX 9.0 \cite{hb971,hb972}."
+ Ehe spectra. of SNe11979€ (Cappollaro.Turatto&Fernley.1995)... 199211 and 1993. (courtesy D. Leonard) were fit at several epochs.," The spectra of 1979C \cite{iuesne}, 1992H and 1993W (courtesy D. Leonard) were fit at several epochs."
+ The resultant fits of these spectra provided a nice match to the observed. data and. vield a crude SED template for the SNellUs., The resultant fits of these spectra provided a nice match to the observed data and yield a crude SED template for the II's.
+ The major emphasis of this elfort was to model the UV as well as possible (this portion of the SED dominates in high-z searches using filters bluer than yma). since the simple assumption of blackbody SEDs has been shown to be quite inaccurate (see fig.," The major emphasis of this effort was to model the UV as well as possible (this portion of the SED dominates in $z$ searches using filters bluer than $1\mu\rmn{m}$ ), since the simple assumption of blackbody SEDs has been shown to be quite inaccurate (see fig."
+ 10 of GND)., 10 of GNP).
+ We can now combine the two mass models from 82.2. and the various SNe rate moclels from 82.3. to estimate the number counts of SNe with and without the presence of a strong-lensing cluster., We can now combine the two mass models from $\S$ \ref{modelclusters} and the various SNe rate models from $\S$ \ref{modelsne} to estimate the number counts of SNe with and without the presence of a strong-lensing cluster.
+ As with blank-feld surveys. searches such as. those outlined here require two images of the same rich cluster. but with the two epochs of observation separated by a suitable amount to allow the SNe to become visible.," As with blank-field surveys, searches such as those outlined here require two images of the same rich cluster, but with the two epochs of observation separated by a suitable amount to allow the SNe to become visible."
+ When estimating the simple count rate of lensed high-z SNe. therefore. both forward (SNe in second image) and backward (SNe in first image) searches can be included. and to allow for the time dilation of the higsh-z SNe. we assume a time of at least 2vvrs between the two exposures.," When estimating the simple count rate of lensed $z$ SNe, therefore, both forward (SNe in second image) and backward (SNe in first image) searches can be included, and to allow for the time dilation of the $z$ SNe, we assume a time of at least yrs between the two exposures."
+ Table 1 eives the counts expected to survey [limits of {κι=26.0 27.0. assuming τομ=0.2&0.6.," Table \ref{count_table} gives the counts expected to survey limits of $I_{814}=26.0$ $27.0$ , assuming $z_{\rmn{cluster}}=0.2\,\&\,0.6$."
+ When comparing lensed and unlensed: counts. one of )0 Κον factors is the slope of the SNe number counts. a—dlogV/dim.," When comparing lensed and unlensed counts, one of the key factors is the slope of the SNe number counts, $\alpha=\rmn{d}\log N/\rmn{d}m$."
+ For az0.4. we expect lensing to increase 1e number counts. while for a«OA. lensing will have 16 opposite clleet (depletion).," For $\alpha>0.4$, we expect lensing to increase the number counts, while for $\alpha<0.4$, lensing will have the opposite effect (depletion)."
+ The slopes in our. models epend on i) the intrinsic SNe properties. contained in the models of GNP. and ii) our assumed. redshift evolution of 10 SNe number density.," The slopes in our models depend on i) the intrinsic SNe properties, contained in the models of GNP, and ii) our assumed redshift evolution of the SNe number density."
+ “Phe models of GNP include the xst-determined: uncertainties in the SNe properties. for example in the dispersion of the peak SNe magnitude. and our simulations therefore account for this first uncertainty.," The models of GNP include the best-determined uncertainties in the SNe properties, for example in the dispersion of the peak SNe magnitude, and our simulations therefore account for this first uncertainty."
+ The redshift evolution of the SNe numbers depends on which SELL we use. and. for Wa. the value of the delay time. 7.," The redshift evolution of the SNe numbers depends on which SFH we use, and, for Ia, the value of the delay time, $\tau$."
+ ὃν running simulations for the unlensec case. we find that tvpical values of a for LE for SELI-L are àστ0.35. and for SELI-LE a20.32.," By running simulations for the unlensed case, we find that typical values of $\alpha$ for II for SFH-I are $\alpha\simeq0.35$, and for SFH-II $\alpha\simeq0.32$."
+ TFhere is little variation in these values between simulation runs., There is little variation in these values between simulation runs.
+ The number-redshift. relation is plotted in Fig. 3.., The number-redshift relation is plotted in Fig. \ref{magnz}.
+ As one would expect. the figure shows there is little cilfercnee in the rate of SNe up to zai," As one would expect, the figure shows there is little difference in the rate of SNe up to $z_{\rmn{cluster}}$."
+ From tenuis to 21. the lensecl counts are lower than those appropriate for blank field searches.," From $z_{\rmn{cluster}}$ to $z\simeq1$, the lensed counts are lower than those appropriate for blank field searches."
+ This is due to the clleet of depletion mentioned above. which unfortunately negates the lensing benefits.," This is due to the effect of depletion mentioned above, which unfortunately negates the lensing benefits."
+ Llowever. the important differences begin at z1. where we sample a fainter. more numerous population of SNe.," However, the important differences begin at $z>1$, where we sample a fainter, more numerous population of SNe."
+ Lensing jas the principal ellect of increasing the proportion of SNe ound at high-z. bringing fainter. magnified examples into view.," Lensing has the principal effect of increasing the proportion of SNe found at $z$, bringing fainter, magnified examples into view."
+ Phere is then a corresponding extension in the recshilt distribution., There is then a corresponding extension in the redshift distribution.
+ For example. for the z=0.2 NEW cluster. he mean redshift. of detection changes from z0.70 to 20.50. and the number of detections at z>1 increases rom 20 to 31 per cent.," For example, for the $z=0.2$ NFW cluster, the mean redshift of detection changes from $z\simeq0.70$ to $z\simeq0.85$, and the number of detections at $z>1$ increases from 20 to 31 per cent."
+ The numbers drop sharply by z=2. even in the lensed case. due to the UV cdrop-olf in the SNe spectrum. being redshifted into the {κι filter. which prevents the detection of higher-z SNe.," The numbers drop sharply by $z=2$, even in the lensed case, due to the UV drop-off in the SNe spectrum being redshifted into the $I_{814W}$ filter, which prevents the detection of $z$ SNe."
+We note that the use of aand an infra-red search would allow the detection of lensed 2 SNe which would be undetectable in the unlensed Case.,We note that the use of and an infra-red search would allow the detection of lensed $z>2$ SNe which would be undetectable in the unlensed case.
+ Clearly. the number of detectable SNe depends. on many," Clearly, the number of detectable SNe depends on many"
+to coincide with the visible solar limb.,to coincide with the visible solar limb.
+" The NaD, line is often called a chromospheric diagnostic but 3-dimensional non-local thermodynamic equilibrium (nonLTE) radiative transfer calculations of magnetoconvection simulations have demonstrated that the line is to a large degree photospheric. especially in magnetic flux concentrations (Leenaarts 2009)."," The $_{1}$ line is often called a chromospheric diagnostic but 3-dimensional non-local thermodynamic equilibrium (nonLTE) radiative transfer calculations of magnetoconvection simulations have demonstrated that the line is to a large degree photospheric, especially in magnetic flux concentrations \citep{Leenaarts+others2009}."
+. For SP data we use 10 raster scans with exposure times of eight seconds or more (1.e.. deep magnetogram mode) at varying positions on the solar disk (see Table 1)) reduced using the Solar Software package routines andsimages-sbsp.," For SP data we use 10 raster scans with exposure times of eight seconds or more (i.e., deep magnetogram mode) at varying positions on the solar disk (see Table \ref{tab:nfi-sp-dat}) ) reduced using the Solar Software package routines and."
+"pro. The step size is 0.146 "" in all scans but number 5 where it is 0.29 ”.", The step size is 0.149 $^{\prime \prime}$ in all scans but number 5 where it is 0.29 $^{\prime \prime}$.
+ The reduced data consist of maps of circular polarization and linear polarization in. the 630.1 (effective g=1.67) and 630.2 nm (effective g=2.5) Fe | lines. both combined and separate.," The reduced data consist of maps of circular polarization and linear polarization in the 630.1 (effective $g$ =1.67) and 630.2 nm (effective $g$ =2.5) Fe I lines, both combined and separate."
+ For details of the routines see Litesetal.(2008)., For details of the routines see \cite{Lites+others2008}.
+". As in NaD, data. pointing errors are corrected in maps where the solar limb ts visible."," As in $_{1}$ data, pointing errors are corrected in maps where the solar limb is visible."
+ Flux concentrations (coherent features with circular polarization signal clearly above the noise level) are identified manually m each NFI and SP circular polarization image., Flux concentrations (coherent features with circular polarization signal clearly above the noise level) are identified manually in each NFI and SP circular polarization image.
+ For the identification we use images of the absolute value of the flux instead of the signed flux to reduce the bias towards choosing bipolar features near the limb., For the identification we use images of the absolute value of the flux instead of the signed flux to reduce the bias towards choosing bipolar features near the limb.
+ A central pixel of each identified feature is chosen and a vector connecting the disk center and limb passing through this point is defined., A central pixel of each identified feature is chosen and a vector connecting the disk center and limb passing through this point is defined.
+ This vector is the radial cut of the feature., This vector is the radial cut of the feature.
+ The SP circular and linear polarization radial cuts are made by averaging over 7 pixels perpendicular to the vector for the 630.1 and 630.2 nm., The SP circular and linear polarization radial cuts are made by averaging over 7 pixels perpendicular to the vector for the 630.1 and 630.2 nm.
+ Unless otherwise specified we use the average of the two lines., Unless otherwise specified we use the average of the two lines.
+ For NFI data the radial cuts (width 7 pixels) are made of the Stokes V filter signal., For NFI data the radial cuts (width 7 pixels) are made of the Stokes $V$ filter signal.
+ No correction for foreshortening is made., No correction for foreshortening is made.
+ Finally. peaks (location and amplitude) in the radial cuts are identified manually in each cut (see Fig. 1)).," Finally, peaks (location and amplitude) in the radial cuts are identified manually in each cut (see Fig. \ref{fig:nad-e1}) )."
+ We have employed two independent sets of models. the simple thin-tube approximation and realistic 3-D simulations.," We have employed two independent sets of models, the simple thin-tube approximation and realistic 3-D simulations."
+" 0"" order in radius. r. for the vertical component of the magnetic field. B-. and 1 order in the radial component of the field. B,."," $0^{th}$ order in radius, $r$, for the vertical component of the magnetic field, $B_{z}$, and $1^{st}$ order in the radial component of the field, $B_{r}$."
+ The external non-magnetic atmosphere 1s given by the semi-empirical SRPM model of the quiet Sun low chromosphere (Fontenlaetal.2007) and internal atmosphere by the recent facular and plage models of the photosphere (Fontenlaetal.2006) which have been extended in height by adding the FALF (facular model. Fontenlaetal.1993)) model chromosphere.," The external non-magnetic atmosphere is given by the semi-empirical SRPM model of the quiet Sun low chromosphere \citep{SRPM2007} and internal atmosphere by the recent facular and plage models of the photosphere \citep{newFAL2006} which have been extended in height by adding the FALF (facular model, \citealt{FAL1993}) ) model chromosphere."
+ We produce 2-dimensional atmospheric slabs (x.2) by varying the flux tube/sheet parameters: radius at zzO km (70). field strength at zzO km (Bo). and the internal flux tube/sheet atmosphere (facular or plage).," We produce 2-dimensional atmospheric slabs $x,z$ ) by varying the flux tube/sheet parameters: radius at $z$ =0 km $r_{0}$ ), field strength at $z$ =0 km $B_{0}$ ), and the internal flux tube/sheet atmosphere (facular or plage)."
+ Observables are synthesized under the assumption of LTE using the Nicole-code of Socas-Navarro (Socas-Navarro2001)., Observables are synthesized under the assumption of LTE using the Nicole-code of Socas-Navarro \citep{Melanie}.
+". Arguably. especially the NaD, line is influenced by nonLTE effects. but since we consider the Stokes V signal which has its? maximum value away from the line core and. more specifically. in the ratio of the Stokes V amplitudes. we chose to do the synthesis in LTE."," Arguably, especially the $_{1}$ line is influenced by nonLTE effects, but since we consider the Stokes $V$ signal which has its' maximum value away from the line core and, more specifically, in the ratio of the Stokes $V$ amplitudes, we chose to do the synthesis in LTE."
+" We are not interested in comparing absolute values of the amplitudes or details of the line profiles. which may well be affected by nonLTE effects. but rather in the ratio between the amplitudes of the center- and limb-sides of magnetic concentrations, nt which should be independent of nonLTE effects. at least to first order."," We are not interested in comparing absolute values of the amplitudes or details of the line profiles, which may well be affected by nonLTE effects, but rather in the ratio between the amplitudes of the center- and limb-sides of magnetic concentrations, $\frac{a_c}{a_l}$, which should be independent of nonLTE effects, at least to first order."
+ We compute |-dimensional rays passing through the slabs at various line of sight (LOS) angles and produce synthetic observables., We compute 1-dimensional rays passing through the slabs at various line of sight (LOS) angles and produce synthetic observables.
+ Synthetic cuts are created from the Stokes V amplitudes and areas (sign determined by the sign of the blue lobe)., Synthetic cuts are created from the Stokes $V$ amplitudes and areas (sign determined by the sign of the blue lobe).
+" This is done also for the NaD, line instead of using the filter transmission profile.", This is done also for the $_1$ line instead of using the filter transmission profile.
+" By doing this the NaD, synthetic observables pick up the maximum signal at each point.", By doing this the $_1$ synthetic observables pick up the maximum signal at each point.
+ We compute the using the original tube/sheet nodel resolution (10km in the horizontal direction. rays placed at 10 km intervals) and also for the spectral (90 ffor NFL 23 Που SP) and spatial (theoretical point spread function (PSF) for a 50 cm telescope. 0.08” and 0.16” pixels size for NFI and SP. respectively) resolutions comparable to those of Hinode.," We compute the using the original tube/sheet model resolution (10km in the horizontal direction, rays placed at 10 km intervals) and also for the spectral (90 for NFI, 23 for SP) and spatial (theoretical point spread function (PSF) for a 50 cm telescope, $^{\prime \prime}$ and $^{\prime \prime}$ pixels size for NFI and SP, respectively) resolutions comparable to those of Hinode."
+ For the purposes of this paper the difference between flux tubes and sheets lies in their linear expansion rates., For the purposes of this paper the difference between flux tubes and sheets lies in their linear expansion rates.
+ In the thin tube approximation the expansion is driven by the magnetic flux conservation with height: BA= const. Where B is the field strength and A 1s the tube’s cross-section.," In the thin tube approximation the expansion is driven by the magnetic flux conservation with height: $BA=const$ , where $B$ is the field strength and $A$ is the tube's cross-section."
+ As B drops, As $B$ drops
+For each of our simulated samples. we have run a Ilxolmogorov-5miürnov (est. which provides the probability that. a model is derived from the same underlving population as the observed sample.,"For each of our simulated samples, we have run a Kolmogorov-Smirnov test, which provides the probability that a model is derived from the same underlying population as the observed sample."
+ Not surprisineglv. none of our models have produced. a significance level higher5 than156.. the highest.5 being5 lor model L. However. it is interesting5 to examine more closely (he source of the discrepancy in (he low-eccentricitv {ο«0.1) and high-eecentricity (ο> 0.6) regimes.," Not surprisingly, none of our models have produced a significance level higher than, the highest being for model L. However, it is interesting to examine more closely the source of the discrepancy in the low-eccentricity $e<0.1$ ) and high-eccentricity $e>0.6$ ) regimes."
+ In most of our simulations. the Ixozai mechanism (ends to overproduce planets with very low orbital eccentricities.," In most of our simulations, the Kozai mechanism tends to overproduce planets with very low orbital eccentricities."
+ The lowest quartile of final eecentricities in anv of (he models is much less (han 0.1. whereas in (he observed sample this is 0.14.," The lowest quartile of final eccentricities in any of the models is much less than 0.1, whereas in the observed sample this is 0.14."
+ There are several reasons for this overabundance of low eccentricities in our model svstems., There are several reasons for this overabundance of low eccentricities in our model systems.
+ First. since we do not have anv observational constraints on relative inclination angles. we have assumed an isotropic distribution of 7).," First, since we do not have any observational constraints on relative inclination angles, we have assumed an isotropic distribution of $i_0$."
+ This implies that 23% of the svstems have /j«Z4. resulting in no Ixozai oscillation.," This implies that $23\%$ of the systems have $i_0<i_{\rm crit}$, resulting in no Kozai oscillation."
+" However. in the total observed. sample. planets with ο,<0.1 are only of the total (or if we exclude multi-planet svstems and bot Jupiters with a,<0.1 AU)."," However, in the total observed sample, planets with $e_1<0.1$ are only of the total (or if we exclude multi-planet systems and hot Jupiters with $a_1<0.1\,$ AU)."
+ Svstems with sufficient initial relative inclination angles still need to overcome other hurdles to achieve hiehlv eccentric orbits., Systems with sufficient initial relative inclination angles still need to overcome other hurdles to achieve highly eccentric orbits.
+ LD many of the binary companions are substellar or in very wide orbits. lxozai periods become so long that the eccentricity oscillation are either suppressed by GR precession. or not completed within the age of (he svstem (or both).," If many of the binary companions are substellar or in very wide orbits, Kozai periods become so long that the eccentricity oscillation are either suppressed by GR precession, or not completed within the age of the system (or both)."
+ This can result in an additional of planets remaining in nearly circular orbits., This can result in an additional of planets remaining in nearly circular orbits.
+ Even when the orbits of the planets do undergo eccentricity oscillations. about just happen to be observed al low eccentricities.," Even when the orbits of the planets do undergo eccentricity oscillations, about just happen to be observed at low eccentricities."
+ Thus. our results suggest that the observed. simple was a remarkably small population of planets in nearly circular orbits. aud other ανασα] processes must Clearly be invoked to perturb their orbits.," Thus, our results suggest that the observed sample has a remarkably small population of planets in nearly circular orbits, and other dynamical processes must clearly be invoked to perturb their orbits."
+ Among the most likely mechanisms is planet.planet scattering in multi-planet svstems. which can easily perturb eccentricities (o nodest. values in the intermediate range ~0.2—0.6 (Rasio&Ford1996:WeidenschillingMarzari1996:&Weidenschilling 2002).," Among the most likely mechanisms is planet–planet scattering in multi-planet systems, which can easily perturb eccentricities to modest values in the intermediate range $\sim 0.2-0.6$ \citep{rasioford96,weiden96,marz02}."
+. Clear evidence that planet scattering nust have occurred in the ο Andromedae system has been presented by Ford. Lvstad. Rasio (2005).," Clear evidence that planet–planet scattering must have occurred in the $\upsilon$ Andromedae system has been presented by Ford, Lystad, Rasio (2005)."
+ Even in most of the svstems in which only one giant planet has been detected so lar. the second planet could have been ejected as a result of the scattering. or it could have been retained in a much wider. eccentric orbit. making it hard to detect by Doppler spectroscopy.," Even in most of the systems in which only one giant planet has been detected so far, the second planet could have been ejected as a result of the scattering, or it could have been retained in a much wider, eccentric orbit, making it hard to detect by Doppler spectroscopy."
+ In the high-eecentricity region. where e420.6. our models show much better agreement with the observed distribution.," In the high-eccentricity region, where $e_1\ga 0.6$, our models show much better agreement with the observed distribution."
+ The lXozai mechanism predicts a small excess of svstenis al the highest eccentricities (ο> 0.9). although it should be noted that the observed eccentricity distribution in this range is not vel well constrained.," The Kozai mechanism predicts a small excess of systems at the highest eccentricities $e>0.8$ ), although it should be noted that the observed eccentricity distribution in this range is not yet well constrained."
+ It. is evident that the, It is evident that the
+"We ascribe the upper kHz QPO frequency profile to the X- flux strong variations from transition layer, as shown in the upper panel of Fig.(1)).","We ascribe the upper kHz QPO frequency profile to the X-ray flux strong variations from transition layer, as shown in the upper panel of \ref{model}) )."
+" Thus, the range of the upper frequency (dv) should be the indication of transition layer radial extent (ór or in terms of the scaled quantity X)."," Thus, the range of the upper frequency $\delta \nu$ ) should be the indication of transition layer radial extent $\delta r$ or in terms of the scaled quantity $X$ )."
+" =3,", = =.
+" If we set the QPO FWHM w as corresponding to the scale of transition layer width, i.e. w~dy, then the quality factor for upper frequency (hereafter upper Q-factor) can be written as, We consider a transition layer (see upper panel of Fig. (1)))"," If we set the QPO FWHM $w$ as corresponding to the scale of transition layer width, i.e. ${\rm w \sim \delta \nu}$, then the quality factor for upper frequency (hereafter upper Q-factor) can be written as, We consider a transition layer (see upper panel of Fig. \ref{model}) ))"
+" between the innermost Keplerian orbit and the magnetosphere (Elsner Lamb 1977; 2010ab), in which the transition for radial velocity of accretion flow from a Keplerian to a corotation with the NS may occur (Titarchuk, Lapidus Muslimov 1999 hereafter TLM; Titarchuk Osherovich 1999)."," between the innermost Keplerian orbit and the magnetosphere (Elsner Lamb 1977; ), in which the transition for radial velocity of accretion flow from a Keplerian to a corotation with the NS may occur (Titarchuk, Lapidus Muslimov 1999 hereafter TLM; Titarchuk Osherovich 1999)."
+ But in this region the formation of kinks and shocks due to the inhomogeneous density of flow and supersonic motions of accretion flow may be responsible for a super-Keplerian rotation (TLM)., But in this region the formation of kinks and shocks due to the inhomogeneous density of flow and supersonic motions of accretion flow may be responsible for a super-Keplerian rotation (TLM).
+" After a self-adjustment which is the function of this layer, the coupling of the sub-Keplerian flow with the super-Keplerian rotation still leads to a Keplerian zone which contributes to the radial drift."," After a self-adjustment which is the function of this layer, the coupling of the sub-Keplerian flow with the super-Keplerian rotation still leads to a Keplerian zone which contributes to the radial drift."
+" When the accreted matter fall into the innermost Keplerian orbit and transition zone, these plasma may strike the magnetospheric boundary and bend the field lines, which lead to the change of the magnetospheric shape due to the strong ram pressure and can excited the Alfvén wave oscillation."," When the accreted matter fall into the innermost Keplerian orbit and transition zone, these plasma may strike the magnetospheric boundary and bend the field lines, which lead to the change of the magnetospheric shape due to the strong ram pressure and can excited the $\grave{e}$ n wave oscillation."
+" In the meantime, the plasma is threaded by the field lines and fall onto the polar cap."," In the meantime, the plasma is threaded by the field lines and fall onto the polar cap."
+" However, on the magnetospheric boundary, the accreted matter carries different velocities and has different density from the former arrivals, which bring into some instabilities, such as Kelvin-Helmholtz instability, Rayleigh-Taylor instability, and the inhomogeneities caused by the azimuthal component of the field trapped inside the inner regions of the disk (Romanova, Kulkarni Lovelace 2007)."," However, on the magnetospheric boundary, the accreted matter carries different velocities and has different density from the former arrivals, which bring into some instabilities, such as Kelvin-Helmholtz instability, Rayleigh-Taylor instability, and the inhomogeneities caused by the azimuthal component of the field trapped inside the inner regions of the disk (Romanova, Kulkarni Lovelace 2007)."
+" The Kelvin-Helmholtz instability seems to be less significant (Rastátter Schindler 1999), and we just focus on the Rayleigh-Taylor instability (Kulkarni Romanova 2008)."," The Kelvin-Helmholtz instability seems to be less significant $\ddot{a}$ tter Schindler 1999), and we just focus on the Rayleigh-Taylor instability (Kulkarni Romanova 2008)."
+ The unstable disc-magnetosphere boundary results in the penetration of the magnetosphere by the disc matter in kinds of forms., The unstable disc-magnetosphere boundary results in the penetration of the magnetosphere by the disc matter in kinds of forms.
+" In addition, the disc matter treading the field lines into the magnetosphere from polar cap may move to the equatorial region and swell the magnetosphere."," In addition, the disc matter treading the field lines into the magnetosphere from polar cap may move to the equatorial region and swell the magnetosphere."
+" On the grounds of this idea and the geometry of accretion disc boundary layer (Regev Hougerat 1988), the radial extent direction can be written as (Elsner Lamb 1977; 2010ab), where the quantity (r-R) stands for the scale of the magnetosphere-disk boundary to stelar surface, and a<1 is a constant ratio coefficient."," On the grounds of this idea and the geometry of accretion disc boundary layer (Regev Hougerat 1988), the radial extent direction can be written as (Elsner Lamb 1977; ), where the quantity (r-R) stands for the scale of the magnetosphere-disk boundary to stelar surface, and $\alpha < 1$ is a constant ratio coefficient."
+" Therefore, we can obtain, Using the definition of X and Eq. (3)),"," Therefore, we can obtain, Using the definition of X and Eq. \ref{upper}) ),"
+" we can write Q» with respect to the upper frequency, For Πιτthe Gssoa)detected twin kHz QPOs, the mass density parameter A is found to be about 0.7 (e.g. Sco X-1) (Zhang 2004; Zhang et al."," we can write $_2$ with respect to the upper frequency, For the detected twin kHz QPOs, the mass density parameter $A$ is found to be about 0.7 (e.g. Sco X-1) (Zhang 2004; Zhang et al."
+" 2007), except for the two unusual X-ray millisecond pulsar cases SAX J1804.5-3654 and XTE J1807-294, A=0.45 and 0.4 respectively (Zhang et al."," 2007), except for the two unusual X-ray millisecond pulsar cases SAX J1804.5-3654 and XTE J1807-294, A=0.45 and 0.4 respectively (Zhang et al."
+ 2010)., 2010).
+" In most cases (except Cir X-1), the position parameter X=R/r lies in the range from 0.7 to 0.92, or the kHz QPO emission position radius is from r=1.1R to r=1.4R (Zhang et al."," In most cases (except Cir X-1), the position parameter X=R/r lies in the range from 0.7 to 0.92, or the kHz QPO emission position radius is from r=1.1R to r=1.4R (Zhang et al."
+ 2010)., 2010).
+" Averagely, we have the following expression to evaluate the upper Q-factor, To make a comparison between the model and the detections of the upper Q-factor, we firstly put the sources whose spectrum of kHz QPOs displayed the high coherence together in Fig. 2."," Averagely, we have the following expression to evaluate the upper Q-factor, To make a comparison between the model and the detections of the upper Q-factor, we firstly put the sources whose spectrum of kHz QPOs displayed the high coherence together in Fig. \ref{Q}."
+" For the Atoll source 4U 1728-34, there is an abrupt drop at v~1050 Hz in the Q-v plot, then it is argued as an effect of innermost stable"," For the Atoll source 4U 1728-34, there is an abrupt drop at $\nu \sim 1050$ Hz in the $\nu$ plot, then it is argued as an effect of innermost stable"
+generate cometary globules and trigger gravitational collapse. but it is more difficult to form pillars like those in M16 because gas must accumulate to a high density in the shadowed tail region.,"generate cometary globules and trigger gravitational collapse, but it is more difficult to form pillars like those in M16 because gas must accumulate to a high density in the shadowed tail region."
+ ? showed how photoionisation of multiple clumps which partially shadow each other leads to dense gas accumulating in shadowed ail regions: it was suggested by ? that multiple clumps are required to quantitatively match observations of the pillars in 116., \citet{LimMel03} showed how photoionisation of multiple clumps which partially shadow each other leads to dense gas accumulating in shadowed tail regions; it was suggested by \citet{PouKanRyuEA07} that multiple clumps are required to quantitatively match observations of the pillars in M16.
+ Recent models (222??.MLIO) showed that elephant trunks can form quite naturally from the photoionisation of a clumpy density field under a range of initial conditions.," Recent models \citep[][ML10]{MelArtHenEA06, RagHenVasEA09, LorRagEsq09,
+ GriNaaWalEA09, GriBurNaaEA10} showed that elephant trunks can form quite naturally from the photoionisation of a clumpy density field under a range of initial conditions."
+ ?. extended our orevious. work (Μ.Ο) which used static initial conditions by considering the photoionisation of dynamic density tields generated by isothermal decaying turbulence. measuring the formation of dillar-like structures as a function of turbulent Mach number.," \citet{GriBurNaaEA10} extended our previous work (ML10) which used static initial conditions by considering the photoionisation of dynamic density fields generated by isothermal decaying turbulence, measuring the formation of pillar-like structures as a function of turbulent Mach number."
+ Filamentary structure in an apparently helical geometry was ‘ound in the shadowed tail regions of a number of elephant citepCarkriGah9Ys.CarGahKri02.. CarGahKri03.. and. it) was suggested this could arise due to twisting of magnetic field lines.," Filamentary structure in an apparently helical geometry was found in the shadowed tail regions of a number of elephant \\citep{CarKriGah98, CarGahKri02, CarGahKri03}, and it was suggested this could arise due to twisting of magnetic field lines."
+ Velocity protiles for some trunks were shown to be consistent with solid-body rotation (2).. again consistent with a magnetic origin of the observed structure. although the actual magnetic field orientation and strength has not yet been measured for these structures.," Velocity profiles for some trunks were shown to be consistent with solid-body rotation \citep{GahCarJohEA06}, again consistent with a magnetic origin of the observed structure, although the actual magnetic field orientation and strength has not yet been measured for these structures."
+ The magnetic field in cometary globules has been measured by optical polarimetry (222).. showing in two cases a field orientation along the head-tail axis of the globule. but in one case (7). a perpendicular field was found.," The magnetic field in cometary globules has been measured by optical polarimetry \citep{SriBhaRaj96,Bha99,BhaMahMan04}, showing in two cases a field orientation along the head–tail axis of the globule, but in one case \citep{Bha99} a perpendicular field was found."
+ ? used near-IR xolarimetry of background stars to measure the magnetic field in ΝΤΟ. finding an ordered large-scale field in the H region. but within the pillars the field is aligned with the pillar axes and significantly misaligned with the ambient field by &~ 30-40," \citet{SugWatTamEA07} used near-IR polarimetry of background stars to measure the magnetic field in M16, finding an ordered large-scale field in the H region, but within the pillars the field is aligned with the pillar axes and significantly misaligned with the ambient field by $\theta \sim 30$ $40^{\circ}$."
+ They suggest this should constrain the magnetic field strength since it has not been strong enough to resist reorientation during the formation and evolution of the pillars. a suggestion we explore in more detail in this work.," They suggest this should constrain the magnetic field strength since it has not been strong enough to resist reorientation during the formation and evolution of the pillars, a suggestion we explore in more detail in this work."
+ Theoretical ealeulations of the effects of magnetie fields on the expansion of H regions were first considered by ?.., Theoretical calculations of the effects of magnetic fields on the expansion of H regions were first considered by \citet{Las66b}.
+ Using analytic calculations ?/— studied ionisation fronts with a plane-parallel magnetic field in the plane of the ionisation front.," Using analytic calculations \citet{RedWilDysEA98}
+ studied ionisation fronts with a plane-parallel magnetic field in the plane of the ionisation front."
+ They found that the velocity separation between R-type and D-type solutions decreases with increasing field strength. and that D-critical ionisation fronts also advance into the neutral gas more rapidly for increasing field strength.," They found that the velocity separation between R-type and D-type solutions decreases with increasing field strength, and that D-critical ionisation fronts also advance into the neutral gas more rapidly for increasing field strength."
+ In 2? jump conditions for ionisation fronts with oblique magnetic fields were presented. together with 1D numerical models showing how an ionisation front could progress from fast-R-type through fast-D-type. slow-R-type. and finally to slow-D-type.," In \citet{WilDysHar00} jump conditions for ionisation fronts with oblique magnetic fields were presented, together with 1D numerical models showing how an ionisation front could progress from fast-R-type through fast-D-type, slow-R-type, and finally to slow-D-type."
+ These extra modes are allowed because fast and slow shocks detach from the ionisation front at different times and propagate into the neutral medium., These extra modes are allowed because fast and slow shocks detach from the ionisation front at different times and propagate into the neutral medium.
+ This work was extended by ?.. who calculated the internal structures of stationary [D ionisation fronts.," This work was extended by \citet{WilDys01}, who calculated the internal structures of stationary 1D ionisation fronts."
+ They found that oblique fields could produce significant transverse velocities and regions where the transverse field component is signiticantly modified., They found that oblique fields could produce significant transverse velocities and regions where the transverse field component is significantly modified.
+ ? used 2D slab-symmetrie radiation-magnetohydrodynamics (R-MHD) simulations to study the photoevaporation flows from magnetised globules., \citet{Wil07} used 2D slab-symmetric radiation-magnetohydrodynamics (R-MHD) simulations to study the photoevaporation flows from magnetised globules.
+ In these models clumps with an initial density of ng=2.10em were placed in an ambient medium 10 less dense.," In these models clumps with an initial density of $n_\mathrm{H}=2\times10^5\,\mathrm{cm}^{-3}$ were placed in an ambient medium $10\times$ less dense."
+ For simplicity a uniform magnetic field with various strengths and orientations was used., For simplicity a uniform magnetic field with various strengths and orientations was used.
+ Plane-parallel radiation was assumed. with a thermal model where the temperature relaxed to a value between LOOK and 10000K according to its ionisation fraction.," Plane-parallel radiation was assumed, with a thermal model where the temperature relaxed to a value between $100\,$ K and $10\,000\,$ K according to its ionisation fraction."
+ It was found that for a weak field the ionised gas pressure dominates the dynamics and the tield was swept into a configuration where it was parallel to the column of neutral gas behind the dense clump., It was found that for a weak field the ionised gas pressure dominates the dynamics and the field was swept into a configuration where it was parallel to the column of neutral gas behind the dense clump.
+" For a sufficiently strong field. however. ye field determined the dynamics and made the flow almost one-""imensional along field lines."," For a sufficiently strong field, however, the field determined the dynamics and made the flow almost one-dimensional along field lines."
+ The first 3D R-MHD calculation including non-equilibrium shotoionising radiative transfer was performed by ?.., The first 3D R-MHD calculation including non-equilibrium photoionising radiative transfer was performed by \citet*{KruStoGar07}.
+ They used ye code witha new ray-tracing module to simulate the expansion of an H region around a point source in a uniform magnetic field., They used the code witha new ray-tracing module to simulate the expansion of an H region around a point source in a uniform magnetic field.
+ The overall expansion is now axisymmetrie rather jan spherically symmetric (cf.2). with a dense shell forming in ocirections parallel to the field. and a possibly numerical instability eveloping for expansion perpendicular to the field.," The overall expansion is now axisymmetric rather than spherically symmetric \citep[cf.][]{Las66a} with a dense shell forming in directions parallel to the field, and a possibly numerical instability developing for expansion perpendicular to the field."
+" Using similar methods. ?. model the photoionisation of a ""ense clump of gas in 3D with an initially uniform magnetic field."," Using similar methods, \citet{HenArtDeCEA09} model the photoionisation of a dense clump of gas in 3D with an initially uniform magnetic field."
+ They use the photon-conserving C7-ray ray-tracing algorithm (2) which allows large timesteps to be taken without loss of accuracy., They use the photon-conserving $^2$ -ray ray-tracing algorithm \citep{MelIliAlvEA06} which allows large timesteps to be taken without loss of accuracy.
+ They found that the evolution of a photoionised globule can be significantly altered by the presence of a strong field. and in some cases a recombining shell forms at the termination shock of the photoevaporation flow when it is confined by a transverse field.," They found that the evolution of a photoionised globule can be significantly altered by the presence of a strong field, and in some cases a recombining shell forms at the termination shock of the photoevaporation flow when it is confined by a transverse field."
+ The implosion phase is strongly asymmetric since the clump compresses much more readily along tield lines than across them., The implosion phase is strongly asymmetric since the clump compresses much more readily along field lines than across them.
+ These authors introduce a detailed thermal model to model the dynamics as realistically as possible for conditions in the Orion Nebula. finding that X-ray heating keeps the neutral gas temperature at 750K. but the cooling in neutral gas is somewhat stronger than that considered in MLIO.," These authors introduce a detailed thermal model to model the dynamics as realistically as possible for conditions in the Orion Nebula, finding that X-ray heating keeps the neutral gas temperature at $T\gtrsim50\,$ K, but the cooling in neutral gas is somewhat stronger than that considered in ML10."
+ In this work we add magnetic fields of various strengths and orientations to some of the models considered in MLIO to study their effects on the dynamics of the dense neutral gas., In this work we add magnetic fields of various strengths and orientations to some of the models considered in ML10 to study their effects on the dynamics of the dense neutral gas.
+ The simulation code is described in refsec:numerics:: refssec:MHD describes the MHD dynamics algorithm: refssec:MP reviews the microphysical processes included: presents a brief comparison of results obtained with our code to the test problem described by ?.., The simulation code is described in \\ref{sec:numerics}: \\ref{ssec:MHD} describes the MHD dynamics algorithm; \\ref{ssec:MP} reviews the microphysical processes included; \\ref{ssec:KSG07} presents a brief comparison of results obtained with our code to the test problem described by \citet{KruStoGar07}.
+ The 3D simulations presented here are introduced in refsec:simulations and our results presented in refsec:results.., The 3D simulations presented here are introduced in \\ref{sec:simulations} and our results presented in \\ref{sec:results}.
+ refssec:projections shows the evolution of the projected gas density and magnetic field orientation: refssec:emission shows the emission from ionised gas: refssec:RSR8a evaluates. boundary effects in strong field simulations., \\ref{ssec:projections} shows the evolution of the projected gas density and magnetic field orientation; \\ref{ssec:emission} shows the emission from ionised gas; \\ref{ssec:R8R8a} evaluates boundary effects in strong field simulations.
+ A calculation explaining the presence of a bright ionised linear ridge/ribbon in the strong field simulations is presented in refsec:BarFormation.., A calculation explaining the presence of a bright ionised linear ridge/ribbon in the strong field simulations is presented in \\ref{sec:BarFormation}.
+ These results are compared to observations of H regions and their magnetic fields in refseerdiscussion.. where we also discuss our work in the context of other recent computational research.," These results are compared to observations of H regions and their magnetic fields in \\ref{sec:discussion}, where we also discuss our work in the context of other recent computational research."
+ Our conclusions are presented ins refsee:conclusions.. and some technical details and tests for the simulations are given in the appendices.," Our conclusions are presented in \\ref{sec:conclusions}, , and some technical details and tests for the simulations are given in the appendices."
+ These calculations are performed on a uniform finite-volume grid. with the hydrodynamics and ray-tracing/microphysics modules as previously described in MLIO.," These calculations are performed on a uniform finite-volume grid, with the hydrodynamics and ray-tracing/microphysics modules as previously described in ML10."
+ The simulations presented here use, The simulations presented here use
+and in the Alagellanie Clouds (see. e.&. Dattinelli&De-mers(2004):Groenewegenetal. (2000))).,"and in the Magellanic Clouds (see e.g. \citet{batti, gro09}) )."
+ The analyses bv Croenewegen&deJong(1993) and Alarigoοἳal. (1999).. based on svnthetic PP-ACGB models. indicate that the extent of the EDU required to reproduce the luminosity function of the MCs should be larger than predicted by the standard models. where TDU is found in the context of the Schwarzschilel criterion.," The analyses by \citet{gro93} and \citet{paola03}, based on synthetic TP-AGB models, indicate that the extent of the TDU required to reproduce the luminosity function of the MCs should be larger than predicted by the standard models, where TDU is found in the context of the Schwarzschild criterion."
+ With respect to mass loss along the AGB. several prescriptions. based on either theoretical or empirical erounds. have been proposed in the literature.," With respect to mass loss along the AGB, several prescriptions, based on either theoretical or empirical grounds, have been proposed in the literature."
+ Among them we recall: Reimers’ classical law (1975) was a fit of measured niass-loss rates as a function of AJAL: Vassiliadis&Wood(1003) (hereinafter. VW93) was based. on the observed correlation between mass-loss rates ancl periods of pulsating AGB stars: Blocker(1995). proposed. a miocification to the teimers’ law based on fitting models from Bowen(LOSS): Bowen&Willson(1991). presented a prescription. based on dynamical models of pulsating oxygen-rich atmospheres including the formation of dust: basing on those models Dedijn(1988). derived a mass-Ioss formalism as a function of basic stellar parameters. AZ. #2 and. Zia: Wachteret(2002.2008) proposed. a formula based. on dust-driven superwind models of C stars: vanLoonetal.(2005) presented. an empirical calibration based. on observations of dust-enshrouded: red supergiants ancl oxvgen-rich. ACD stars: Schródeor&Cuntz(2005). introduced a correction to the Reimers’ formula. based. on some physical arguments: Stranieroetal.(2006) (hereinafter. SOG) formulated. a revised. calibration of the mass. loss-period. relation for pulsating AGB stars.," Among them we recall: Reimers' classical law (1975) was a fit of measured mass-loss rates as a function of $M R/L$; \citet{VW93} (hereinafter VW93) was based on the observed correlation between mass-loss rates and periods of pulsating AGB stars; \citet{blo} proposed a modification to the Reimers' law based on fitting models from \citet{bowen88}; \citet{bowen91} presented a prescription based on dynamical models of pulsating oxygen-rich atmospheres including the formation of dust; basing on those models \citet{bedijn88} derived a mass-loss formalism as a function of basic stellar parameters, $M$, $R$ , and $T_{\rm eff}$; \citet{wachter02, wachter08}
+ proposed a formula based on dust-driven superwind models of C stars; \citet{jacco} presented an empirical calibration based on observations of dust-enshrouded red supergiants and oxygen-rich AGB stars; \citet{schr05} introduced a correction to the Reimers' formula based on some physical arguments; \citet{oscar2} (hereinafter S06) formulated a revised calibration of the mass loss-period relation for pulsating AGB stars."
+ In AGB models. an additional source of uncertainty is related to the adopted low-P radiative opacities. which quantify the removal of radiation energy as photons pass through the L-rich mantle.," In AGB models, an additional source of uncertainty is related to the adopted low-T radiative opacities, which quantify the removal of radiation energy as photons pass through the H-rich mantle."
+ The carly predictions by Marigo(2002) were confirmed. by Cristalloetal.(2008)... who showed how the physical ancl chemical evolution. of low-niass. low-Z stars are allected by the adoption of the correct Opacity treatment.," The early predictions by \citet{paola02} were confirmed by \citet{sergio}, who showed how the physical and chemical evolution of low-mass, low-Z stars are affected by the adoption of the correct opacity treatment."
+ The models presented. in. this paper were calculated by means of the ALTON code for stellar evolution (Alazzitelli 1979).. a full description of which can be found in Ventura (1998).," The models presented in this paper were calculated by means of the ATON code for stellar evolution \citep{italo}, a full description of which can be found in \citet{paolo1}."
+ Convection was modelled according to the Full Spectrum of Purbulence (hereinafter EST) prescription by Canuto&Mazzitelli (1991)., Convection was modelled according to the Full Spectrum of Turbulence (hereinafter FST) prescription by \citet{canuto}.
+. Mixing of chemicals and nuclear burning were coupled by means of a diffusive approach. Following the scheme presented in Cloutman&Koll(1976): accordingly. convective overshoot was modelled by an exponential decay of convective velocities bevond the formal borders. with an e-lding decay of/=Hp.," Mixing of chemicals and nuclear burning were coupled by means of a diffusive approach, following the scheme presented in \citet{clout}; accordingly, convective overshoot was modelled by an exponential decay of convective velocities beyond the formal borders, with an e-folding decay of $l=\zeta H_P$."
+ During the two main phases of core nuclear burning and in occurrence of the second clredge-up we used ὁ=0.02. in agreement with the calibration based on the width of the main sequences of open clusters given in Venturaetal.(1998).," During the two main phases of core nuclear burning and in occurrence of the second dredge-up we used $\zeta=0.02$, in agreement with the calibration based on the width of the main sequences of open clusters given in \citet{paolo1}."
+. During the TP-ACXGD phase. given the uncertainties outlined in the previous section. no overshoot was considered.," During the TP-AGB phase, given the uncertainties outlined in the previous section, no overshoot was considered."
+ ‘To investigate the sensitivity of the results obtained on he mass loss prescription. our choice was to compare the indings obtained with two cilferent prescriptions. namely Blocker(1995) and Stranieroetal.(2006).," To investigate the sensitivity of the results obtained on the mass loss prescription, our choice was to compare the findings obtained with two different prescriptions, namely \citet{blo} and \citet{oscar2}."
+. Lo the former case it is assumed. based on hvedrocdsnamical simulations BowenLOSS)... a steep increase of mass loss with luminosity as the star enters the ACB phase.," In the former case it is assumed, based on hydrodynamical simulations \citep{bowen88}, a steep increase of mass loss with luminosity as the star enters the AGB phase."
+ This is simulated: by multiplying the canonical Reimers’ formula by a luminosity )»xver (L)., This is simulated by multiplying the canonical Reimers' formula by a luminosity power $^{2.7}$ ).
+" ""Phe free parameter. entering the Reiner’s oescription is gg=0.02. according to the calibration of he luminosity function of lithiume-rich stars observed in the Magellanic Clouds given in Venturactal.(2000)."," The free parameter entering the Reimer's prescription is $\eta_R=0.02$, according to the calibration of the luminosity function of lithium-rich stars observed in the Magellanic Clouds given in \citet{paolo2}."
+. In Stranieroetal.(2006). the mass-loss rate is made vary with the pulsation period in the fundamental mocle: compared to VW93 the empirical revision bv SOG (see their figure 5). based on a compilation of more recent data. predicts larger mass-loss rates for pulsation periods in the range 100900 davs. whilst lower rates are associated to periods larger than 500 cays.," In \citet{oscar2} the mass-loss rate is made vary with the pulsation period in the fundamental mode: compared to VW93 the empirical revision by S06 (see their figure 5), based on a compilation of more recent data, predicts larger mass-loss rates for pulsation periods in the range $100-300$ days, whilst lower rates are associated to periods larger than 500 days."
+ For the equation of state (EOS) we adopted: the latest version of the OPAL EOS (2005) where available. superseded by the Saumonctal.(1995) EOS in the partial ionization regime.," For the equation of state (EOS) we adopted the latest version of the OPAL EOS (2005) where available, superseded by the \citet{saumon} EOS in the partial ionization regime."
+ The EOS is extended to the high-density. high-tempcrature regime according to the description given in Stoltzmann&Blocker(2000).," The EOS is extended to the high-density, high-temperature regime according to the description given in \citet{blo3}."
+. The radiative opacities were calculated following the OPAL prescription. according to Iglesias&Rogers(1996).," The radiative opacities were calculated following the OPAL prescription, according to \citet{iglesias}."
+ The conductive opacities were taken from Poteckin, The conductive opacities were taken from Poteckin.
+ As to the low-temperature opacities. ie. 1500Ix 100001. we used the same large data set. of tables as in Ventura&Alarigo(2009)... which were computed with the ESOPUS tool (Alarigo Avinger to consistently follow the significant changes in the CNO surface abundances caused by the TDU and UBB during the TP-ACD evolution.," As to the low-temperature opacities, i.e. $1\,500\,{\rm K} \le T \le 10\,000\,{\rm K}$ , we used the same large data set of tables as in \citet{paolo09}, which were computed with the SOPUS tool (Marigo Aringer to consistently follow the significant changes in the CNO surface abundances caused by the TDU and HBB during the TP-AGB evolution."
+ The opacity set was. designed. to.account for. the complex interplay between convection and nucleosynthesis. so that variations in the C/O ratio can be driven by positive/negative changes in. both carbon. and. oxygen abundances.," The opacity set was designed toaccount for the complex interplay between convection and nucleosynthesis, so that variations in the C/O ratio can be driven by positive/negative changes in both carbon and oxygen abundances."
+ For instance. while an increment of C/O is," For instance, while an increment of C/O is"
+the S H]/Lo ratio. which can be used to distinguish shocked nebulae from photoionized gas (Fesenetal.1985).,"the [S $\alpha$ ratio, which can be used to distinguish shocked nebulae from photoionized gas \citep{fesen85:Evolved_SNRs}."
+" A related. concern is that some SNRs might be ""Lost? inside superbubbles formed by several nested SN explosions (MacLow&AMeCray1988)..", A related concern is that some SNRs might be “lost” inside superbubbles formed by several nested SN explosions \citep{maclow88:superbubbles}.
+ In. practice. SNRs evolving inside low-density cavities do not disappear they just »come fainter and evolve more slowly. at least until the Mast. wave reaches the cavity edges.," In practice, SNRs evolving inside low-density cavities do not disappear – they just become fainter and evolve more slowly, at least until the blast wave reaches the cavity edges."
+ The largest. most prominent superbubble complex in the LMC. 30. Doradus. iàrbors several well-stuclicd SNRs. including 6910 (NI57B) Clownsleyetal.2006).. ancl many other objects isted in “Table 1. are associated: with superbubbles.," The largest, most prominent superbubble complex in the LMC, 30 Doradus, harbors several well-studied SNRs, including $-$ 6910 (N157B) \citep{townsley06:30Dor}, and many other objects listed in Table 1 are associated with superbubbles."
+ As long as the surveys are sensitive enough to find faint SNRs. this should not be a major issue.," As long as the surveys are sensitive enough to find faint SNRs, this should not be a major issue."
+ In the absence of à svstematic re-analysis of all the available data with a set of homogeneous criteria to identify SNRs in both Clouds. we cannot guarantee that our sample is LOO% complete.," In the absence of a systematic re-analysis of all the available data with a set of homogeneous criteria to identify SNRs in both Clouds, we cannot guarantee that our sample is $100\%$ complete."
+ Such an analysis is outside the scope of the present work. but all the evidence seenis to indicate that our compilation should not be missing a large number of objects. and it should include few. if any. spurious ones.," Such an analysis is outside the scope of the present work, but all the evidence seems to indicate that our compilation should not be missing a large number of objects, and it should include few, if any, spurious ones."
+ We conclude that our list MC SNlts is at least. complete (which is enough for our goals). and that the paucity of SNRs below 50 mJy is probably real. ancl not due to any observational selection cllects.," We conclude that our list of MC SNRs is at least complete (which is enough for our goals), and that the paucity of SNRs below $\sim 50$ mJy is probably real, and not due to any observational selection effects."
+ This suggests that any MC SNRs with Ilux densities between the racio [oor and the detection limits must fade relatively quickly., This suggests that any MC SNRs with flux densities between the radio floor and the detection limits must fade relatively quickly.
+ As we shall discuss in 6.. the results of Paper 11. can be used to argue independently that the SNR sample we have compiled. has a high degree of completeness.," As we shall discuss in \ref{sec:disc}, the results of Paper II can be used to argue independently that the SNR sample we have compiled has a high degree of completeness."
+ ligure 3. shows the cumulative and dillerential histograms of remnant sizes in the LMC and in the SMC., Figure \ref{sizedist} shows the cumulative and differential histograms of remnant sizes in the LMC and in the SMC.
+ In. both Clouds. the cumulative distributions are roughly linear in reninant size. Le.. with an equal number of remnants. per equal size bin in the cillerential distributions. up to a cutoll at a physical radius rou30 pc.," In both Clouds, the cumulative distributions are roughly linear in remnant size, i.e., with an equal number of remnants per equal size bin in the differential distributions, up to a cutoff at a physical radius $r_{\rm cut}\sim 30$ pc."
+ Because the data points in a cumulative histogram are not independent of one another. these curves cannot be fitted in the usual was (using the x statistic).," Because the data points in a cumulative histogram are not independent of one another, these curves cannot be fitted in the usual way (using the $\chi^{2}$ statistic)."
+ A robust. quantitative estimate for the slope of the distributions can be obtained by performing instead a maximum-likelihood fit. as follows (seeAlaoz&Rix1993.forasimilartreatmentiopliedtoadillerent. problem)..," A robust, quantitative estimate for the slope of the distributions can be obtained by performing instead a maximum-likelihood fit, as follows \citep[see][for a similar treatment applied to a
+different problem]{maoz93:grav_lensing_statistics}."
+ Suppose that a particular model predicts a size distribution δαν= n(r). which integrates to {PounCr)dr= Nowhere N is the total number of remnants up to rou.," Suppose that a particular model predicts a size distribution $dN/dr=n(r)$ , which integrates to $\int _0^{r_{\rm cut}} n(r) dr = N$, where $N$ is the total number of remnants up to $r_{\rm cut}$."
+ Lowe bin our data into many small bins between r=0 and rau. cach of width dr. most bins will have zero SNRs. and some will have one SNR.," If we bin our data into many small bins between $r=0$ and $r_{\rm cut}$, each of width $\delta r$, most bins will have zero SNRs, and some will have one SNR."
+ Given the model. the Poisson probability of finding j remnants in the 7th bin. for whieh the model predicts n(r;) remnants. is We will define the likelihood. of a given. model. as the product. of these probabilities.," Given the model, the Poisson probability of finding $j$ remnants in the $i$ th bin, for which the model predicts $n(r_i)$ remnants, is We will define the likelihood of a given model as the product of these probabilities."
+ The logarithm of the likelihood. considering that j abways equals either 0 or 1. simplifies to where the summation is only over the specilie data values of the SNI radii.," The logarithm of the likelihood, considering that $j$ always equals either 0 or 1, simplifies to where the summation is only over the specific data values of the SNR radii."
+ A power law of index à. having the proper nornialization. will have the form Inserting in Iq. 2..," A power law of index $\alpha$, having the proper normalization, will have the form Inserting in Eq. \ref{loglike},"
+ dillerentiating In with respect to o. and equating to zero to find the maximum gives the maxiniun likelihood solution. with an uncertainty on à of This procedure vields a maximum likelihood index of a=014+0.18 for the LMC. and a=0.32+0.28 for the SAIC.," differentiating $\ln L$ with respect to $\alpha$, and equating to zero to find the maximum gives the maximum likelihood solution, with an uncertainty on $\alpha$ of This procedure yields a maximum likelihood index of $\alpha=0.14\pm 0.18$ for the LMC, and $\alpha=0.32\pm 0.28$ for the SMC."
+ Thus. the LAIC appears to indeed have a SNR size distribution that is close to uniform.," Thus, the LMC appears to indeed have a SNR size distribution that is close to uniform."
+ In the SMC. the best fit is intermediate between a flat. distribution and one that rises linearly with radius. but given the smaller number of SNRs. it is consistent with both slopes.," In the SMC, the best fit is intermediate between a flat distribution and one that rises linearly with radius, but given the smaller number of SNRs, it is consistent with both slopes."
+ From Figure 3.. it appears that the steeper slope in the SMC is driven by the small-radius side of the distribution.," From Figure \ref{sizedist}, it appears that the steeper slope in the SMC is driven by the small-radius side of the distribution."
+ Indeed. if we fit separately the first S points and the following 14 points. the maximun likelihood solution is à=1.7d1.0 at small raclii. and à=0.17£0.23 thereafter.," Indeed, if we fit separately the first 8 points and the following 14 points, the maximum likelihood solution is $\alpha=1.7\pm 1.0$ at small radii, and $\alpha=0.17\pm 0.23$ thereafter."
+ This result. confirms the visual impression. but formally it is still consistent with a slope close to zero at all radii at the 1.77 level.," This result confirms the visual impression, but formally it is still consistent with a slope close to zero at all radii at the $1.7\sigma$ level."
+ We conclude that the SNR size distribution in both Cloucls is consistent with being roughly uniform. although there are hints for a deviation at small raciii in the SMC.," We conclude that the SNR size distribution in both Clouds is consistent with being roughly uniform, although there are hints for a deviation at small radii in the SMC."
+ Lt is possible that a deficit at small racii is also present in the LMC distribution (sce Figure 3)). but the Poisson errors are too large to claim that the data require it.," It is possible that a deficit at small radii is also present in the LMC distribution (see Figure \ref{sizedist}) ), but the Poisson errors are too large to claim that the data require it."
+ The linear cumulative distribution of SN. sizes in the ALC's (and also the Milky. Way) has been previously noted and discussed by Mathewsonetal.(1984)... (1984).. Cureen (1984).. Hughesetal. (1984).. Pusco-Femiano&Preite-Martinez (1984).. Berkhuijsen (1987).. nicutt. (1988).. ancl most. recently by Bandiera&Petruk (2010)..," The linear cumulative distribution of SNR sizes in the MCs (and also the Milky Way) has been previously noted and discussed by \citet{mathewson84:SNR-Magellanic-Clouds}, \citet{mills84:Radio_SNRs}, \citet{green84:SNR_Statistics}, \citet{hughes84:SNR_N_D}, \citet{fusco-femiano84:LMC_SNRs}, \citet{berkhuijsen87:SNRs_N_D}, \citet{chu88:LMC_SNRs_environments}, and most recently by \citet{bandiera10:Statistics_SNRs}."
+ Several of these papers also pointed: out eutolfs in the distribution., Several of these papers also pointed out cutoffs in the distribution.
+ ALL of these papers considered. smaller SNR samples. often based. on much shallower radio data and. with few exceptions. did not include multiwavelength observations.," All of these papers considered smaller SNR samples, often based on much shallower radio data and, with few exceptions, did not include multi-wavelength observations."
+ Some of these authors interpreted the observed size distribution as evidence that most MC SNRs are in their “free expansion” phase. during which the shock velocity is constant. and inferred that these SNRs expand into an extremely low-censity αποπα.," Some of these authors interpreted the observed size distribution as evidence that most MC SNRs are in their “free expansion” phase, during which the shock velocity is constant, and inferred that these SNRs expand into an extremely low-density medium."
+ Alternatively. Green(1984) and Hughesetal.(1984) warned that the observed distribution was the result of selection elfects: most objects they discussed. hack been selected. in. X-rays. and the X-rav Hux limits then led to the exclusion. of Larger. and fainter remnants. and their faint radio counterparts.," Alternatively, \citet{green84:SNR_Statistics} and \citet{hughes84:SNR_N_D} warned that the observed distribution was the result of selection effects; most objects they discussed had been selected in X-rays, and the X-ray flux limits then led to the exclusion of larger and fainter remnants, and their faint radio counterparts."
+ Our present compilation is bigger. it incorporates the most recent multi-wavelength data. it extends to largersizes. and most importantly. it is sensitive to racio Hux densities two orders of magnitude below the observed. luminosity floor.," Our present compilation is bigger, it incorporates the most recent multi-wavelength data, it extends to largersizes, and most importantly, it is sensitive to radio flux densities two orders of magnitude below the observed luminosity floor."
+ With these data we now confirm the luminosity Door. the uniform size distribution. and the cutolf at rau~30 pe in the Magellanic Cloud SNRs.," With these data we now confirm the luminosity floor, the uniform size distribution, and the cutoff at $r_{\rm cut} \sim 30$ pc in the Magellanic Cloud SNRs."
+owing to feedback ellects could harden the source spectrum. within the ionization bubbles. leading to potentially higher temperatures than we have found here.,"owing to feedback effects could harden the source spectrum within the ionization bubbles, leading to potentially higher temperatures than we have found here."
+ Although there is πο straightforward way within the scope of our semi-analvtic work to directly reproduce local fi. and densitv-feedback ellects from numerical simulations. we mention these caveats and note that we able to reproduce the results of 7. by lowering the N-rav/UV boundary to 30 eV or placing the I-fronts closer to the source. both of which elfectivelv harden the local ionizing spectra.," Although there is no straightforward way within the scope of our semi-analytic work to directly reproduce local $f_{\rm esc}$ and density-feedback effects from numerical simulations, we mention these caveats and note that we able to reproduce the results of \citet{chen08} by lowering the X-ray/UV boundary to 30 eV or placing the I-fronts closer to the source, both of which effectively harden the local ionizing spectra."
+" Last. 2.. 7.. and ? have a ""ully cosmological calculation that keeps track of the evolving spectra of stellar and QSO populations."," Last, \citet{ricotti05}, \citet{chen08}, and \citet{Thomas:08} have a fully cosmological calculation that keeps track of the evolving spectra of stellar and QSO populations."
+ Therefore. soft. X-ravs from high-z sources are recshilteck and can become important for hard UV ionization at. later epochs.," Therefore, soft X-rays from $z$ sources are redshifted and can become important for hard UV ionization at later epochs."
+ Our current results do not factor in this effect. but we plan to extend this work in the near future to fully cosmological calculations that include realistic galaxy profiles anc the redshift evolution of galaxy halos and their radiation fields.," Our current results do not factor in this effect, but we plan to extend this work in the near future to fully cosmological calculations that include realistic galaxy profiles and the redshift evolution of galaxy halos and their radiation fields."
+ We have examined the cllects of X-rays from high-redshift quasars and stars when acting in combination with hard UV ionizing radiation from these sources., We have examined the effects of X-rays from high-redshift quasars and stars when acting in combination with hard UV ionizing radiation from these sources.
+ We Lined that. relative to hard UV radiation. X-rays may not dominate the ionization and thermal history of the IGM. and contribute modest. increases to the LGAL ionization at. z LO and contribute of order. 10 101 to he ICM. temperatures.," We find that, relative to hard UV radiation, X-rays may not dominate the ionization and thermal history of the IGM, and contribute modest increases to the IGM ionization at $z \sim$ 10 and contribute of order $10^3$ $10^5$ K to the IGM temperatures."
+" This is in. contrast with some earlier. works in which X-ravs could cause IGM heating up o 10"" Ix and. near-total reionization at 2=1020.", This is in contrast with some earlier works in which X-rays could cause IGM heating up to $^6$ K and near-total reionization at $z$ =10–20.
+ While some of this may be due to our simplified mocdels. we believe that most of the dilference between our results (where we include the X-rays coming from incividual sources). and those of other works deriving high ICM temperatures and ionization from X-rays at z = 1020. arise from the Iatter's assumption of strong X-ray production that is tied to the star formation rate at high τοΚΙ.," While some of this may be due to our simplified models, we believe that most of the difference between our results (where we include the X-rays coming from individual sources), and those of other works deriving high IGM temperatures and ionization from X-rays at $z$ = 10–20, arise from the latter's assumption of strong X-ray production that is tied to the star formation rate at high redshift."
+ We also examined the 21 em signatures of various cases involving combinations of stars and black hole masses. and ind that the 21 cm signal of X-ray versus UV ionization could be distinct. resulting from their dillering contributions o the topology of reionization.," We also examined the 21 cm signatures of various cases involving combinations of stars and black hole masses, and find that the 21 cm signal of X-ray versus UV ionization could be distinct, resulting from their differing contributions to the topology of reionization."
+ We find that the brightness empoerature emission expected. from. X-rays alone occur at smaller scales than that from UV radiation., We find that the brightness temperature emission expected from X-rays alone occur at smaller scales than that from UV radiation.
+" Phe clillerent spatial scales at which they manifest may therefore ""blur he 21 em signature of the percolation of reionization around carly halos. depending on whether a cosmic X-ray or UV ickeround: built up first."," The different spatial scales at which they manifest may therefore “blur"" the 21 cm signature of the percolation of reionization around early halos, depending on whether a cosmic X-ray or UV background built up first."
+ An N-rav. background. may not significantly. precede a UV. background. as à typical X-ray xhotoionization timescale exceeds the Hubble time for z 10.," An X-ray background may not significantly precede a UV background, as a typical X-ray photoionization timescale exceeds the Hubble time for $z \ga$ 10."
+ From our simpilied treatment. it is unclear whether there isa cosmological epoch when the IGMs thermodsnamic and ionization properties are determined mostly by. X-rays.," From our simpified treatment, it is unclear whether there is a cosmological epoch when the IGM's thermodynamic and ionization properties are determined mostly by X-rays."
+ The role of X-rays versus hard UV. radiation can also be tested through their interactions with the CMD. where the relative strengths of their contributions to reionization as well as the redshifts that they dominantly. contribute at can. be constrained through the €MD polarization power spectrum at large angular scales.," The role of X-rays versus hard UV radiation can also be tested through their interactions with the CMB, where the relative strengths of their contributions to reionization as well as the redshifts that they dominantly contribute at can be constrained through the CMB polarization power spectrum at large angular scales."
+ Phe currently operating all-sky CMD mission Planck max be able to distinguish such scenarios., The currently operating all-sky CMB mission $Planck$ may be able to distinguish such scenarios.
+ For sullicientlv hard raciation from sources. the LI 11 and Le LEE fronts maw lic very close to each other.," For sufficiently hard radiation from sources, the H II and He III I-fronts may lie very close to each other."
+ Although our calculation is 1D in nature. this result will impact the escape fraction of ionizing radiation from. primordial ealaxies. and the geometry of bubbles and. chimneys as ionization proceeds from these galaxies.," Although our calculation is 1D in nature, this result will impact the escape fraction of ionizing radiation from primordial galaxies, and the geometry of bubbles and chimneys as ionization proceeds from these galaxies."
+ We hope to examine this problem in a future work., We hope to examine this problem in a future work.
+ To further explore the evolution of first-light. sources and the LGAL with recdshift. we will extend our current calculations using the Galacticus code to a Cully cosmological framework that includes evolving dark matter halos. galaxy/BIL formation.uz and evolving stellar/QSO populations with time-dependent radiation fields.," To further explore the evolution of first-light sources and the IGM with redshift, we will extend our current calculations using the Galacticus code to a fully cosmological framework that includes evolving dark matter halos, galaxy/BH formation, and evolving stellar/QSO populations with time-dependent radiation fields."
+ This will permit a self-consistent. caleulation of IGM. reionization. ancl allow us to derive predictions for the growth. and evolution of a cosmologically representative clistribution of ionizecL bubbles as a function. of redshift.," This will permit a self-consistent calculation of IGM reionization, and allow us to derive predictions for the growth and evolution of a cosmologically representative distribution of ionized bubbles as a function of redshift."
+ We can also calculate the bubbles’ thermal properties. as well as the statistical properties of the bubble population. such as the mean size of ionized. and neutral regions and. power spectra of 21 cm emission. or absorption relative to the CAMB (utilizing the known correlation properties. of the dark matter halos which host the sources).," We can also calculate the bubbles' thermal properties, as well as the statistical properties of the bubble population, such as the mean size of ionized and neutral regions and power spectra of 21 cm emission or absorption relative to the CMB (utilizing the known correlation properties of the dark matter halos which host the sources)."
+ Such predictions be tested by data from CALB2 space telescopes: such asPlanck and grouncd-basecl radio telescopes that are designed to map the percolation of reionization around first-light) sources.," Such predictions will be tested by data from CMB space telescopes such as, and ground-based radio telescopes that are designed to map the percolation of reionization around first-light sources."
+ These. observations. coupled: with our etailed theoretical predictions. will additionally place strong constraints on the populaions of ionizing sources at intermediate to high. redshifts and. therefore. on. the properties of carly ecnerations of ealaxics and AGN.," These observations, coupled with our detailed theoretical predictions, will additionally place strong constraints on the populations of ionizing sources at intermediate to high redshifts and, therefore, on the properties of early generations of galaxies and AGN."
+ The resulting improvenients in our understanding of these carly objects will permit more robust predictions to be made for other observing programs. such as those of theTelescope... which will probe similar galaxy/QSO populations.," The resulting improvements in our understanding of these early objects will permit more robust predictions to be made for other observing programs, such as those of the, which will probe similar galaxy/QSO populations."
+ AV gratefully acknowledges support [rom Research Corporation through the Single. Investigator Cottrell College Science. Aware. ancl from the University of San Francisco Faculty Development Fund.," AV gratefully acknowledges support from Research Corporation through the Single Investigator Cottrell College Science Award, and from the University of San Francisco Faculty Development Fund."
+ AJB acknowledges the support of the Gordon Betty Moore Foundation., AJB acknowledges the support of the Gordon Betty Moore Foundation.
+ We thank the referee for a constructive report. ancl Massinio ]ücotti. Nucler Chen and Steve Purlanetto for useful input.," We thank the referee for a constructive report, and Massimo Ricotti, Xuelei Chen and Steve Furlanetto for useful input."
+When this fit is extrapolated to lower energies. a large soft excess becomes evident.,"When this fit is extrapolated to lower energies, a large soft excess becomes evident."
+ We model this by adding a (redshiftecl) black body. and a local photoclectric absorption o account both for the Galactic one and. any possible accditional absorption in the source (note that we expect his to be ~107em.> based on the Jalmer decrement of he BLR).," We model this by adding a (redshifted) black body, and a local photoelectric absorption to account both for the Galactic one and any possible additional absorption in the source (note that we expect this to be $\sim 10^{22}\, {\rm cm}^{-2}$ based on the Balmer decrement of the BLR)."
+ The fit results in a 47=439.93 for 413 degrees of freedom ancl a probability of the model not being able o describe the data of only83%., The fit results in a $\chi^2=439.93$ for 413 degrees of freedom and a probability of the model not being able to describe the data of only.
+. We tried to fit a plasma emission mocel (IBavimoncd-Smith) instead of the black boc. »ut no wood Lit could be achieved.," We tried to fit a plasma emission model (Raymond-Smith) instead of the black body, but no good fit could be achieved."
+ The rather high black body temperature eV) suggests a low black hole mass Alen~ ÀA£.. which is consistent with a luminosity of —3ergs rom such a black hole radiating at near the Iddington limit.," The rather high black body temperature $kT_{BB}\sim 140$ eV) suggests a low black hole mass $M_{BH}\sim 10^4-10^5\, M_{\odot}$ , which is consistent with a luminosity of $\sim 3\times 10^{43}\, {\rm erg}\,
+{\rm s}^{-1}$ from such a black hole radiating at near the Eddington limit."
+" It is remarkable that the [fi o the X-rav. data does not require any photoclectric absorption in excess to the ealactic value (Nowe~1.36«107""em. 7)."," It is remarkable that the fit to the X-ray data does not require any photoelectric absorption in excess to the galactic value $N_{Gal}\sim
+1.36\times 10^{20}\, {\rm cm}^{-2}$ )."
+ In fact we have frozen Ny to this value ancl added: an extra. photoelectric absorption Component at the redshift of the target 2= 0.0653., In fact we have frozen $N_H$ to this value and added an extra photoelectric absorption component at the redshift of the target $z=0.0653$ .
+ lte-tting the data to this model finds its best value at no intrinsic absorption with a 3 sigma upper limit of 14«107em7.," Re-fitting the data to this model finds its best value at no intrinsic absorption with a 3 sigma upper limit of $1.4 \times 10^{20}\, {\rm cm}^{-2}$."
+ This value is 70 times smaller than the minimum predicted. from the BL )Jalmer decrement interpreted. in terms of dust. reddening and for a normal eas to dust ratio., This value is 70 times smaller than the minimum predicted from the BLR Balmer decrement interpreted in terms of dust reddening and for a normal gas to dust ratio.
+ Ht is even 7 times smaller than the value implied by the NLR reddening. but this might just be due to small covering factor of an internally reddened NLR.," It is even 7 times smaller than the value implied by the NLR reddening, but this might just be due to small covering factor of an internally reddened NLR."
+ The existence. of an ionised absorber is ruled. out. by the N-rav data., The existence of an ionised absorber is ruled out by the X-ray data.
+ “Phe addition of a multiplicative model with a fixed column density of Αμ~107em does not improve the v," The addition of a multiplicative model with a fixed column density of $N_{H}\sim 10^{22}\, {\rm cm}^{-2}$ does not improve the $\chi^2$."
+ Only very [iuge values of the ionisation parameter £1000 or very high temperatures of the absorber (KZ~LOT IX) can reach a us as good as (but not smaller than) the one without an lonisect absorber (see fig 6)).," Only very large values of the ionisation parameter $\xi>1000$ or very high temperatures of the absorber $kT \sim 10^8
+\, {\rm K}$ ) can reach a $\chi^2$ as good as (but not smaller than) the one without an ionised absorber (see fig \ref{noabsori}) )."
+ Fie 7. shows the resulting EPLC-pn spectrum alter the overall fit was performed. together with the contributions to the X7 from each channel.," Fig \ref{Xray-spec} shows the resulting EPIC-pn spectrum after the overall fit was performed, together with the contributions to the $\chi^2$ from each channel."
+ There. are no resicluals at the energies of the most prominent absorption edges., There are no residuals at the energies of the most prominent absorption edges.
+ The existence of a slight positive residual at the hard energy end is uncertain as the background. level is very high at these energles., The existence of a slight positive residual at the hard energy end is uncertain as the background level is very high at these energies.
+ ‘There is à negative residual at around. 0.7 keV in the data (see Fie 7))., There is a negative residual at around 0.7 keV in the data (see Fig \ref{Xray-spec}) ).
+ The feature extends for 0.1-0.2 keV with an amplitude of 10 per cent., The feature extends for 0.1-0.2 keV with an amplitude of $\sim 10$ per cent.
+ At the moment. individual calibration of the IEEPIC-pn instrument gives residuals of ~3 per cent or less. bu joint. calibration of all ΧΑΛΙΛΝοτο instruments leaves large discrepancies between IEPIC pn and EPLC MOS in the energy range 0.3-1.3 keV of up ο 15 per cent.," At the moment, individual calibration of the EPIC-pn instrument gives residuals of $\sim 3$ per cent or less, but joint calibration of all XMM-Newton instruments leaves large discrepancies between EPIC pn and EPIC MOS in the energy range 0.3-1.3 keV of up to 15 per cent."
+ Besides that. a negative residual similar to ours is seen in many EPIC-pn spectra at around. 0.7. keV. (S. Sembay. private communication) and therefore we believe it to bea calibration artifact.," Besides that, a negative residual similar to ours is seen in many EPIC-pn spectra at around 0.7 keV (S. Sembay, private communication) and therefore we believe it to be a calibration artifact."
+ Fig 8 shows the model [fitted with the various components labelled., Fig \ref{Xray-model} shows the model fitted with the various components labelled.
+ Phe flux. corrected for. photoelectric absorption (assumed. galactic) of the source is ~10ereem7s5! in the 0.5-2 keV band. and 2.00101orecn7s tin the 2-10 keV. band.," The flux, corrected for photoelectric absorption (assumed galactic) of the source is $\sim 1.62\times 10^{-12}\, {\rm
+erg}\, {\rm cm}^{-2}\, {\rm s}^{-1}$ in the 0.5-2 keV band and $2.09\times 10^{-12}\, {\rm erg}\, {\rm cm}^{-2}\, {\rm s}^{-1}$ in the 2-10 keV band."
+ Phe luminosity is 3A107eres ‘in the 0.5-2 keV band and 3.91075ergs in the 2-10 keV band.," The luminosity is $3.1\times 10^{43}\, {\rm erg}\,
+{\rm s}^{-1}$ in the 0.5-2 keV band and $3.9\times 10^{43}\, {\rm
+erg}\, {\rm s}^{-1}$ in the 2-10 keV band."
+ About of the 0.5-2 keV luminosity (7o10eres T) is contributed by the soft excess that we," About of the 0.5-2 keV luminosity $\sim 7\times 10^{42}\, {\rm erg}\, {\rm s}^{-1}$ ) is contributed by the soft excess that we"
+is absent.,is absent.
+ Other spatial derivative ternis decay as e?2725 for the expected AVTD limits of the variables as 75»x |Z: so that an AVTD sineularity is consistent., Other spatial derivative terms decay as $e^{\Lambda -2 \tau - 2z}$ for the expected AVTD limits of the variables as $\tau \to \infty$ \cite{ber-bkbvm2}: so that an AVTD singularity is consistent.
+ Fig., Fig.
+" 6 shows log,44|V| vs 7 for typical spatial points.", \ref{ber-fig6} shows $\log_{10}|V|$ vs $\tau$ for typical spatial points.
+ Thus we see the expected exponetial decay., Thus we see the expected exponetial decay.
+ In Fie. 7..," In Fig. \ref{ber-fig7},"
+ we see that the maxi values of the chanecs with time of the quautitites expected to be constant in an AVTD regine also decay expoucutially as expected., we see that the maximum values of the changes with time of the quantitites expected to be constant in an AVTD regime also decay exponentially as expected.
+ These results will be discussed in detail elsewhere [23].., These results will be discussed in detail elsewhere \cite{ber-bkbvm2}.
+ We emphasize here that the polarized ((1) models present numerical difficulties., We emphasize here that the polarized $U(1)$ models present numerical difficulties.
+ The absence of Ve. means that spiky features do not develop., The absence of $V_{\nabla \omega}$ means that spiky features do not develop.
+ The situation unfortunately changes for generic (uupolanzed) ((1) models., The situation unfortunately changes for generic (unpolarized) $U(1)$ models.
+ It appears that consistency areuimeuts which sugecst au AVTD sineularity in polarized C(1) models and restrict e to [0.1] in the Cowes models fail in," It appears that consistency arguments which suggest an AVTD singularity in polarized $U(1)$ models and restrict $v$ to $[0,1]$ in the Gowdy models fail in"
+asvinmetric.,asymmetric.
+ However. due to the stability. of the lower end of the cistributions. the mode is consistent. between ~1.10Hem and 410em for all models considered.," However, due to the stability of the lower end of the distributions, the mode is consistent, between $\sim 1\times 10^{14}cm$ and $4\times 10^{14}cm$ for all models considered."
+ ]t is apparent from Fig., It is apparent from Fig.
+ 4 that the assumptions for photometric error and. the direction. of the transverse velocity. are important factors. particularly for a larger assumed mean microlens mass.," \ref{diff} that the assumptions for photometric error and the direction of the transverse velocity are important factors, particularly for a larger assumed mean microlens mass."
+ Like the assumptions for microlens mass. assumptions for the error and transverse velocity. direction οσο the measurement of transverse velocity. anc hence source size.," Like the assumptions for microlens mass, assumptions for the error and transverse velocity direction effect the measurement of transverse velocity, and hence source size."
+ Upper limits are more sensitive than lower limits which are set. primarily by the stellar velocity dispersion., Upper limits are more sensitive than lower limits which are set primarily by the stellar velocity dispersion.
+ These elfects are also enhanced by the presence of a component of smooth matter., These effects are also enhanced by the presence of a component of smooth matter.
+ The co-dependence of mean microlens mass and source size is now treated more rigorously., The co-dependence of mean microlens mass and source size is now treated more rigorously.
+" In Sec 3.2 the ποιασος probability distributions. passim). that the continuum source size is between S4 and 9,|dS, given that the mean microlens mass is (mn were obtained."," In Sec \ref{source_size} the normalised probability distributions, $p_s(S_{\rmn s} | \langle m \rangle)$, that the continuum source size is between $S_{\rmn s}$ and $S_{\rmn s} + {\rmn d}S_{\rmn s}$ given that the mean microlens mass is $\langle m \rangle$, were obtained."
+ The probability that the source size lies in the above range and the mean microlens mass is between m) and n3|dom) is then given by where pun) is the probability that the mean microlens mass is between Gn) and Gra}|clones., The probability that the source size lies in the above range and the mean microlens mass is between $\langle m \rangle$ and $\langle m \rangle + {\rmn d} \langle m \rangle$ is then given by where $p_m(\langle m \rangle)$ is the probability that the mean microlens mass is between $\langle m \rangle$ and $\langle m \rangle + {\rmn d} \langle m \rangle$.
+"pu, (oa) is a known function (given in ΑΛΑΤο). that is approximately independent of the source size assumed and was determined under the same assumptions for the Bavesian prior (Lat ane logarithmic) of galactic ellective transverse velocity as the function ο(οτανlent)."," $p_m(\langle m \rangle)$ is a known function (given in WWTc), that is approximately independent of the source size assumed and was determined under the same assumptions for the Bayesian prior (flat and logarithmic) of galactic effective transverse velocity as the function $p_{v}(v_{tran}|\langle m\rangle)$."
+ However unlike ο(ΌτανΕΣ). Pu(or) is dependent on the prior assumed (see WW'Tc).," However unlike $p_{v}(v_{tran}|\langle m\rangle)$, $p_m (\langle m \rangle)$ is dependent on the prior assumed (see WWTc)."
+ Figure 5. shows contour plots of the two-dimensional distribution pass.) For the assumed. source orientations. smooth matter. fraction ancl photometric errors.," Figure \ref{contour} shows contour plots of the two-dimensional distribution $p_{m,s}(S_s,\langle m\rangle)$ for the assumed source orientations, smooth matter fraction and photometric errors."
+ ALL possible combinations of the two values of each parameter are included. so that eight models are shown.," All possible combinations of the two values of each parameter are included, so that eight models are shown."
+ The contours are at 3.6. 14. 26. 61 per cent of the peak height. and so the extrema of the relevant contours represent de. 3o. 20. and lo limits on the single variables.," The contours are at 3.6, 14, 26, 61 per cent of the peak height, and so the extrema of the relevant contours represent $4\sigma$, $3 \sigma$, $2\sigma$, and $1\sigma$ limits on the single variables."
+ The light and dark lines correspond. to the assumptions of logarithmic and Hat priors for effective transverse velocity., The light and dark lines correspond to the assumptions of logarithmic and flat priors for effective transverse velocity.
+" The most ikelv model is that £m?20.05 O.2AL. (dependent on the model svstematies of trajectory direction. smooth matter content and photometric uncertainty) and S,272tot cm (approximately independent of model)."," The most likely model is that $\langle m \rangle \simeq 0.05-0.2 M_{\odot}$ (dependent on the model systematics of trajectory direction, smooth matter content and photometric uncertainty) and $S_{\rmn s} \simeq 2\times10^{14}$ cm (approximately independent of model)."
+" The most important benefit of combining the distributions for 9, and (mj is that an upper limit can x: placed on the mean microlens mass.", The most important benefit of combining the distributions for $S_{\rmn s}$ and $\langle m \rangle$ is that an upper limit can be placed on the mean microlens mass.
+ The limits on in are subject to both random: ancl systematic. uncertainties: 10 measured mass increases with photometric errors and ecreases with the fraction of smooth matter., The limits on $\langle m \rangle$ are subject to both random and systematic uncertainties; the measured mass increases with photometric errors and decreases with the fraction of smooth matter.
+ Xssuming that 16 chosen values for these parameters bracket. the likely models. upper ancl lower limits for 6m) can be found by using 10 most extreme values of the contour limits.," Assuming that the chosen values for these parameters bracket the likely models, upper and lower limits for $\langle m \rangle$ can be found by using the most extreme values of the contour limits."
+ Vhis vields the 95 per cent result that QOLAL.Sim?X1M..., This yields the 95 per cent result that $0.01 M_{\odot} \la \langle m \rangle \la 1 M_{\odot}$.
+ The upper imit. which agrees with that obtained [rom simultaneous consideration of probabilities for £m)? and Vio; (WAVΤο) is ower than that obtained by Lewis Irwin (1996). although sulliciently high that it does not conlliet with any popular models of galactic halos or bulges.," The upper limit, which agrees with that obtained from simultaneous consideration of probabilities for $\langle m\rangle$ and $V_{tran}$ (WWTc) is lower than that obtained by Lewis Irwin (1996), although sufficiently high that it does not conflict with any popular models of galactic halos or bulges."
+" Similarly. the contour maxima viele the result that there is a 95 per cent chance of 5.107 em xiS,X107 enm."," Similarly, the contour maxima yield the result that there is a 95 per cent chance of $5 \times 10^{12}$ cm $\la S_{\rmn s} \la 10^{15}$ cm."
+ At a given confidence level. both the upper and lower limits of source size. and the lower limit of mean microlens mass have smaller values in Fig.," At a given confidence level, both the upper and lower limits of source size, and the lower limit of mean microlens mass have smaller values in Fig."
+ 5 than those obtained [from probabilities for single variables in previous sections anc WWTe., \ref{contour} than those obtained from probabilities for single variables in previous sections and WWTc.
+ This can be attributed to the asymmetry of the distribution in both dimensions., This can be attributed to the asymmetry of the distribution in both dimensions.
+ The observed. peak in the Dight-curve. for. image (X of Q2237|0305 in LOSS appears to be one half. of a double peaked event that is characteristic of a source having passed, The observed peak in the light-curve for image A of Q2237+0305 in 1988 appears to be one half of a double peaked event that is characteristic of a source having passed
+within which the cluster deusitv is 200 times the can density required to close the universe).,within which the cluster density is 200 times the mean density required to close the universe).
+ This sale region has been mapped in J. IL aud Is. using the WIRC iustimment on the Palomar in telescope.," This same region has been mapped in J, H, and $_{s}$ using the WIRC instrument on the Palomar $\,$ in telescope."
+ We also take advantage of spectra from SDSS DR? which we add to our own data from WYN., We also take advantage of spectra from SDSS DR7 which we add to our own data from WIYN.
+ The Near-IR magnitudes are driven bw old stars. which comprise the bulk of the ealaxv lnass and are therefore paramount in our later derivation of the galaxy stellar mass.," The Near-IR magnitudes are driven by old stars, which comprise the bulk of the galaxy mass and are therefore paramount in our later derivation of the galaxy stellar mass."
+ The Near-IR photometry was gathered at the Palomar 200i telescope using the Widefield Iufahled Camera (AVIRC) during Έπος nights of observations between April 23-25. 2008.," The Near-IR photometry was gathered at the Palomar 200in telescope using the Widefield InfraRed Camera (WIRC) during three nights of observations between April 23-25, 2008."
+" As detailed in Table 1. the nights were clear with seeing between 0.9-1.3""7 and most of the region covered bySpitzer was duaged. as shown iu Figure 1.."," As detailed in Table 1, the nights were clear with seeing between $^{\prime\prime}$ and most of the region covered by was imaged, as shown in Figure \ref{spitzer_cov}."
+ Fields in J were observed iu sequences of [Os. and for 308 in Π. Iu ky. 3 coadds of 13s cach were taken to avoid saturation.," Fields in J were observed in sequences of 40s, and for 30s in H. In $_{s}$, 3 coadds of 13s each were taken to avoid saturation."
+" The total iutegration times for J. IT. aud IX, are 250.0. 187.0. and 217.0 minutes. respectively."," The total integration times for J, H, and $_{s}$ are 250.0, 187.0, and 247.0 minutes, respectively."
+ Dark frames were also obtained for cach iutegration time., Dark frames were also obtained for each integration time.
+ The same basic reduction technique was applied to all three bands using the pipeline of Ton Jarvet5, The same basic reduction technique was applied to all three bands using the pipeline of Tom Jarret.
+5 A απού dark frame was subtracted from the data frames. aud correction terius for the flux non-linearity are applied.," A median dark frame was subtracted from the data frames, and correction terms for the flux non-linearity are applied."
+ A inediau sky nade of a niaxiuun of LO frames was calculaed aud subtracted from the data frame before flat fielding., A median sky made of a maximum of 10 frames was calculated and subtracted from the data frame before flat fielding.
+ By accounting for the rotationa offset of he telescope using a list of known stars from the Two Micron All Sky Survey (Skrutskic et al., By accounting for the rotational offset of the telescope using a list of known stars from the Two Micron All Sky Survey (Skrutskie et al.
+ 2006: rereatter. 2ALASS) the astrometry was neasured and verified by eve.," 2006; hereafter, 2MASS) the astrometry was measured and verified by eye."
+ The backeround leve ancl fiux das corrections were then subtracted fro he rales., The background level and flux bias corrections were then subtracted from the frames.
+αρ}. was used to mosaic the frames ogether., \citep{ber07} was used to mosaic the frames together.
+ A final flux calibration ou the mosaic. relative to the 2\TASS catalog. was performed.," A final flux calibration on the mosaic, relative to the 2MASS catalog, was performed."
+" The Near-IR photometry reaches a depth of J19.5. II—18.5. aud Ik,~ 17.) with uucertaiuties of 1.9. 2.7. and respectively."," \nocite{skr06}
+ The Near-IR photometry reaches a depth of $\sim$ 19.5, $\sim$ 18.5, and $_{s}\sim$ 17.0 with uncertainties of 1.9, 2.7, and, respectively."
+ Figure 2 demonstrates how this is roughly two maeutiles deeper than those available with 2\LASS., Figure \ref{jhk2mass} demonstrates how this is roughly two magnitudes deeper than those available with 2MASS.
+" SExtractor (7) was usec to ideutifv sources and measure their magnituOs,", SExtractor \citep{ber96} was used to identify sources and measure their magnitudes.
+ To Clsüuro Wwe extracted both simall and oxended objects; we used aμοι of 5 aud of 2.0.," To ensure we extracted both small and extended objects, we used a of 5 and of 2.0."
+ To eusure we found the sinaller objects that ofteu surround larger objects (ex., To ensure we found the smaller objects that often surround larger objects (ex.
+ near the brightest cluster galaxy) we used aiinincont of 0.00005., near the brightest cluster galaxy) we used a of 0.00005.
+ We have collected archival photometry frou theSpitzer Space Telescope (PIs: Fazio. Ricke. Witavama).," We have collected archival photometry from the Space Telescope (PIs: Fazio, Rieke, Kitayama)."
+ The data were taken in 2003. 9001. aud 2007 aud all sets were observed for at least 1.5 hours (see Table 13).," The data were taken in 2003, 2004, and 2007 and all sets were observed for at least 1.8 hours (see Table \ref{obs}) )."
+ The MIPS. instiiuneut provides Far-Infrared images at 2tpau. TO;," The MIPS instrument provides Far-Infrared images at $\mu$ m, $\mu$ m"
+LI (8) measurements which will be presented later (Fig. 4)),LRG $w(\theta )$ measurements which will be presented later (Fig. \ref{fig:2SLAQ_LRG_3sur}) )
+ and it will be taken into account there., and it will be taken into account there.
+ Finally. Fig.," Finally, Fig."
+ 3. shows our results when we cross-correlate PQZ QSOs with spectroscopic 25LAQ LRGs., \ref{fig:2D_2qz} shows our results when we cross-correlate 2QZ QSOs with spectroscopic 2SLAQ LRGs.
+ No fibre collision elfect is expected in this case since 2072 QSOs had higher priority than 2dGIU galaxies for spectroscopic observations., No fibre collision effect is expected in this case since 2QZ QSOs had higher priority than 2dFGRS galaxies for spectroscopic observations.
+ Thus there is no issue of fibre incompleteness (see also. ee. Myers et al.," Thus there is no issue of fibre incompleteness (see also, e.g., Myers et al."
+ 2003)., 2003).
+ The reason for performing these measurements is to see if we can detect any signal in the case of the overlappec redshift range. (triangles)., The reason for performing these measurements is to see if we can detect any signal in the case of the overlapped redshift range (triangles).
+ Although the results whenwe use 20Z QSOs in 1.0;X2.2 range (open circles) seem to give a null average signal as expected. in the case of the common QSO-LRG redshift range (triangles) the sunples are. small and. the results appear too noisy to draw anv statistically significant conclusions.," Although the results whenwe use 2QZ QSOs in $1.0\leq z\leq 2.2$ range (open circles) seem to give a null average signal as expected, in the case of the common QSO-LRG redshift range (triangles) the samples are small and the results appear too noisy to draw any statistically significant conclusions."
+ 3efore we proceed to measure the redshilt-space cross-correlation function. we shall repeat our w(@) measurements. Cross-correlating our spectroscopic QSO samples from SDSS. 207 and 28LACO with the photometric LRG samples. which were presented in Section 2.2 and Table 2..," Before we proceed to measure the redshift-space cross-correlation function, we shall repeat our $w(\theta )$ measurements, cross-correlating our spectroscopic QSO samples from SDSS, 2QZ and 2SLAQ with the photometric LRG samples, which were presented in Section 2.2 and Table \ref{Table:photo_qso}."
+ The purpose of these measurements is to see how the LltGs correlate with bright and faint QSOs and then use Limber's formula to convert these 2-D. measurements into 3-D. real-space measurements and compare them with our results from the spectroscopic samples (Section 4)., The purpose of these measurements is to see how the LRGs correlate with bright and faint QSOs and then use Limber's formula to convert these 2-D measurements into 3-D real-space measurements and compare them with our results from the spectroscopic samples (Section 4).
+ The results using 25LAQ photometric LRCs are shown in Fie., The results using 2SLAQ photometric LRGs are shown in Fig.
+ 4 (filled circles are corrected for fibre collisions. fron 0.1!/«8100. Le. including the bump in Fig. 2))," \ref{fig:2SLAQ_LRG_3sur} (filled circles are corrected for fibre collisions, from $0.1'<\theta<100'$, i.e. including the bump in Fig. \ref{fig:2D_2slaq}) )"
+ and the results using the ANOmeea photometric LRGs are shown in Fie. 5.., and the results using the AAOmega photometric LRGs are shown in Fig. \ref{fig:AAO_LRG_3sur}.
+ We should note that neither the AAOmega nor the 28LACQ results have been corrected for stellar contamination of the LRGs (215'4 in AXOmoesga and z5'4 in 28LAQJ., We should note that neither the AAOmega nor the 2SLAQ results have been corrected for stellar contamination of the LRGs $\simeq15\%$ in AAOmega and $\approx5\%$ in 2SLAQ).
+ In both cases. triangles show the results using SDSS QSOs. filled circles using 2SLAQ QSOs and open circles using 2QZ QSOs.," In both cases, triangles show the results using SDSS QSOs, filled circles using 2SLAQ QSOs and open circles using 2QZ QSOs."
+" The 6,4 and ry values from the fits (in the range O.1’- 100 ) to these measurements are also shown.", The $\theta _0$ and $r_0$ values from the fits (in the range $'$ $'$ ) to these measurements are also shown.
+ Both results show a slightly steeper slope for the brightest QSO sample (ολο. = O.S+0.1) compared to the other QSO samples.," Both results show a slightly steeper slope for the brightest QSO sample (SDSS, $\gamma =-0.8\pm0.1$ ) compared to the other QSO samples."
+ Comparing our (0) results. from 320 QSO-spectroscopic LRGs (triangles. of Fig. 3)), Comparing our $w( \theta)$ results from 2QZ QSO-spectroscopic LRGs (triangles of Fig. \ref{fig:2D_2qz}) )
+ with those rom 2QZ QSO-photometrie LRGs (open circles of Fig. 4)), with those from 2QZ QSO-photometric LRGs (open circles of Fig. \ref{fig:2SLAQ_LRG_3sur}) )
+ =ve note that the latter shows a positive signal whereas 1ο former shows no signal. at least on small scales. ancl appears to be noisy.," we note that the latter shows a positive signal whereas the former shows no signal, at least on small scales, and appears to be noisy."
+ This is probably because the 2QZ QSO sample is much larger (z3: ) in the photometric LRG area (Table 2)) than it is in the spectroscopic LRG area CLable 1)., This is probably because the 2QZ QSO sample is much larger $\approx3\times$ ) in the photometric LRG area (Table \ref{Table:photo_qso}) ) than it is in the spectroscopic LRG area (Table \ref{Table:spec_qso}) ).
+ In this Section we present our results for the QSO-LBCG redshift-space cross-correlation function. £(s). for cillerent ο and LRG samples.," In this Section we present our results for the QSO-LRG redshift-space cross-correlation function, $\xi(s)$, for different QSO and LRG samples."
+ Fig., Fig.
+ 6 shows the results using our 25LAQ QSO-28L (spectroscopic) LG samples., \ref{fig:xis_2slaq} shows the results using our 2SLAQ QSO-2SLAQ (spectroscopic) LRG samples.
+ Our 25LAO QSO sample consistsAQ of QSOs with the same average redshift as the 28LAOQ LIGs. be. 0.35xz 0.75.," Our 2SLAQ QSO sample consists of QSOs with the same average redshift as the 2SLAQ LRGs, i.e. $0.35\leq z\leq 0.75$ ."
+ Filled, Filled
+more direct Incasurciments for the nine objects of Allen et al.,more direct measurements for the nine objects of Allen et al.
+ for the aualvsis that will follow in Sect. L., for the analysis that will follow in Sect. \ref{discussion}.
+ The accretion power Pp=WALne is deduced. assuming a steady. spherically saunietric model for accretion outo a gravitational body (Boucdi1952).," The accretion power $P_B = \eta \dot M_B c^2$ is deduced assuming a steady, spherically symmetric model for accretion onto a gravitational body \citep{bondi52}."
+".. For the sake of simplcity we have assumed an efficiency 4j 1l: the accretion rate. Mp. is defined at the Boudi radius rp=2GALIpye (ος Is the sound speed. Mp the mass of the SMDBII) as where pp. sp. aud Tp are the mass and ΠΙΟ deusities. aud the temperature at rp. respectively,"," For the sake of simplicity we have assumed an efficiency $\eta =
+1$ ; the accretion rate, $\dot M_B$, is defined at the Bondi radius $r_B = 2 G
+M_{BH}/c^2_s$ $c_s$ is the sound speed, $M_{BH}$ the mass of the SMBH) as where $\rho_B$, $n_B$, and $T_B$ are the mass and number densities, and the temperature at $r_B$, respectively."
+ The mass of the SMDII has been estimated frou the stellar velocity dispersion. available for all sources from the IIvperLbeda database. using the Tremaineetal.(2002) relation.," The mass of the SMBH has been estimated from the stellar velocity dispersion, available for all sources from the HyperLeda database, using the \citet{tremaine02} relation."
+ Iu three CoreG of the Virgo sample. the mass of the SMDBII has been measured directly aud we adopted the vaue reported by Marconi&IInut(2003).," In three CoreG of the Virgo sample, the mass of the SMBH has been measured directly and we adopted the value reported by \citet{marconi03}."
+. The temperature and deusitv were deduced from. the properties of the hot ISM. through Chandra ταν observatious., The temperature and density were deduced from the properties of the hot ISM through Chandra X-ray observations.
+ The data analysis aud the results of these observations are prescuted in detail iu the appendix. where we derive the deprojected profiles of temperature and cleasity across the coronae for the four VCC sources.," The data analysis and the results of these observations are presented in detail in the appendix, where we derive the deprojected profiles of temperature and density across the coronae for the four VCC sources."
+ The temperature profiles across the coronal regions are essentially constant. with typical values of Z2:0.1.0.7 keV. For the caleulation of the Dondiaccretion rate we assumed that Tp is equal to the temperature at the innermost shell.," The temperature profiles across the coronal regions are essentially constant, with typical values of $T \approx 0.4 - 0.7$ keV. For the calculation of the Bondiaccretion rate we assumed that $T_B$ is equal to the temperature at the innermost shell."
+" The value of 5p was extrapolated down to rp assunius a relation p(r)Xr"". with a deduced from a power-law fit to the density profile of tle ISM (see the appendix for details)."," The value of $n_B$ was extrapolated down to $r_B$ assuming a relation $n(r) \propto r^{-\alpha}$, with $\alpha$ deduced from a power-law fit to the density profile of the ISM (see the appendix for details)."
+ The values of these quantities. with the derived accretionand jet powers. are reported iu Table 2.. which also includes the data for the other Virgo objects from," The values of these quantities, with the derived accretionand jet powers, are reported in Table \ref{tab2}, , which also includes the data for the other Virgo objects from"
+For the total raciated X-rav energv. {ων we find a correlation of (Fig.,For the total radiated X-ray energy $E_X$ we find a correlation of (Fig.
+" 1 bottom right). which is expected to scale as Ex(1,)DpοpixT, (Eq."," 1 bottom right), which is expected to scale as $E_X(T_p) \propto EM_p(T_p) \tau_f(T_p)
+\propto T_p^\delta$ (Eq."
+ 7) with 4.541.326.3. which indeed agrees with the best fit. ὁ=6.1x0.5.," 7) with $\delta = \alpha + \beta = 4.5 + 1.8 = 6.3$ , which indeed agrees with the best fit, $\delta = 6.1 \pm 0.5$."
+ We cannot caleulate the thermal [lare energv. £p (Eq., We cannot calculate the thermal flare energy $E_T$ (Eq.
+" 10) with current. data. since we need the knowledge of the flare peak density πρ. which is generally not independently measured from EM, and Tj."," 10) with current data, since we need the knowledge of the flare peak density $n_p$, which is generally not independently measured from $EM_p$ and $T_p$."
+" The values quoted in Table 4 of Giiddel (2004) have a mean value of n,=101775 * and show no temperature dependence.", The values quoted in Table 4 of Güddel (2004) have a mean value of $n_p = 10^{11.5 \pm 0.6}$ $^{-3}$ and show no temperature dependence.
+" The flare peak density n,Pp can be caleulated with relation (Ec.", The flare peak density $n_p$ can be calculated with relation (Eq.
+" 10). if one uses some theoretical assumptions (e.g.. RTV scaling law. see 84.1) which predicts (he temperature dependence of (he electron density ny(1,)."," 10), if one uses some theoretical assumptions (e.g., RTV scaling law, see 4.1) which predicts the temperature dependence of the electron density $n_p(T_p)$."
+" In the EM,—T, scaling law (Eq.", In the $EM_p-T_p$ scaling law (Eq.
+" 11). the observables represent the peak emission measure EM,=EM(I—Hl.T—T,) and peak temperature T,=T(!/,,) measured at the peak time /—/, of stellar flares."," 11), the observables represent the peak emission measure $EM_p=EM(t=t_p,T=T_p)$ and peak temperature $T_p=T(t=t_p)$ measured at the peak time $t=t_p$ of stellar flares."
+ HLowever. since both observables EAM(1) and T(/) change during the flare as a function of time. we inquire how closely the temperature 74 predicted by the statistical scaling law (Ex.," However, since both observables $EM(t)$ and $T(t)$ change during the flare as a function of time, we inquire how closely the temperature $T_S$ predicted by the statistical scaling law (Eq."
+" 11). matches the observed peak temperature T,, αἱ the peak time /—/,. and how the peak temperature T,, and the statistically predicted peak temperature Z5 matches the maximum temperature 73; of observed stellar Πάνος,"," 11), matches the observed peak temperature $T_p$ at the peak time $t=t_p$, and how the peak temperature $T_p$ and the statistically predicted peak temperature $T_S$ matches the maximum temperature $T_M$ of observed stellar flares."
+ For this purpose we collected evolutionary. pliase diagrams of the temperature 7(/) versus (he emission measure £M(/) observed in 8 different flares (Fig., For this purpose we collected evolutionary phase diagrams of the temperature $T(t)$ versus the emission measure $EM(t)$ observed in 8 different flares (Fig.
+ 2)., 2).
+ The curves of 5 stellar flares observed with are Irom (he stus 10090 (or V7LL Tau] (Stern 1996). 1EQIS39.64-8002 (Pan οἱ al.," The curves of 5 stellar flares observed with are from the stars HR1099 [or V711 Tau] (Stern 1996), 1EQ1839.6+8002 (Pan et al."
+ 1997). H Peg (Doyle et al.," 1997), II Peg (Doyle et al."
+ 1991). UN Avi CEsuru οἱ al.," 1991), UX Ari (Tsuru et al."
+ 1989). ancl Algol (Stern et al.," 1989), and Algol (Stern et al."
+ 1992): (wo observations of AB Dor flares are observed withBeppoSAX (Maggio et al., 1992); two observations of AB Dor flares are observed with (Maggio et al.
+ 2000). and one observation of Proxima Centauri was obtained with (Reale et al.," 2000), and one observation of Proxima Centauri was obtained with (Reale et al."
+ 2003)., 2003).
+ The evolutionary curves are shown in Fig., The evolutionary curves are shown in Fig.
+ 2. including error bars in temperature. and some are shown with confidence regions for both derived (EM and T) parameters.," 2, including error bars in temperature, and some are shown with confidence regions for both derived (EM and T) parameters."
+" The peak emission measures EM, are indicated with", The peak emission measures $EM_p$ are indicated with
+The amounts of the πο clements Li. Be. aud D (LiBeB) at present are thought to be the παπι of the products of the bie baug nucleosvuthesis and the subsequent cosiunic-ray spallation reactions.,"	The amounts of the light elements Li, Be, and B (LiBeB) at present are thought to be the sum of the products of the big bang nucleosynthesis and the subsequent cosmic-ray spallation reactions."
+ The contribution from the big bane nucleosvuthesis to ‘Li is thought to be most prouounced among these three elements., The contribution from the big bang nucleosynthesis to $^7$ Li is thought to be most pronounced among these three elements.
+" Though some fractions of Li have been depleted Inside some cool stars. its abundances on the surfaces of metal-poor stars are the only information that we can use to identify the contribution from the big bane uucleosvuthesis (ee.Ryan.Norris,&Beers1990)."," Though some fractions of Li have been depleted inside some cool stars, its abundances on the surfaces of metal-poor stars are the only information that we can use to identify the contribution from the big bang nucleosynthesis \citep[e.g.,][]{Ryan_99}."
+ Thus it is important to kuow the contribution from the cosluic-ray spallation reactions to Li in the carly stages of the Galaxy evolution during which metal-poor stars were formed., Thus it is important to know the contribution from the cosmic-ray spallation reactions to Li in the early stages of the Galaxy evolution during which metal-poor stars were formed.
+ With this knowledge. we cau constrain the cosmological paraincters to svuthesize light οσοι. iu 1ο big hang nucleosvuthesis (Ryanctal.2000).," With this knowledge, we can constrain the cosmological parameters to synthesize light elements in the big bang nucleosynthesis \citep{Ryan_00}."
+. Iu addition. the imvestigation of the evolution of these enients iu the carly stages of the CGalaxw enables us o understand the origin of cosmüc-ravs responsible for ie liebt. clement uucleosvuthesis.," 	 	In addition, the investigation of the evolution of these elements in the early stages of the Galaxy enables us to understand the origin of cosmic-rays responsible for the light element nucleosynthesis."
+ Receut observations of extremely unctal-poor stars apparcutly sugeest that the xunary (not secondary) spallation process dominated ιο light clement nucleosvuthesis in the carly Galaxy (Dunucau.Lambert.&Lemke1992)., Recent observations of extremely metal-poor stars apparently suggest that the primary (not secondary) spallation process dominated the light element nucleosynthesis in the early Galaxy \citep{Duncan_92}.
+. In other ποτ». cosnuic-ravs composed of C/O interacting with protons and/or Ue nuclei in the interstellar πουαι (ISM) have predominantly produced LiBeB. A primary wwechanisun. iu which LiBeB were produced in supernova ejecta-euriclied superbubbles. was suggested to explain the Be evolution in the carly Galaxy (Iigdon.Liugeufelter.&Ramaty1998).," In other words, cosmic-rays composed of C/O interacting with protons and/or He nuclei in the interstellar medium (ISM) have predominantly produced LiBeB. A primary mechanism, in which LiBeB were produced in supernova ejecta-enriched superbubbles, was suggested to explain the Be evolution in the early Galaxy \citep{Higdon_98}."
+. Since then. two-component models in which helt clemeuts are produced by standard Galactic cosmic-ravs aud moetal-enriched particles in superbubbles have becu investigated bv several authors (Ramatyetal.2000:Fields.Olive.Vangioni-Flhuu.&Cassé2000:SuzukiYoshii 2001).," Since then, two-component models in which light elements are produced by standard Galactic cosmic-rays and metal-enriched particles in superbubbles have been investigated by several authors \citep{Ramaty_00, Fields_00, Suzuki_01}."
+. These studies concluded that a primary imnechnauisni is necded to explain the observed abundance treuds of Be and DB with O aud Fe., These studies concluded that a primary mechanism is needed to explain the observed abundance trends of Be and B with O and Fe.
+ Recently. Suzuki&Yoshii(2001) investigated the chemical evolution of LiBeB in the carly epoch of the Galaxy using an iuhomoseuecous cheuica evolution model developed by Tsujimoto.Shigevama.&Yoshi(1999) in which star formation is assunied to be induced by superuova (SN) explosions.," Recently, \citet{Suzuki_01} investigated the chemical evolution of LiBeB in the early epoch of the Galaxy using an inhomogeneous chemical evolution model developed by \citet{Tsujimoto_99} in which star formation is assumed to be induced by supernova (SN) explosions."
+ Their mode succeeded in reproducing uot oulv the observed metallicity distribution of Galactic halo stars but also the observe abundance correlations of heavy elements., Their model succeeded in reproducing not only the observed metallicity distribution of Galactic halo stars but also the observed abundance correlations of heavy elements.
+ Suzuki&Yoshii also considered two origins of Calactic cosiic- that svuthesize ght clemenuts., \citet{Suzuki_01} also considered two origins of Galactic cosmic-rays that synthesize light elements.
+ One is from the ISN, One is from the ISM
+ ,"hence $g(a/x)<g(a_{max}/x_{max})$, and vice -versa)."
+This clearly largeasgu fo ——————puldi!). Thuswe can write p do£2 Fusnguoi4. 4 ol — fy/2 | ," This gives, after some rearrangement and utilizing \ref{eps_def}) ) Inequality \ref{q_int4}) ) must be satisfied for all choices of the scaling factor $x$ ."
+us J/ fyτω£2pO↜∣ oU Pavo )di ⊔∐↲∐↓≀↕↴⇀↸↕∐∐∐∐≀↕↴∐∪∖∖⊽≼↲≺⇂⋟∖⊽≼↲↕↽≻≀↧↴↕⋅≀↧↴∐∪∐⋟∖⊽↕∐≺∢≼↲≀↧↴↕⋅≼↲≀↧↴↥≺∐⋟∖⊽⊔⋅↕∣↽≻∏∐∪∐∪↓≯≼↲↕∐↲↕⋅≸↽↔↴∡∖↽," For smaller $x$ $x \lesssim x_{max}$ ) $\epsilon$ can be negative, but we want to show that as $x$ increases eventually $\epsilon$ must become positive."
+∐↓∏⋟∖⊽↥⋟∖⊽≀↧↴∐⋟∖⊽↓⋟⋡∖↽⊔↱≻↕⋝⊔⋟∪↕⋅≀↧↴∐↴∏↥↕⋟∖⊽∏↕↽≻↕↽≻≼↲↕⋅∣↽≻∪∏↕∐⇂≺⇂≼↲↕↽≻≼↲↕∐⇂⋟∖⊽∪∐⊔∐↲⋟∖⊽≀↧↴∐↓↕↽≻∐∐≸≟↓⋟∏∐≺∢∐∪∐≀↕↴∐≼⊔∐↖≺↽↔↴≼↲∐≼↲↕⋅≀↕↴↥∖∖⇁↕∐∪∖↽≼↲↕⋅⊣↲⋟∖⊽∐∐↓≀↕↴∩↲ further. but we can seethat therange of possibleis," Later we will choose a more restrictive distribution of positive energy to better illustrate quantum interest, but first we will show that (at least when $d=4$ ) the total amount of positive energy is strictly greater than the total negative energy that passes the observer."
+ most strongly inlluenced by y.," To do so, evaluate \ref{q_int4}) ) in the limit as $x \rightarrow 
+\infty$."
+ Ify=1 (we have a state(hat actually achieves the minimum allowedby 9(/))then the only wav (14))," In this limit for $t=a/x$ and $t=b/x$ we can accurately evaluate $g(t)$ in a Taylor series about $t=0$: There are no odd powers because of the assumed symmetry in $g$, but even if we don't require $g$ to be symmetric there will not be any $t$ term in the series because of the peak at $t=0$ (which also forces $g''(0)$ to be negative)."
+ or(15)) can be satüislied is ifr =1:Le. positive energy," Thus \ref{q_int4}) ) can be written as In the limit $x\rightarrow\infty$ , $\epsilon \rightarrow 0$ and when the dimension $d=4$, $\epsilon$ is strictly greater than $0$ for $x$ sufficiently large."
+" must znanedialely dO 1 (p) ∣∼≓≐⋅⊥∕∕⋗⋗⊥⋅ | 16)"" ⋅hme UYQO) Ey, """, In 2dimensional Minkowski space we can only conclude that $\epsilon$ is at least zero for arbitrary fluxes using the large $x$behavior of the inequality \ref{q_int4}) ).
+A new set of parton distribution functions (PDF) was constructed in the framework of a statistical approach of the nucleon |1].. which has the following characteristic - For quarks (antiquarks). the building blocks are the helicity dependent distributions q- (qz) and we deline q—q6 and Ag=q—q4q (similarly for - At the initial-lp energy. scale taken at (Q22=Γον”.-D these distributions- are given- by the sum of two ternis. a quasi Fermi-Dirac function and a helicity independent diffractive contribution. which leads (ο a universal behavior for all flavors al very low vr. - The flavor asvinmetry [or the light sea. d>4. observed in the data is built in.,"A new set of parton distribution functions (PDF) was constructed in the framework of a statistical approach of the nucleon \cite{bbs1}, which has the following characteristic - For quarks (antiquarks), the building blocks are the helicity dependent distributions $q_{\pm}$ $\bar q_{\pm}$ ) and we define $q= q_+ +q_-$ and $\Delta q = q_+ -q_-$ (similarly for - At the initial energy scale taken at $Q^2_0= 4 \mbox{GeV}^2$, these distributions are given by the sum of two terms, a quasi Fermi-Dirac function and a helicity independent diffractive contribution, which leads to a universal behavior for all flavors at very low $x$ - The flavor asymmetry for the light sea, $\bar d > \bar u$, observed in the data is built in."
+ This is clearly understood in terms of (he Pauli exclusion principle. based on the fact that the proton contains (wo 4 quarks ancl only one d - The chiral properties of QCD lead to strong relations between qancl q. For example. it is found that the well estalished result. Aw>0 implies Av>0 and similarly Ad«0 leads to Ad«0. - Concerning the gluon. the unpolarized eliou distribution is given in terms of a (quasi Bose-Einstein function. wilhparameter. but for simplicity. one assumes zero gluon polarization. AGGr.Q7)=0. at the initial energy- All unpolarized ancl polarized distributions depend uponeight [ree parameters. which were determined in 2002 (See [1])). from an NLO fit of a selected set of accurate DIS For illustration. the + light quark (antiquark) distributions are displaved on Fig.," This is clearly understood in terms of the Pauli exclusion principle, based on the fact that the proton contains two $u$ quarks and only one $d$ - The chiral properties of QCD lead to strong relations between $q$and $\bar q$ For example, it is found that the well estalished result $\Delta u >0 $ implies $\Delta 
+\bar u >0$ and similarly $\Delta d <0$ leads to $\Delta \bar d <0$ - Concerning the gluon, the unpolarized gluon distribution is given in terms of a quasi Bose-Einstein function, with, but for simplicity, one assumes zero gluon polarization, $\Delta G(x,Q_0^2)=0$, at the initial energy- All unpolarized and polarized distributions depend upon free parameters, which were determined in 2002 (See \cite{bbs1}) ), from an NLO fit of a selected set of accurate DIS For illustration, the $\pm$ light quark (antiquark) distributions are displayed on Fig."
+ 1 and we clearly notice the essential features mentioned above7., 1 and we clearly notice the essential features mentioned above.
+. More recently. new tests against experimental (unpolarized and polarized) data turned out to be very satisfactory. in particular in hacronic reactions 2.3].," More recently, new tests against experimental (unpolarized and polarized) data turned out to be very satisfactory, in particular in hadronic reactions \cite{bbs2,bbs3}."
+. The statistical approach has been extended to (he interesting5 case where the PDF have. in addition to the usual Bjorken 2 dependence. an explicit transverse momentum hp1 dependence [4] ancl (his might5 be used in future calculations with no Ay1 integration.," The statistical approach has been extended to the interesting case where the PDF have, in addition to the usual Bjorken $x$ dependence, an explicit transverse momentum $k_T$ dependence \cite{bbs4} and this might be used in future calculations with no $k_T$ integration."
+5 Concerning the strange quark aud antiquark distributions. a simplibving assumption consists 3.i». ≓⋅⋅ ⋅ ⋅ ⋅⋅⋅ ↥∪↥≀↧↴↳↽≼↲↼↦⋅⋜↜⋅∣⋮⋅≬⊋−∣⋝∶↼↦⋅⋜⋡⋅∣⋮⋅↙↽⊋−∣⋝≀↕↴∐≼⇂⋟∖⊽∐⊔∐≀↧↴↕⋅↥∡∖↽↓∪↕⋅⊔∐↲≺∢∪↕⋅↕⋅≼↲⋟∖⊽↕↽≻∪↕∐∐∐≸≟↕↽≻∪↥≀↧↴∐∠≼↲≼⇂≼∐⋟∖⊽⊔⋅↕∣↽≻⋯↕∪∐⋟∖⊽," Concerning the strange quark and antiquark distributions, a simplifying assumption consists to take $s(x,Q^2)=\bar {s}(x,Q^2)$ and similarly for the corresponding polarized distributions $\Delta s(x,Q^2)=\Delta \bar {s}(x,Q^2)$."
+" ""⋅∆↼⇁⋅≼⋝⋅≀⋮⋅∁⊋−∣⋟∶∆↼↦⋅≼⋝⋅≀⋮⋅↙↽⊋−∣⋝⋅∐∪∖∖↽≼↲∖↽≼↲↕⋅≀↧↴≺∢≀↧↴↕⋅≼↲↓∏↥≀↧↴∐≀↧↴↥⋡∖↽⊳∖⇁↕⊳∖⊽∪↓⊔∐↲≼⇂≀↧↴↥≀↧↴↥≼↲≼⇂∏⊳∖⇁∩≻⊔∐↲≺∢∪∐≺∢∐↕⊳∖⇁↕∪∐ negative ⊔⋯↴↥↼↦⋅⋖⋡⋅≀⋮⋅↙↽⊋⇉↕⋝⊋≜∕↼⋚⋖↜⋅≀⋮⋅≬⊋⋅−⋟↕⋝≀↧↴∐≺⇂⊔∐↲≺∢∪↕⋅↕⋅≼↲⋟∖⊽↕↽≻∪∐≼∐∐≸≟↕↽≻∪↥≀↕↴↕⋅↕∠≼↲≺⇂≺∐⋟∖⊽∏⋅↕∣↽≻∏∐∪∐⋟∖⊽≀↕↴↕⋅≼↲∏∐≼↲≺⇂∏≀↕↴↥⋅⋟∖⊽∐⋯↴∐ and ο]."," However a careful analysis of the data led us to the conclusion that $s(x,Q^2)\neq \bar {s}(x,Q^2)$ and the corresponding polarized distributions are unequal, small and negative \cite{bbs5}."
+ Now let us come back to the miportant prediction of the statistical approach. namely At>O0 and Ad< 0. whieh contrasts with the flavor svmametric assiümplion made. for example. in Ref. ον. ," Now let us come back to the important prediction of the statistical approach, namely $\Delta \bar u >0$ and $\Delta \bar d <0$ , which contrasts with the flavor symmetric assumption $\Delta \bar u = \Delta \bar d = \Delta s = \Delta \bar s$ made, for example, in Ref. \cite{hks}. ."
+With Chis assumption. the Ag don't," With this assumption, the $\Delta \bar q$ don't"
+seven of (he eight RQQs in the sample listed in Table 1 now have been imaged with the VLBA. so we may sav something about their global properties.,"Seven of the eight RQQs in the sample listed in Table \ref{tab:sample}
+ now have been imaged with the VLBA, so we may say something about their global properties."
+" These objects actually have a verv wide range of radio powers al 5 Gllz. as expected [from their similar flux. densities and wide range of redshifts. vet their brightness temperature values or limits are remarkably similar,"," These objects actually have a very wide range of radio powers at 5 GHz, as expected from their similar flux densities and wide range of redshifts, yet their brightness temperature values or limits are remarkably similar."
+ The actual values/limits are subject to a strong selection ellect. namely the Lact that sources with a few millijansky of compact flix detected on scales of one to a few milliareseconds inevitably must have brightness temperatures of 103 IX or greater.," The actual values/limits are subject to a strong selection effect, namely the fact that sources with a few millijansky of compact flux detected on scales of one to a few milliarcseconds inevitably must have brightness temperatures of $10^8$ K or greater."
+ Still. it is a kev result (hat essentially all the ROQQs. including those with flat ancl steep spectra. have some Components with minimum brightness temperatures of =LO” IN. (ef.," Still, it is a key result that essentially all the RQQs, including those with flat and steep spectra, have some components with minimum brightness temperatures of $\gtrsim 10^8$ K (cf."
+ Table 4))., Table \ref{tab:tb}) ).
+ This implies that all have nonthermal emitting cores. rather than being dominated by thermal Lori such as the one inferred for NGC LOGS bv Gallimore et al. (," This implies that all have nonthermal emitting cores, rather than being dominated by thermal tori such as the one inferred for NGC 1068 by Gallimore et al. ("
+1997. 2004).,"1997, 2004)."
+ Further. (he four objects that we observed at niultiple frequencies show no evidence for low-frequency. radio (urnovers. as shown in Figure 2..," Further, the four objects that we observed at multiple frequencies show no evidence for low-frequency radio turnovers, as shown in Figure \ref{fig:spect}."
+ This implies that Iree-Iree absorption al ~1.5 GlIIz is not significant in most RQQs. ancl (hus cannot easily account for flat radio spectra al eigaliertz frequencies.," This implies that free-free absorption at $\sim 1.5$ GHz is not significant in most RQQs, and thus cannot easily account for flat radio spectra at gigahertz frequencies."
+ Why do RQQs contain such compact. aud vel relatively low-liminosily (compared {ο RLQs) radio emission?," Why do RQQs contain such compact, and yet relatively low-luminosity (compared to RLQs) radio emission?"
+ A likely possibility. of course. is that (his weak emission is somehow related to the process of accretion on supermassive black holes at the centers of the quasar host galaxies. ancl to the ellicieney of radiation from that accretion process.," A likely possibility, of course, is that this weak emission is somehow related to the process of accretion on supermassive black holes at the centers of the quasar host galaxies, and to the efficiency of radiation from that accretion process."
+ A kev element to consider here is (he overall spectral enerev distributions of the RQQs., A key element to consider here is the overall spectral energy distributions of the RQQs.
+ Figure 3. plots the radio/X-rav. ratio versus (he racdio/oplical ratio lor the RQQs in our sample., Figure \ref{fig:sed} plots the radio/X-ray ratio versus the radio/optical ratio for the RQQs in our sample.
+ The racio/optical ratio. Ho. is defined as (he ratio of observed 5 Giz ancl D band flux densities. where the conversion from £ magnitude to flux densitv is given by Schmidt&Green(1933).," The radio/optical ratio, $R_O$, is defined as the ratio of observed 5 GHz and $B$ band flux densities, where the conversion from $B$ magnitude to flux density is given by \citet{sch83}."
+. We computed the raclio/X-rav ratio. Ay. converting from the observed energy. bands to keV fluxes using N(GE)xE7.," We computed the radio/X-ray ratio, $R_X$, converting from the observed energy bands to 2--10 keV fluxes using $N(E)\propto E^{-2}$."
+ This twpically reduces the X-ray fluxes by a factor of order 2: eiven (his fact and the lack of knowledge of the spectra. Jy values may be in error by factors of 2.3.," This typically reduces the X-ray fluxes by a factor of order 2; given this fact and the lack of knowledge of the spectra, $R_X$ values may be in error by factors of 2–3."
+ For the five RQQs that we observed with the VLBA. the 5-Gllz radio power [rom the unresolved milliaresecond core has been used: otherwise we use the total VLA flix density (Table 1)).," For the five RQQs that we observed with the VLBA, the 5-GHz radio power from the unresolved milliarcsecond core has been used; otherwise we use the total VLA flux density (Table \ref{tab:sample}) )."
+" Diagonal lines ave plotted. corresponding to O/X ratios of 1 and 100. where this ratio is defined as the ratio of 7/5, at ito the 210 keV fIux."," Diagonal lines are plotted corresponding to O/X ratios of 1 and 100, where this ratio is defined as the ratio of $\nu S_\nu$ at to the 2–10 keV flux."
+ The radio-loud quasars 3C 273 and 3€ 454.3 also are plotted. with the (wo points corresponding to their total 5 GllIz flux densities and the flux densities οἱ," The radio-loud quasars 3C 273 and 3C 454.3 also are plotted, with the two points corresponding to their total 5 GHz flux densities and the flux densities of"
+ , 
+"volume orders of magnitude larger than Hm"". which would enhance the chances of detection ol the kaon decay neutrinos.","volume orders of magnitude larger than $1 {\rm km}^3$, which would enhance the chances of detection of the kaon decay neutrinos."
+ The neutrinos produced from kaon decavs max be detected as a neutrino backeround. so let us estimate the this of the neutrino background roughly.," The neutrinos produced from kaon decays may be detected as a neutrino background, so let us estimate the flux of the neutrino background roughly."
+" It can be seen from Figure 3 that the thence of neutrinos €2.N(e,) at 10. eV from this example is ~ LO” GeV. If we adopt (his value. the resulting neutrino backeround level at LOMI eV is ~3xLOH GeV ? Fur to where we use the local GRB rate. 0.4 oE ft. without correction due to jet collimation (Guettaοἱal.2005)n and the age of the universe. ~LO10 ve."," It can be seen from Figure 3 that the fluence of neutrinos $\epsilon_\nu^2
+N(\epsilon_\nu)$ at $10^{17-18}$ eV from this example is $\sim$ $^{53}$ GeV. If we adopt this value, the resulting neutrino background level at $10^{17-18}$ eV is $\sim 3 \times 10^{-11}$ GeV $^{-2}$ $^{-1}$ $^{-1}$, where we use the local GRB rate, 0.4 $^{-3}$ $^{-1}$, without correction due to jet collimation \citep{guetta05} and the age of the universe, $\sim 10^{10}$ yr."
+ This background level is [ar below the detection limit expected. from 3 vears of IceCube operation., This background level is far below the detection limit expected from 3 years of IceCube operation.
+ Since our exaniple is an optimistec case for neutrino emission. a more realistic background level may become lower 1 or 2 orders of magnitude than above.," Since our example is an optimistic case for neutrino emission, a more realistic background level may become lower 1 or 2 orders of magnitude than above."
+ Dominant sources of background neutrinos in the highest energv band are not sure so far., Dominant sources of background neutrinos in the highest energy band are not sure so far.
+ If we can observe the background neutrinos al 10!*.1I eV with a highly efficient detector in the future. the characteristic (Ia spectrum predicted [roi GRB sources is distinguishable from spectra of other sources. such as GRB early-alterglows or blazars (see. e.g.. Schneider οἱ al.," If we can observe the background neutrinos at $10^{17-18}$ eV with a highly efficient detector in the future, the characteristic flat spectrum predicted from GRB sources is distinguishable from spectra of other sources, such as GRB early-afterglows or blazars (see, e.g., Schneider et al."
+ 2002)., 2002).
+" Razzaqueetal.(2004b) showed that very high energv. eamma rays [rom. à""-decay (=10!* eV) can escape trom shells without electron-positron pair creation.", \citet{raz04b} showed that very high energy gamma rays from $\pi^0$ -decay $\gtrsim 10^{17}$ eV) can escape from shells without electron-positron pair creation.
+ However. this result largely depends on the assumption of spectra in (he lower energv region.," However, this result largely depends on the assumption of spectra in the lower energy region."
+ A recent observation (Blakeetal.|2005) shows brighter optical emission than (2004b) assumed., A recent observation \citep{bla05} shows brighter optical emission than \citet{raz04b} assumed.
+ On the other hand. AT also has some decaying modes into neutral pions.," On the other hand, $K^0_{\rm L}$ also has some decaying modes into neutral pions."
+ So. there is a possibility that delaved. gamma rays come from the decays of AT.," So, there is a possibility that delayed gamma rays come from the decays of $K^0_{\rm L}$."
+ We can investigate the validity of this statement by checking whether Aj can decay alter GRB photons escape from shells., We can investigate the validity of this statement by checking whether $K^0_{\rm L}$ can decay after GRB photons escape from shells.
+" Since the mean lifetime of AT is 5.17x10 s in the kaon rest frame. the mean lifetime of AT. Ad. of energy ej in the shell rest. [rame becomes s. On the other hand. the typical dvnamical timescale of GRBs in the shell rest lrame is 3 (D/100)(/,,/30ms) s. So. we expect that gamma rays produced from kaons with energv larger than 10P(D/100) eV can escape from shells."," Since the mean lifetime of $K^0_{\rm L}$ is $\times 10^{-8}$ s in the kaon rest frame, the mean lifetime of $K^0_{\rm L}$, $\Delta t_{\rm k}$, of energy $\epsilon_{\rm k}$ in the shell rest frame becomes $\Delta t_{\rm k} = 3.3 (\epsilon_{\rm k} / 10^{16.5} {\rm eV})$ s. On the other hand, the typical dynamical timescale of GRBs in the shell rest frame is 3 $\Gamma/100$ $t_{\rm var}/30 {\rm ms}$ ) s. So, we expect that gamma rays produced from kaons with energy larger than $10^{18.5} (\Gamma/100)$ eV can escape from shells."
+ Ol course. such 2gamma ravs cannot travel more than ~| Alpe| because of electron- pair creation wilh a radio backeround5 (ProtheroeandBiermann1996).," Of course, such gamma rays cannot travel more than $\sim 1$ Mpc because of electron-positron pair creation with a radio background \citep{pro96}."
+. The seconcary enission via inverse Compton of cosmic microwave background photons may be observed as GeV-TeV |photons., The secondary emission via inverse Compton of cosmic microwave background photons may be observed as GeV-TeV photons.
+ Razzaqueletal.(2004b) concluded that the corresponding|2 time delavs are in (he range of 10-100 5 and the duration timescale may be much longer., \citet{raz04b} concluded that the corresponding time delays are in the range of 10-100 s and the duration timescale may be much longer.
+ These (imescales are niuch longer than the lifetime of AT in the observer's frame., These timescales are much longer than the lifetime of $K^0_{\rm L}$ in the observer's frame.
+" Therefore. unfortunately. the GeV-TeV secondary photons from kaons may not be distinguishable from other sources. such as z""-decay."," Therefore, unfortunately, the GeV-TeV secondary photons from kaons may not be distinguishable from other sources, such as $\pi^0$ -decay."
+ There is a possibility that failed GRBs occur near/in our galaxy., There is a possibility that failed GRBs occur near/in our galaxy.
+ As, As
+"according to (Sternbergetal.2002) where A is the characteristic overdensity of dark matter halos. p, is the mean barvonic density of the universe. iy.=yyfry ls Che concentration parameter and f, is the mass prolile as given in Nichols&Blancd-Lawthorn(2009).","according to \citep{Sternberg2002}
+ where $\Delta$ is the characteristic overdensity of dark matter halos, $\rho_u$ is the mean baryonic density of the universe, $x_{\rm vir}\equiv{}R_{\rm vir}/r_s$ is the concentration parameter and $f_m$ is the mass profile as given in \citet{Nichols2009}."
+". The characteristic overdensity is given bv (Brvan&Norman1998) and the mean barvonic density by For an Einasto prolile. the concentration parameter. νι, also referred (o as c. is In this growing potential (here is no angular dependence and therefore the angular monientuni is conserved: such that where L is the per unit mass angular momentum squared and 7j, and r, ave the peri and apogalacticon radii that would occur if the potential was static."," The characteristic overdensity is given by \citep{Bryan1998}
+ and the mean baryonic density by For an Einasto profile, the concentration parameter, $x_{\rm vir}$, also referred to as $c$, is \citep{Duffy2008}
+ In this growing potential there is no angular dependence and therefore the angular momentum is conserved; such that where $L^2$ is the per unit mass angular momentum squared and $r_p$ and $r_a$ are the peri and apogalacticon radii that would occur if the potential was static."
+ We present our results in terms of both orbit circularity aud eccentricitv in recognition of the fact that both are in common use when cliseussing CDM model., We present our results in terms of both orbit circularity and eccentricity in recognition of the fact that both are in common use when discussing CDM model.
+ We define the eccentricity of the orbit in terms of the ratio of the apogalacticon to the perigalacticon. Γρ). and the corresponding circularity to be jj=v1—€.," We define the eccentricity of the orbit in terms of the ratio of the apogalacticon to the perigalacticon, $\epsilon=(r_a-r_p)/(r_a+r_p)$ , and the corresponding circularity to be $\eta=\sqrt{1-\epsilon^2}$."
+ These parameters are not conserved throughout the orbit due to the growing potential. ancl when referred (o we use the final state of these parameters.," These parameters are not conserved throughout the orbit due to the growing potential, and when referred to we use the final state of these parameters."
+medium-passband filter at nnm. which was used to measure continuum emission.,"medium-passband filter at nm, which was used to measure continuum emission."
+ Integration times were 20mmin (He). mmin r1]. and mmin (continuum). each split into 5 dithered exposures to compensate for the gaps between the CCDs. bad pixels. and cosmics.," Integration times were min $\alpha$ ), min ]), and min (continuum), each split into 5 dithered exposures to compensate for the gaps between the CCDs, bad pixels, and cosmics."
+ Data reduction was standard. starting with bias frame subtraction and division by a flat field frame (a combination of dome and sky flats): bad pixels were masked and excluded from further processing.," Data reduction was standard, starting with bias frame subtraction and division by a flat field frame (a combination of dome and sky flats); bad pixels were masked and excluded from further processing."
+ Then mosaics were made from the single exposures as follows., Then mosaics were made from the single exposures as follows.
+" First. a “positional reference frame"" was constructed from the continuum exposures by accurately registering and combining the data for each chip separately. resulting in one average Image for each chip."," First, a “positional reference frame” was constructed from the continuum exposures by accurately registering and combining the data for each chip separately, resulting in one average image for each chip."
+ Large enough dithering steps had been chosen such that these images overlapped. allowing an accurate determination of the positior of each chip with respect to the others.," Large enough dithering steps had been chosen such that these images overlapped, allowing an accurate determination of the position of each chip with respect to the others."
+ The averaged images for each of the 8 chips were then registered. rotated. shifted. and finally averaged into one large master image which served as positional reference frame for registering the other exposures.," The averaged images for each of the 8 chips were then registered, rotated, shifted, and finally averaged into one large master image which served as positional reference frame for registering the other exposures."
+ All individual exposures through all three filters were thei registered to this positional reference frame., All individual exposures through all three filters were then registered to this positional reference frame.
+ The final mosaics were then constructed by taking the median of the rotated anc shifted single exposures to reject cosmic ray events., The final mosaics were then constructed by taking the median of the rotated and shifted single exposures to reject cosmic ray events.
+ Further processing of the Ha image was applied to show the newly discovered HH objects more clearly., Further processing of the $\alpha$ image was applied to show the newly discovered HH objects more clearly.
+ The Ha frame was lightly gaussian smoothed so that the PSF matched that of the continuum frame: the latter was then appropriately scaled and subtracted from the Ha image., The $\alpha$ frame was lightly gaussian smoothed so that the PSF matched that of the continuum frame; the latter was then appropriately scaled and subtracted from the $\alpha$ image.
+ Since the wavelength difference between two filters was quite large. different colours of stars led to many small residual images which were removed by hand.," Since the wavelength difference between two filters was quite large, different colours of stars led to many small residual images which were removed by hand."
+ Finally. a large-scale NW-SE gradient in the Ha emission was subtracted from the image to enhance the contrast.," Finally, a large-scale NW-SE gradient in the $\alpha$ emission was subtracted from the image to enhance the contrast."
+ Figure | shows a ὃν (62.7 « ppe) section of the Ha image., Figure \ref{hh46ha} shows a $\times$ $\sim$ $\times$ pc) section of the $\alpha$ image.
+" The position of the young star driving the flow (the only protostellar object known in the area) is marked by a cross,", The position of the young star driving the flow (the only protostellar object known in the area) is marked by a cross.
+ It is seen to be embedded in a globule 1)) which ts illuminated and evaporated from the north-west by the exciting stars of the Gum nebula region and several degrees ppc) to the north-west., It is seen to be embedded in a globule ) which is illuminated and evaporated from the north-west by the exciting stars of the Gum nebula region and several degrees pc) to the north-west.
+ In the south-eastern part of the image. a number of apparently connected illuminated globules. rims. and fingerlike structures can be seen: these may be the remains of a larger cloud that has been evaporated from the north-west.," In the south-eastern part of the image, a number of apparently connected illuminated globules, rims, and fingerlike structures can be seen: these may be the remains of a larger cloud that has been evaporated from the north-west."
+ The edges of these globules are outlined by rims of Ha emission., The edges of these globules are outlined by rims of $\alpha$ emission.
+ Smooth extended Ho emission is seen to fill the spaces between the globules and to the north-west., Smooth extended $\alpha$ emission is seen to fill the spaces between the globules and to the north-west.
+ The jet system consists of much more compact. knotty. and filamentary Hoa structures.," The jet system consists of much more compact, knotty, and filamentary $\alpha$ structures."
+ Superimposed on the smooth background. a number of additional compact. knotty patehes of He emission can also be seen i1 the north-easterr and western corners of the image (marked with arrows).," Superimposed on the smooth background, a number of additional compact, knotty patches of $\alpha$ emission can also be seen in the north-eastern and south-western corners of the image (marked with arrows)."
+ Omm Figure 2 shows the same field with the continuum emission and large-scale Πα gradient subtracted in order to show the compact Ha features better., 0mm Figure \ref{hh46hausm} shows the same field with the continuum emission and large-scale $\alpha$ gradient subtracted in order to show the compact $\alpha$ features better.
+ The main features are the rims around the globules whose structure suggests illumination from the north-west. with bright emission at their north-eastern tips and streamers extending to the south-east.," The main features are the rims around the globules whose structure suggests illumination from the north-west, with bright emission at their north-eastern tips and streamers extending to the south-east."
+ In addition to harbouring and à comparably-sized second globule.2.. at the eastern edge of the image (Reipurth 1983). many smaller and less pronounced features are seen. including a very narrow finger pointing to the north-west.," In addition to harbouring and a comparably-sized second globule, at the eastern edge of the image (Reipurth 1983), many smaller and less pronounced features are seen, including a very narrow finger pointing to the north-west."
+ The features in the north-east and south-west corners marked by arrows in also stand out clearly now., The features in the north-east and south-west corners marked by arrows in \\ref{hh46ha} also stand out clearly now.
+ We argue that these are Herbig-Haro type structures as follows:, We argue that these are Herbig-Haro type structures as follows:
+spectrum from models.,spectrum from models.
+ A summary of the main results is elven in Section T., A summary of the main results is given in Section 7.
+ We use a Hubble constant of Ho=100knis+Mpe+., We use a Hubble constant of $H_0 = 100~h {\rm km~s^{-1}~Mpc^{-1}}$.
+" The AbellACO catalogue of clusters of galaxies. CXboll 1958. Abell Corwin, Olowin 1989) is presently the largest. available source. of the large-scale cistribution of matter in the Universe covering the whole sky outsice the δν Way zone of avoidance."," The Abell–ACO catalogue of clusters of galaxies (Abell 1958, Abell, Corwin, Olowin 1989) is presently the largest available source of the large-scale distribution of matter in the Universe covering the whole sky outside the Milky Way zone of avoidance."
+ We use for the present study a recent compilation of measured. redshifts for these clusters by Ancdernach. Tago. Stengler-Larrea (1997).," We use for the present study a recent compilation of measured redshifts for these clusters by Andernach, Tago, Stengler-Larrea (1997)."
+ This compilation gives redshifts or a total of about 2000 AbellACO clusters (including supplementary. or S-clusters).," This compilation gives redshifts for a total of about 2000 Abell--ACO clusters (including supplementary, or S-clusters)."
+ We USc¢ the 1995 version of the compilation. omitted. all S-clusters and used only clusters with measured redshifts up lo2—0.12.," We used the 1995 version of the compilation, omitted all S-clusters and used only clusters with measured redshifts up to $z=0.12$."
+" To this sample we added: all clusters with photometric redshift estimates 2,x0.12."," To this sample we added all clusters with photometric redshift estimates $z_{est}\le
+0.12$."
+ Our ‘ull sample contains 1304 AbellACO clusters of galaxies. 869 of which have measured redshifts.," Our full sample contains 1304 Abell–ACO clusters of galaxies, 869 of which have measured redshifts."
+ We have included: clusters of richness class O in our study., We have included clusters of richness class 0 in our study.
+ About half ofall clusters in the nearby region stucie ave of this richness class and the number of objects is crucia in the present work., About half of all clusters in the nearby region studied are of this richness class and the number of objects is crucial in the present work.
+ Abell clusters of richness class 0 are X-rav ολτους and hosts of cD galaxies with extended haloes as often as clusters of higher richness., Abell clusters of richness class 0 are X-ray emitters and hosts of cD galaxies with extended haloes as often as clusters of higher richness.
+ Both facts sugges that these elusters are physical objects which can be usec to trace the large-scale structure., Both facts suggest that these clusters are physical objects which can be used to trace the large-scale structure.
+ Possible projection ellects discussed. by Sutherland (1988). Dekel CI989). aux others are not crucial for the present study as we are mostly interested in the distribution of clusters on large scales.," Possible projection effects discussed by Sutherland (1988), Dekel (1989) and others are not crucial for the present study as we are mostly interested in the distribution of clusters on large scales."
+ A small excess of cluster pairs on small separations noted. by Sutherland ancl Dekel can be considered as an adclitional selection ellect., A small excess of cluster pairs on small separations noted by Sutherland and Dekel can be considered as an additional selection effect.
+ This sample was used. in Paper L| to derive a new catalogue of superclusters anc to. study. their. spatial distribution., This sample was used in Paper I to derive a new catalogue of superclusters and to study their spatial distribution.
+ In the present. paper we use both the cluster sample and the supercluster catalogue., In the present paper we use both the cluster sample and the supercluster catalogue.
+ “Phe use of. the supercluster catalogue gives us the possibility to analyse the distribution. of clusters in) different environments., The use of the supercluster catalogue gives us the possibility to analyse the distribution of clusters in different environments.
+" Superclusters were determined. using a ""friends-ol-friends? technique with neighbourhood. radius 24Alpe.", Superclusters were determined using a “friends-of-friends” technique with neighbourhood radius 24.
+. This radius was chosen on the basis of the multiplicity function which shows that individual superclusters start to become evident at à neighbourhood radius of about 16Mpesz at radii larger than 30 ssuperclusters begin to join into huge agelomerates with dimensions exceeding the characteristic scale of. the supercluster-void. network., This radius was chosen on the basis of the multiplicity function which shows that individual superclusters start to become evident at a neighbourhood radius of about 16; at radii larger than 30 superclusters begin to join into huge agglomerates with dimensions exceeding the characteristic scale of the supercluster-void network.
+ Thus the neighbourhood radius must lie within these boundaries., Thus the neighbourhood radius must lie within these boundaries.
+ The inlluence of this radius on our results for the correlation Function shall be studied below (Section 4.3)., The influence of this radius on our results for the correlation function shall be studied below (Section 4.3).
+ In Paper L superclusters were divided. into richness classes according to their multiplicity (the number of member clusters in superclusters)., In Paper I superclusters were divided into richness classes according to their multiplicity (the number of member clusters in superclusters).
+ I0 was also shown that the overall distribution of superelusters of different: richness is rather similar: superclusters are located. in chains that form a Fairlv. regularὃν network., It was also shown that the overall distribution of superclusters of different richness is rather similar: superclusters are located in chains that form a fairly regular network.
+ The mean diameter of voids between superclusters 1s ~100Alpe., The mean diameter of voids between superclusters is $\sim 100$.
+ The skeleton of the supercluster-void network is formed bv very rich superclusters., The skeleton of the supercluster-void network is formed by very rich superclusters.
+ Poor and medium rich superclusters as well as dsolated clusters. are scattered around them. leaving void interiors empty. of rich clusters.," Poor and medium rich superclusters as well as isolated clusters are scattered around them, leaving void interiors empty of rich clusters."
+ “Phe distribution. of superclusters in void wedis depends on the supercluster richness: the mean separation between poor ancl medium rich superclusters is small ancl has a smooth clistribution whereas the separation |»ebween very Lich superclusters is much larger ancl its clistribution is peaked: over 75 of very rich superelusters are located at separations 110.150 oon opposite sides of vokS., The distribution of superclusters in void walls depends on the supercluster richness: the mean separation between poor and medium rich superclusters is small and has a smooth distribution whereas the separation between very rich superclusters is much larger and its distribution is peaked: over 75 of very rich superclusters are located at separations $110 - 150$ on opposite sides of voids.
+ This finding. motivated us to study the correlation function of clusters of οιdaxies located. in superclusters of different richness., This finding motivated us to study the correlation function of clusters of galaxies located in superclusters of different richness.
+ As in Paper Lowe divide cluster samples into populations using he supercluster richness as the parameter which determines of clusters (see Frisch 1995)., As in Paper I we divide cluster samples into populations using the supercluster richness as the parameter which determines of clusters (see Frisch 1995).
+ In contrast to Paper Lowe ¢ivide superclusters into only. two richness classes with variable richness threshold., In contrast to Paper I we divide superclusters into only two richness classes with variable richness threshold.
+ We shall use the following nomenclature of cluster samples., We shall use the following nomenclature of cluster samples.
+ The first 3 capital letters ACO denote clusters from he AbellACO catalogue (excluding S-clusters): the following capital letter indicates whether we use the sample of all clusters (A) or the sample of clusters with measure recishitts (lh) the ollowing capital letter denotes custer. samples in high-. or low-density environments. (respectively LL or L): the last number inclicrates the limiting multiplicity δι of superclusters used. to divide the sample into high- and low-density populations.," The first 3 capital letters ACO denote clusters from the Abell--ACO catalogue (excluding S-clusters); the following capital letter indicates whether we use the sample of all clusters (A) or the sample of clusters with measured redshifts (R); the following capital letter denotes cluster samples in high-, or low-density environments (respectively H or L); the last number indicates the limiting multiplicity $N_{cl}$ of superclusters used to divide the sample into high- and low-density populations."
+ Clusters belonging to superclusters with Nop members were attributed to the high-density population. ancl isolated clusters as well as clusters in superclusters withfess than Ay members to the low-clensity population.," Clusters belonging to superclusters with $N_{cl}$ members were attributed to the high-density population, and isolated clusters as well as clusters in superclusters with than $N_{cl}$ members to the low-density population."
+ lo calculate the correlation function of clusters. of galaxies we generate Poisson samples of test. particles with the same shape and selection function as the real samples., To calculate the correlation function of clusters of galaxies we generate Poisson samples of test particles with the same shape and selection function as the real samples.
+ The selection ellects depend on Galactic absorption. on the dilliculty to find. lower richness clusters at. large distances. on the decrease in the fraction of clusters with measured redshifts with distance. the dillerences in the mean density of clusters in the Abell arwl ACO catalogues. etc.," The selection effects depend on Galactic absorption, on the difficulty to find lower richness clusters at large distances, on the decrease in the fraction of clusters with measured redshifts with distance, the differences in the mean density of clusters in the Abell and ACO catalogues, etc."
+ Poisson saniples nuist be generated. with all. these ellects taken into account., Poisson samples must be generated with all these effects taken into account.
+ We have calculated the selection function as à function of two variables. the Galactic," We have calculated the selection function as a function of two variables, the Galactic"
+lines. which were artificially removed by interpolation prior to ratioing.,"lines, which were artificially removed by interpolation prior to ratioing."
+ To obtain further comparisons. we used the United Ixingdom Infrared. Telescope (UIXIRT) and its [acilitv echelle spectrograph. CGS4 to obtain spectra of three normal carbon stars. TYC 1021 and DD +17° 3325. and DD 4-237 2998. that are not expected to have winds and for which the [le line should be weak or altogether absent.," To obtain further comparisons, we used the United Kingdom Infrared Telescope (UKIRT) and its facility echelle spectrograph, CGS4 to obtain spectra of three normal carbon stars, TYC 1021 and BD $\arcdeg$ 3325, and BD $\arcdeg$ 2998, that are not expected to have winds and for which the He line should be weak or altogether absent."
+ These stars were observed on 2008 June 29. as part of the UIXIRT Service Program. with CGS54s narrow slit. which provided a resolving power of 735.000.," These stars were observed on 2008 June 29, as part of the UKIRT Service Program, with CGS4's narrow slit, which provided a resolving power of $\sim$ 35,000."
+ A further UIXIRT Service observation of the inactive RCB star XX. Cam was obtained on 2009 January 8., A further UKIRT Service observation of the inactive RCB star XX Cam was obtained on 2009 January 8.
+ All of the UIXIRT. spectra were reduced in the manner described above., All of the UKIRT spectra were reduced in the manner described above.
+ Fieure |. contains the spectra of the five IC. stars along with the comparison hot carbon star DD +17° 3325 (top) and cool carbon star TYC 1021 (bottom)., Figure 1 contains the spectra of the five HdC stars along with the comparison hot carbon star BD $\arcdeg$ 3325 (top) and cool carbon star TYC 1021 (bottom).
+ The Πας stars in Fig., The HdC stars in Fig.
+ | are arranged in order of effective temperature (see Table 1 [or values)., 1 are arranged in order of effective temperature (see Table 1 for values).
+ For most of (hese stars several wide-ranging values of 7;;; exist in the literature and thus the ordering may not be exact., For most of these stars several wide-ranging values of $T_{eff}$ exist in the literature and thus the ordering may not be exact.
+ The wavelengths of several atomic lines. including the Ie I doublet. are identified in both figures. as well as some of the numerous lines of CN.," The wavelengths of several atomic lines, including the He I doublet, are identified in both figures, as well as some of the numerous lines of CN."
+ The wavelengths and identifications are from [Πτα(1974).. Wallaceetal. (1993).. and [linkleetal.(1995).," The wavelengths and identifications are from \citet{hir74}, \citet{wal93}, and \citet{hin95}."
+. With the exception of FQ Aqr. the spectra have been shifted in wavelength to provide best overall matches to the atomic lines.," With the exception of FQ Aqr, the spectra have been shifted in wavelength to provide best overall matches to the atomic lines."
+ The derived heliocentrie radial velocities lor the newly observed stars are listed in Table 1., The derived heliocentric radial velocities for the newly observed stars are listed in Table 1.
+ Small deviations are present [rom line to line. but the overall accuracy of the fit is about 3 km +.," Small deviations are present from line to line, but the overall accuracy of the fit is about 3 km $^{-1}$."
+ The radial velocities of the HdC stars aresvstematically about 12 km 1 more negative than those in Lawson&Cottrell (1997).., The radial velocities of the HdC stars aresystematically about 12 km $^{-1}$ more negative than those in \citet{law97}. .
+ Lyνιc1019 Ly(Leg107 J1550—564. , $L_{X}/L_{Edd} \simeq 10^{-10}$ $L_{X}/L_{Edd} \simeq 10^{-2}$ $-$ 
+A supernova remnant (SNR) consists of expelled. material called Cejecta from. the explosion. and a swept-up interstellar matter.,A supernova remnant (SNR) consists of expelled material called “ejecta” from the explosion and a swept-up interstellar matter.
+ Phe. NX-rav. emission. results from the interactions of cjecta and swept-up matter., The X-ray emission results from the interactions of ejecta and swept-up matter.
+ From the X-rav emission of SNRs we may obtain valuable information about the physical properties of the ejecta. swept-up plasma. elemental abundances anc the history of the explosion.," From the X-ray emission of SNRs we may obtain valuable information about the physical properties of the ejecta, swept-up plasma, elemental abundances and the history of the explosion."
+ Shock wave mav be the blast wave associated with the stellar explosion and/or the reverse shock wave. which propagates inwards from the decelerated: blast wave ancl raises. the temperature of the stellar ejecta.," Shock wave may be the blast wave associated with the stellar explosion and/or the reverse shock wave, which propagates inwards from the decelerated blast wave and raises the temperature of the stellar ejecta."
+ Phe voung SNRs are bright in X-rays and dominated by the emission from ejecta., The young SNRs are bright in X-rays and dominated by the emission from ejecta.
+ Thus they provide fruitful information about the elements synthesized by the supernova (SN) explosions., Thus they provide fruitful information about the elements synthesized by the supernova (SN) explosions.
+ Pherefore. the observation of the voung SNRs is the best method to investigate the abundances of the elements synthesized. by the SNe (see e.g. SNIO06 (Ixovamactal.1995)... RCMSG (Bamba.Ixovama&Tomida2000).. Tvcho (Warrenetal.2005) [or the best studied. cjecta-cdominated SNRs in the X-ray band).," Therefore, the observation of the young SNRs is the best method to investigate the abundances of the elements synthesized by the SNe (see e.g. SN1006 \citep {b11}, RCW86 \citep
+{b3}, Tycho \citep {b24} for the best studied ejecta-dominated SNRs in the X-ray band)."
+ (346.6-0.2. (RA(2000)=17104105. Dee.(2000)= 40°11) is a shel-ivpe SNR. located in the Galactic plane.," G346.6-0.2 $\rmn{RA}(2000)=17^{\rmn{h}} 10^{\rmn{m}}
+19^{\rmn{s}}$, $\rmn{Dec.}~(2000)=-40\degr 11\arcmin$ ) is a shell-type SNR located in the Galactic plane."
+ lt was discovered. by Clark.Caswell&Green.(1975). in radio band having an angular size of S arcmin (Whiteoak&Green 1996)., It was discovered by \citet {b5} in radio band having an angular size of 8 arcmin \citep {b25}.
+. In N-rav band. on the other hand. 346.6- was first observed by during its Galactic plane survey (Yamaucjοἱal.2008).," In X-ray band, on the other hand, G346.6-0.2 was first observed by during its Galactic plane survey \citep {b28}."
+. Ht is shown that the size of the X-ray emiissicn [rom €1346.6-0.2 is less extended than its reported radio sructure., It is shown that the size of the X-ray emission from G346.6-0.2 is less extended than its reported radio structure.
+ Five OLICIT20 MlIZ) masers were detected toward this SNR. and they are all located along the southern edge of the remnant (Ixoraleskyetal.1998).," Five OH(1720 MHz) masers were detected toward this SNR and they are all located along the southern edge of the remnant \citep
+{b10}."
+ We are stuving the X-ray emissions [rom the ejecta-dominated SNRs., We are studying the X-ray emissions from the ejecta-dominated SNRs.
+ For this purpose we have chosen several small size Galactic SNRs. one of which is €1346.6-0.2.," For this purpose we have chosen several small size Galactic SNRs, one of which is G346.6-0.2."
+ Detailed: properies of €346.6-0.2. remained unknown so [ar because of νεvy limitecl photon statistics in the PS) data., Detailed properties of G346.6-0.2 remained unknown so far because of very limited photon statistics in the ) data.
+ is the most recent X-ray astronomical satellite (see. Aitsucdactal. (2007))) having a large collecting area :uicl low background., is the most recent X-ray astronomical satellite (see \citet {b17}) ) having a large collecting area and low background.
+" ""Therefore. it is the best instrument for observing cim and diffuse sources."," Therefore, it is the best instrument for observing dim and diffuse sources."
+ We propose here to sucly C346.6-0.2 to understand the origin of, We propose here to study G346.6-0.2 to understand the origin of
+Another class of dwarf novae are the WZ Sagittae stars (UCWWZ). which are extreme examples ol the SU UMa-type stars.,"Another class of dwarf novae are the WZ Sagittae stars (UGWZ), which are extreme examples of the SU UMa-type stars."
+ WZ See stars exhibit large-unplitude outbursts (> 6 mag) that occur less frequently than the those of the UGSU (Osaki1995))., WZ Sge stars exhibit large-amplitude outbursts $\geq$ 6 mag) that occur less frequently than the those of the UGSU \citealt{osaki}) ).
+ WZ See stars do exhibit superhuuips iu their light curves. but do not appear to have any “normal outbursts.," WZ Sge stars do exhibit superhumps in their light curves, but do not appear to have any `normal' outbursts."
+ Here we present observations of the dwarf nova identified with the ROSAT x-ray source RAS J05323L92-621755 (hereafter RN0532+62)., Here we present observations of the dwarf nova identified with the ROSAT x-ray source RXS J053234.9+624755 (hereafter RX0532+62).
+ The discovery of this star. which lies in Cameloparclalis. is described by Poyuer&Shear," The discovery of this star, which lies in Camelopardalis, is described by \citet{poy}."
+s(2006).. Bernhardetal.(2005) classified it as a U Gem-type dwarf uovae with a recurrence time scale of 133 days.," \citet{bernhard}
+ classified it as a U Gem-type dwarf novae with a recurrence time scale of 133 days."
+ The relatively frequent outbursts show that this star is not a COWZ., The relatively frequent outbursts show that this star is not a UGWZ.
+ RN0532+62 uuderweut a well-observed superoutburst iu 2005 March., RX0532+62 underwent a well-observed superoutburst in 2005 March.
+ During this outburst a superhuump period. P4. of 0.0571(2) d (82.2 mins) was reported by Touny Vaninunster (CBA Belgium)7:: Poyner&Shears(2006) independently found a similar. but less accurate. value of Py.," During this outburst a superhump period, $P_{\rm sh}$, of 0.0571(2) d (82.2 mins) was reported by Tonny Vanmunster (CBA Belgium); \citet{poy} independently found a similar, but less accurate, value of $P_{\rm
+sh}$."
+ We undertook observatious of this star to independently determine the orbital period. Por. aud allow determination of the superhump period excess. €.," We undertook observations of this star to independently determine the orbital period, $P_{\rm orb}$, and allow determination of the superhump period excess, $\epsilon$."
+ We took spectra using the MDAL Observatory 1.31 MeCraw-Hill telescope ou itt Peak during two observing runs in 2005 September aud 2006 January., We took spectra using the MDM Observatory 1.3m McGraw-Hill telescope on Kitt Peak during two observing runs in 2005 September and 2006 January.
+ We used the Mark ID spectrograph and a thinned 1021? SITe CCD detector. with a 600 line ! erism that eave spectra spanning from 1500 ito 7000 aal 2.3 |.," We used the Mark III spectrograph and a thinned $1024^2$ SITe CCD detector, with a 600 line $^{-1}$ grism that gave spectra spanning from 4500 to 7000 at 2.3 $^{-1}$."
+" In all. 106 spectra of RX05324-62 were taken with typical exposure times of 305, along with frequent comparison spectra to maintain the waveleneth calibration."," In all, 106 spectra of RX0532+62 were taken with typical exposure times of 480 s, along with frequent comparison spectra to maintain the wavelength calibration."
+ At the beeinuine aud eud of every photometric night. we observed flux staudards to determine au absolute flux calibration.," At the beginning and end of every photometric night, we observed flux standards to determine an absolute flux calibration."
+ Table 1 gives a journal of observatious., Table 1 gives a journal of observations.
+ The system was in quiescence for all the observations reported here., The system was in quiescence for all the observations reported here.
+ We also obtained direct images with the MDM 2.lm Hiltner Telescope., We also obtained direct images with the MDM 2.4m Hiltner Telescope.
+ A SITe 20187 CCD with an image scale of 0.275 arcsec | was used. with filters matching the Johnson UBV and ]xrou-Cousius I passbauds.," A SITe $2048^2$ CCD with an image scale of 0.275 arcsec $^{-1}$ was used, with filters matching the Johnson UBV and Kron-Cousins I passbands."
+ The stars iu the fiekl were calibrated using Lauclolt(1992). standards., The stars in the field were calibrated using \citet{landolt92} standards.
+ Exposures were obtained while the sky appeared photometric. aud the reduced images were processed using the DAOphot program as implemented in," Exposures were obtained while the sky appeared photometric, and the reduced images were processed using the DAOphot program as implemented in"
+was checked with the spectropolarimetric standard 22024-].,was checked with the spectropolarimetric standard 2024-1.
+" On 2007 January 8. 300ss exposures were taken at retarder waveplate angles 0°. 22.5"" 45°. and 67.57."," On 2007 January 8, s exposures were taken at retarder waveplate angles $^\circ$ , $^\circ$, $^\circ$, and $^\circ$."
+ These data were reduced in the same manner as those obtained for our target 992207., These data were reduced in the same manner as those obtained for our target 92207.
+ We obtained the values 7 = 9.65+0.12% and ce Ξ 136.82-E0.34. which are in very good agreement with the previous measurements 7? = 9.53+0.02% and «= 136.75+0.16° by Fossati ((2007). who used numerous spectropolarmmetric observations of 22024-1 retrieved from the ESO archive.," We obtained the values $P$ = $\pm$ and $\psi$ = $\pm$ $^\circ$, which are in very good agreement with the previous measurements $P$ = $\pm$ and $\psi$ = $\pm$ $^\circ$ by Fossati (2007), who used numerous spectropolarimetric observations of 2024-1 retrieved from the ESO archive."
+ Here we describe the characteristics of the observed polarizations and their implications for the source., Here we describe the characteristics of the observed polarizations and their implications for the source.
+" Based on previous studies of H,,. it is not surprising that even in low resolution spectra. the P. Cygnt emission line of H,, of HD 92207 shows significant variability. both in line shape and equivalent width."," Based on previous studies of $_\alpha$, it is not surprising that even in low resolution spectra, the P Cygni emission line of $_\alpha$ of HD 92207 shows significant variability, both in line shape and equivalent width."
+ Figure 1. shows average line profile shapes for the seven spectra with dates labeled.," Figure \ref{fig1}
+ shows average line profile shapes for the seven spectra with dates labeled."
+ The different spectra have been shifted vertically for better display., The different spectra have been shifted vertically for better display.
+ The equivalent widths (EWs) were evaluated between +275 kms |. the terminal speed ος. of the wind(see Tab. 2)).," The equivalent widths (EWs) were evaluated between $\pm 275$ km $^{-1}$, the terminal speed $v_\infty$ of the wind(see Tab. \ref{tab2}) ),"
+ and are listed in Table 1.., and are listed in Table \ref{tab1}.
+ Note that all of the lines show a net emission., Note that all of the lines show a net emission.
+" Figure 2. shows the polarization of H,, and the neighboring continuum in the form of Q-U diagrams for each night.", Figure \ref{fig2} shows the polarization of $_\alpha$ and the neighboring continuum in the form of Q-U diagrams for each night.
+ Each panel displays the date. moving chronologically elockwise from lower left.," Each panel displays the date, moving chronologically clockwise from lower left."
+" The black point in each panel signifies the mean continuum polarization from relatively ""clean"" regions of the spectra that appear to lack lines. including oon the shortward side of H,, and 6569-6580 oon the longward side."," The black point in each panel signifies the mean continuum polarization from relatively “clean” regions of the spectra that appear to lack lines, including on the shortward side of $_\alpha$ and 6569–6580 on the longward side."
+ We determined a typical standard deviation in the continuum polarization to be approximately σ~0.01% for any given night. which is approximately the size of the plotted black squares.," We determined a typical standard deviation in the continuum polarization to be approximately $\sigma \approx 0.04\%$ for any given night, which is approximately the size of the plotted black squares."
+ Combining all seven nights. the average polarizations we (Po)=1.994 and (Pi)=(; however. no attempt has been made to correct the data for interstellar polarization. so these values cannot be considered intrinsic to the star.," Combining all seven nights, the average polarizations are $\langle P_Q\rangle = - 1.99\%$ and $\langle P_U\rangle= -2.58\%$; however, no attempt has been made to correct the data for interstellar polarization, so these values cannot be considered intrinsic to the star."
+ However. our analysis focuses on variable polarization which is intrinsic to the star.," However, our analysis focuses on variable polarization which is intrinsic to the star."
+ For wavelengths within the line. dashed and dotted line types distinguish between blue and red shifts from line center (these appear as blue and red colors in the online version of the figure).," For wavelengths within the line, dashed and dotted line types distinguish between blue and red shifts from line center (these appear as blue and red colors in the online version of the figure)."
+ Again. points considered to be within the line are those between +275 km ! of line center.," Again, points considered to be within the line are those between $\pm 275$ km $^{-1}$ of line center."
+" The central panel is an overplot of the line data for the other seven panels. indicating the full range of polarimetric variations across H,,."," The central panel is an overplot of the line data for the other seven panels, indicating the full range of polarimetric variations across $_\alpha$."
+ There is a change in the continuum polarization over the 3 month time span., There is a change in the continuum polarization over the 3 month time span.
+ The overall amplitude of change ts about0.5%.. which is over 10 times larger than the dispersion for any one of the averages.," The overall amplitude of change is about, which is over 10 times larger than the dispersion for any one of the averages."
+" Most remarkable are the observed variations in the polarization across the H,, P Cygni line.", Most remarkable are the observed variations in the polarization across the $_\alpha$ P Cygni line.
+ The average of the polarization within the line ts typically greater than the dispersion in the continuum knot by a factor of 4 or more., The average of the polarization within the line is typically greater than the dispersion in the continuum knot by a factor of 4 or more.
+ The two main exceptionsare the nights of March 11 that shows very little variation across the line and March 26 that shows the greatest variations. essentially a change in polarization of about 2%..," The two main exceptionsare the nights of March 11 that shows very little variation across the line and March 26 that shows the greatest variations, essentially a change in polarization of about ."
+Reference to the nowextensive body of published average polaruuctry provides several huudred exainples of pulsars whose coual component pairs inve prominently: depolarized outer edges.,Reference to the now extensive body of published average polarimetry provides several hundred examples of pulsars whose conal component pairs have prominently depolarized outer edges.
+ The effect is so widespread. udeed. that it is difficult o identify conipletelv convincingC» examples*t to he coutrary.," The effect is so widespread, indeed, that it is difficult to identify completely convincing examples to the contrary."
+ The stars comprising four of the five nuaiu profile classes(ey... see Paper VI) of conal single (Sq). double (D). triple (T). and five-component (M) virtually all exhibit the phenomenon as do the few stars iu the more restrictedcT aud classes.," The stars comprising four of the five main profile classes, see Paper VI) of conal single ), double ), triple ), and five-component ) virtually all exhibit the phenomenon as do the few stars in the more restricted and classes."
+" It is worth noting hat: eoodT examplesWw"" of! outerv edgeedcdeoe depolarizationYarivati fonudare amongfound starsamong starswith thwith bothinner andinner outer audedgesouter conal.: configurations:.Toysty TMall 00 ofi th ∖stars owEith inner cones discussed in the theforegoing Paper VII [Mitra Raukin (2002)] show the effect. though interestingly. in each case it is more promincut on the trailing than ou the leading The best examples of stars with little or no edge depolarization all either have (or probably have. given that some are vet not well observed) core single or inner-cone tripleT configurations: some are BO355151. 0150155. 05102:3. 055905. 0626|21. 07L028. 083515. 090619 (nain pulse}, L05552. 1322|83. and 173730. and note that. overall. these stars have mich shorter periods than is typical for the normal pulsar population."," It is worth noting that good examples of outer edge depolarization are found among stars with both inner and outer conal configurations; all three of the stars with inner cones discussed in the foregoing Paper VII [Mitra Rankin (2002)] show the effect, though interestingly, in each case it is more prominent on the trailing than on the leading The best examples of stars with little or no edge depolarization all either have (or probably have, given that some are yet not well observed) core single or inner-cone triple configurations; some are B0355+54, 0450+55, 0540+23, 0559--05, 0626+24, 0740–28, 0833–45, 0906–49 (main pulse), 1055–52, 1322+83, and 1737–30, and note that, overall, these stars have much shorter periods than is typical for the normal pulsar population."
+ We then sununumuizethe overall characteristics of the edee-depolarization phenomenon: Iu the remainder of this paper we will explore the causes audcouscquences of these circtustaces. drawing extensively on the emlier articles of this sevies. Papers IVIL (soe References).," We then summarize the overall characteristics of the edge-depolarization phenomenon: In the remainder of this paper we will explore the causes and consequences of these circumstances, drawing extensively on the earlier articles of this series, Papers I–VII (see References)."
+ We will show that these structural characteristics of conal cluission bemusaud therefore well resolved conal component parsare almost certainly the result of subbeam circulation ax iu pulsar DO913|10 (sce Deshpande Raukin 2001)., We will show that these structural characteristics of conal emission beams—and therefore well resolved conal component pairs—are almost certainly the result of subbeam circulation as in pulsar B0943+10 (see Deshpande Rankin 2001).
+ This circulation. dn sweeping a series of polarized subbeams around the maguectic axis aud past our sightline. is responsible for the outer-edge depolarization and (sometimes periodic) modal Huctuations in pulsars where the sielitline traverse cuts the emission beam centrally(e.y.. 0525|21): it is also largely respousible for the very differeut »olarization effects observed m pulsars where the sightline traverse is oblique(οι B0sQ9|71).," This circulation, in sweeping a series of polarized subbeams around the magnetic axis and past our sightline, is responsible for the outer-edge depolarization and (sometimes periodic) modal fluctuations in pulsars where the sightline traverse cuts the emission beam centrally, 0525+21); it is also largely responsible for the very different polarization effects observed in pulsars where the sightline traverse is oblique, B0809+74)."
+ Of course. only subbeanis with particular angular »olarization patterus can produce the particular sorts of depolarized profile forms that are observed. in the. remainder of. this andpaper we cucleavor 3. understand what general features are required. in them.," Of course, only subbeams with particular angular polarization patterns can produce the particular sorts of depolarized profile forms that are observed, and in the remainder of this paper we endeavor to understand what general features are required of them."
+ In the following sectious. we first briefly through our observations aud then consider the contrasting characteristics of stars. first with ative separated conal componcut pairs. aud then LAS counal single S4," In the following sections, we first briefly describe our observations and then consider the contrasting characteristics of stars, first with well separated conal component pairs, and then with conal single profiles."
+ The penultimate cone gives the results of modelling the polarized cnussion∙⋅ beam. aud we conclude with: a sunu of. our results and a discussion. of. their 2.duplications..," The penultimate section gives the results of modelling the polarized emission beam, and we conclude with a summary of our results and a discussion of their implications."
+This interpretation is al odds with the morphology of the radio emission. as discussed bv tothenflugetal.(2004) (seealsothediscussioninMorlinoetal.2009).. but at present it is not possible to rule it out.,"This interpretation is at odds with the morphology of the radio emission, as discussed by \cite{radio} \citep[see also the
+discussion in][]{morlino}, but at present it is not possible to rule it out."
+ [fit turns out to be correct. then there would be no observational constraint on the magnetic field in the shock region.," If it turns out to be correct, then there would be no observational constraint on the magnetic field in the shock region."
+ 1 on the other hand the interpretation based on strong svnchrotron losses is confirmed. the inferred levels of magnetization can be interpreted as a result. of streaming instability induced by cosmic ravs efficiently accelerated. al a parallel shock front (Bell1973:Bell2001:Bell 2004).. although alternative mechanisms have also been put forward Giacalone&Jokipii 2007).," If on the other hand the interpretation based on strong synchrotron losses is confirmed, the inferred levels of magnetization can be interpreted as a result of streaming instability induced by cosmic rays efficiently accelerated at a parallel shock front \citep{bell78,bl2001,bell2004}, although alternative mechanisms have also been put forward \citep[e.g.][]{gj07}."
+. In Table 1 we list the SNRs where evidence has been collected. from X-raw observations. lor a strong magnetic field: wy is the shock velocity. D» is the value of the magnetic field downstream of the shock as inferred from the X-ray. brightness prolile and as is the magnetic enerev density immediately. downstream of the forward shock. in units of the total kinetic pressure at upstream infinity (a2=BZ/(SmpyuZ)).," In Table \ref{tab:SNR} we list the SNRs where evidence has been collected, from X-ray observations, for a strong magnetic field: $u_0$ is the shock velocity, $B_2$ is the value of the magnetic field downstream of the shock as inferred from the X-ray brightness profile and $\alpha_2$ is the magnetic energy density immediately downstream of the forward shock, in units of the total kinetic pressure at upstream infinity $\alpha_{2}=B_2^2/(8\pi\rho_0u_0^2$ ))."
+ The data are from Volk.Berezhko&Ixsenolontov(2005). (numbers in parentheses) and [rom al. (2006)., The data are from \cite{V+05} (numbers in parentheses) and from \cite{P+06}.
+ The efficient acceleration and (he magnetic field amplification are the (wo most impressive manifestations of nonlinear diffusive acceleration al SNR. shocks., The efficient acceleration and the magnetic field amplification are the two most impressive manifestations of nonlinear diffusive acceleration at SNR shocks.
+ Both (these aspects Lave been included in a Monte Carlo scheme by Vlacdimirov.Ellison&Bykov (2006)... with the magnetic field. amplification described accordingly to the phenomenological scenario proposed bv Bell&Lucek(2001).," Both these aspects have been included in a Monte Carlo scheme by \cite{elli06}, , with the magnetic field amplification described accordingly to the phenomenological scenario proposed by \cite{bl2001}."
+. Vladimirov.Ellison&Bykov(2006) find that. when amplification is efficient. (he wave pressure makes (he plasma upstream of the shock less compressible aud the ehiange in the energy density of the magnetic turbulence across the subshock strongly affects. ow. which in turn affects the injection efficiency and the entire process of cosmic ray acceleration.," \cite{elli06} find that, when amplification is efficient, the wave pressure makes the plasma upstream of the shock less compressible and the change in the energy density of the magnetic turbulence across the subshock strongly affects $\Rs$, which in turn affects the injection efficiency and the entire process of cosmic ray acceleration."
+ In a previous paper (Caprioetal.2008a) we used a three-[uid approach (gas. cosmic ravs and Alfvénn waves). to show the verv general nature of the magnetic dvnaamical feedback: the shock dynamics is significantly affected by the turbulence backreaction as soon as (he magnetic pressure becomes comparable to Chat of the gas upstream of the subshock.," In a previous paper \citep{jumpl} we used a three-fluid approach (gas, cosmic rays and Alfvénn waves), to show the very general nature of the magnetic dynamical feedback: the shock dynamics is significantly affected by the turbulence backreaction as soon as the magnetic pressure becomes comparable to that of the gas upstream of the subshock."
+ In particular. for the magnetization levels inferred Lom available observations (see Tab. 1))," In particular, for the magnetization levels inferred from available observations (see Tab. \ref{tab:SNR}) )"
+" aud (he saturation values that are expected Irom different amplification mechanisms proposed in the literature. we found that 2,4 would naturally lead to values in the range 6-10. in good agreement with the values currently inferred. [rom observations."," and the saturation values that are expected from different amplification mechanisms proposed in the literature, we found that $\Rt$ would naturally lead to values in the range 6-10, in good agreement with the values currently inferred from observations."
+ In the present paper we use a kinetic modelfor particle acceleration at a parallel shock, In the present paper we use a kinetic modelfor particle acceleration at a parallel shock
+elliptical galaxy which satisfies the average trend. can undergo either ~ 2-4. major dry-mergers (if only one has mass ratio l:l. the rest being either 1:2 or 1:3). or ~ 2-3 major (if the mass ratio is always L1. re. at the first step we create two galaxies. each one via a merger of two units like model E (2xE). and then we let them merge togheter). before crossing the boundary set by observations.,"elliptical galaxy which satisfies the average trend, can undergo either $\sim$ 2-4 major dry-mergers (if only one has mass ratio 1:1, the rest being either 1:2 or 1:3), or $\sim$ 2-3 major dry-mergers (if the mass ratio is always 1:1, i.e. at the first step we create two galaxies, each one via a merger of two units like model E (2xE), and then we let them merge togheter), before crossing the boundary set by observations."
+ In any case. following the same argument of the previous sections. no arbitrary sequence of dry-mergers of a model E galaxy with similar progenitor can ever form a 10'M. spheroid with the highest observed |<Mg/Fe >| for that mass.," In any case, following the same argument of the previous sections, no arbitrary sequence of dry-mergers of a model E galaxy with similar progenitor can ever form a $10^{12}M_{\odot}$ spheroid with the highest observed $<Mg/Fe>$ ] for that mass."
+ Before drawing our final remarks. we summarize the results obtained in the studied case by means of fig. 5..," Before drawing our final remarks, we summarize the results obtained in the studied case by means of fig. \ref{fig_espl},"
+ where the average [«Mg/Fe >| and the average |<Fe/H >] are plotted versus Jog(c)., where the average $<Mg/Fe>$ ] and the average $<Fe/H>$ ] are plotted versus $log(\sigma)$.
+ In this figure we sketch the allowed paths as solid lines. whereas we plot as dashed lines what is impossible on the basis of stellar population arguments.," In this figure we sketch the allowed paths as solid lines, whereas we plot as dashed lines what is impossible on the basis of stellar population arguments."
+ We compare them with the observations by Nelan et al. (, We compare them with the observations by Nelan et al. (
+2006): the black dotted line being the mean value. and the shaded area brackets the observational scatter reported by the authors.,"2006): the black dotted line being the mean value, and the shaded area brackets the observational scatter reported by the authors."
+ These values are in agreement with the previous analysis by Thomas et al. (, These values are in agreement with the previous analysis by Thomas et al. (
+2005). although in Nelan et al. (,"2005), although in Nelan et al. ("
+2006) the slope of the MFMR is somewhat flatter at high velocity dispersions.,2006) the slope of the MFMR is somewhat flatter at high velocity dispersions.
+ As done in the previous sections. we follow the merger history which links progenitors to the to the as an example of a possible accretion history which may lead to the creation of a massive spheroid from. small subunits.," As done in the previous sections, we follow the merger history which links progenitors to the to the as an example of a possible accretion history which may lead to the creation of a massive spheroid from small subunits."
+ With we name the outcome of the multiple mergers described in Sec., With we name the outcome of the multiple mergers described in Sec.
+ 3.1., 3.1.
+ We already know that some of the chemical properties of such a galaxy will differ from those featured by model F. We assume that the progenitors of type A inhabit the upper (lower) left corner in the upper (lower) panel. without any further specification on their initial cc and on the dynamical outcome of the multiple mergers.," We already know that some of the chemical properties of such a galaxy will differ from those featured by model F. We assume that the progenitors of type A inhabit the upper (lower) left corner in the upper (lower) panel, without any further specification on their initial $\sigma$ and on the dynamical outcome of the multiple mergers."
+ To be conservative. in fact. we assume that they can somehow increase their c7.. but we recall that. were the build-up of model 400xA a sequence ofsystems.. its final cx would be equal to the maximum value of the stellar velocity dispersion among the progenitors (Ciotti et al..," To be conservative, in fact, we assume that they can somehow increase their $\sigma$, but we recall that, were the build-up of model 400xA a sequence of, its final $\sigma$ would be equal to the maximum value of the stellar velocity dispersion among the progenitors (Ciotti et al.,"
+ 2007. see also Nipoti et al.," 2007, see also Nipoti et al."
+ 2003 for the effects of non-homology)., 2003 for the effects of non-homology).
+ This means that the remnant merger would probably lie very close to A in the upper panel of fig. 5..," This means that the remnant merger would probably lie very close to A in the upper panel of fig. \ref{fig_espl},"
+ without moving leftward as we assume in our simple sketch., without moving leftward as we assume in our simple sketch.
+ From the upper panel of Fig. 5..," From the upper panel of Fig. \ref{fig_espl},"
+ we notice that the model 400xA is fairly close to the model that we want to reproduce through the sequence of mergers (F). therefore we mark this path (solid line) in the c plane asellowed.. meaning that. on pure chemical evolution basis. itis possible to create a massive sphoroid which satisfies the MFMR starting from @-enhanced building blocks.," we notice that the model 400xA is fairly close to the model that we want to reproduce through the sequence of mergers (F), therefore we mark this path (solid line) in the $\sigma$ plane as, meaning that, on pure chemical evolution basis, it is possible to create a massive sphoroid which satisfies the MFMR starting from $\alpha$ -enhanced building blocks."
+ On the other hand. the lower panel of Fig. 5..," On the other hand, the lower panel of Fig. \ref{fig_espl},"
+ tells us that the galaxy of type 400xA falls short in reproducing the MMR by more than two orders of magnitude., tells us that the galaxy of type 400xA falls short in reproducing the MMR by more than two orders of magnitude.
+ A path joining A with E and F (dashed line) is impossible if we want to satisfy both the observed relationships., A path joining A with E and F (dashed line) is impossible if we want to satisfy both the observed relationships.
+ We do not show here the paths leading to a final F galaxy from building blocks of type B. It is intuitive that they will be on the metallicity-c plane. but they will be outside the region delimited by the observations in the [Mg/Fe|-» plane.," We do not show here the paths leading to a final F galaxy from building blocks of type B. It is intuitive that they will be on the $\sigma$ plane, but they will be outside the region delimited by the observations in the $\sigma$ plane."
+ If we allow galaxies to form through à sequence of mergers involving different progenitors. given the random nature of the hierarchical assembly it is natural to expect a scatter in the predicted MMR and the MFMR much larger than the observed ones and no slope at all.," If we allow galaxies to form through a sequence of mergers involving different progenitors, given the random nature of the hierarchical assembly it is natural to expect a scatter in the predicted MMR and the MFMR much larger than the observed ones and no slope at all."
+ We expect this to happen also in the most recent models based on the hierarchical assembly and even if the features several gestures toward thedownsizing. as it has indeed been shown by Nagashima et al. (," We expect this to happen also in the most recent models based on the hierarchical assembly and even if the features several gestures toward the, as it has indeed been shown by Nagashima et al. ("
+2005).,2005).
+ This fact and our previous works on the quasi-monolitic formation of massive ellipticals (PMOA) as well as the effect of major wet mergers (Pipino Matteucci. 2006) allow us to derive the first important conclusions:viceversa.," This fact and our previous works on the quasi-monolitic formation of massive ellipticals (PM04) as well as the effect of major wet mergers (Pipino Matteucci, 2006) allow us to derive the first important conclusions:."
+. As expected from the discussion of section 3.4. the outcome (called in Fig. 5) ," As expected from the discussion of section 3.4, the outcome (called in Fig. \ref{fig_espl}) )"
+of a major dry merger involving two massive ellipticals. will occupy the same place of its progenitor and does not move to where it is expected (ie. at F)). therefore we mark this path as Moreover. the model has been designed to represent an average elliptical of medium size. although the observations tell us that there are galaxies of the same mass but lower a-enhacement (as low as 0.1 dex. lower boundary of the shaded region - Fig. 5..," of a major dry merger involving two massive ellipticals, will occupy the same place of its progenitor and does not move to where it is expected (i.e. at ), therefore we mark this path as Moreover, the model has been designed to represent an average elliptical of medium size, although the observations tell us that there are galaxies of the same mass but lower $\alpha$ -enhacement (as low as 0.1 dex, lower boundary of the shaded region - Fig. \ref{fig_espl},"
+ upper panel)., upper panel).
+ This means that. in principle. we should be able to observe also giant ellipticals with >| as low as ~ 0.1 dex.," This means that, in principle, we should be able to observe also giant ellipticals with $<Mg/Fe>$ ] as low as $\sim$ 0.1 dex."
+ Since this is not the case. we have not only the need of a way to increase c without having," Since this is not the case, we have not only the need of a way to increase $\sigma$ without having"
+another smaller one. and there were a close passage of a massive embrvo (seeCha&Navakshin2010).,"another smaller one, and there were a close passage of a massive embryo \citep[see][]{ChaNayakshin10}."
+. We now «pinassume that the cmbrvos are roughly in a solid body rotation with angular frequeney Quy=£i Ono. where 0<ἕμῃ lods à parameter.," We now assume that the embryos are roughly in a solid body rotation with angular frequency $\Omega_0 = \xi_{\rm rot} \Omega_{\rm break}$ , where $0 <
+\xi_{\rm rot} < 1$ is a parameter."
+" From our numerical simulations to cate. andalso Boleyetal.(2010)... £4,~ 0.1."," From our numerical simulations to date, andalso \cite{BoleyEtal10}, $\xi_{\rm rot} \sim 0.1$ ."
+ fe=d corresponds to the maximum rotation frequency., $\xi_{\rm rot} =1$ corresponds to the maximum rotation frequency.
+" We write the mass of the embryo AM=I0Mm, and use the fiducial parameters for the model embryo. [rom Navakshin(2010a).. which vielels ou,=OS AU /1-- where (4=FÉIO07i vears is. the age of⋅ the embryo."," We write the mass of the embryo $M_{\rm emb} = 10 M_J m_1$ and use the fiducial parameters for the model embryo from \cite{Nayakshin10c}, which yields $R_{\rm emb} =
+0.8$ AU $t_4^{-1/2}$, where $t_4 = t/10^4$ years is the age of the embryo."
+ sThe minimum rotation period of an embryo is The minimum angular frequency is likely to be the angular frequeney of disc rotation at the location where the embryos are born., The minimum rotation period of an embryo is The minimum angular frequency is likely to be the angular frequency of disc rotation at the location where the embryos are born.
+ This corresponds to the maximum spin period of Typical rotation period of an embryo is then Zio~ 100 vears.," This corresponds to the maximum spin period of Typical rotation period of an embryo is then $T_{\rm rot,0} \sim$ 100 years."
+" Now let a rocky planet of mass M, be formed by a direct eravitational collapse of a solicd-clominatecl region. such as the ""grain cluster” (Navakshin2010¢) (we first assume that rotational support is not important)."," Now let a rocky planet of mass $M_p$ be formed by a direct gravitational collapse of a solid-dominated region, such as the “grain cluster” \citep{Nayakshin10a} (we first assume that rotational support is not important)."
+" The dust density. pi. of such regions may be higher than that of the gas by the [actor⋅ of⋅ up to f,↝⊥⊐LO. where fj ENis the mass fraction⋅. of dust in the embryo."," The dust density, $\rho_d$, of such regions may be higher than that of the gas by the factor of up to $f_g^{-1/2} \sim 10$, where $f_g$ is the mass fraction of dust in the embryo."
+ For typical opacity values this implies the dust. densities in the range of 100 to LO' & , For typical opacity values this implies the dust densities in the range of $10 ^{-9}$ to $10^{-7}$ g $^{-3}$.
+"The specific angular momentum of the collapsing fragment is 9ἕνατρ. where Ae=(M,Agp,) is the linear dimension of the fragment."," The specific angular momentum of the collapsing fragment is $\sim
+\xi_{\rm rot} \Omega_0 R_{\rm fr}^2$, where $R_{\rm fr} = (3 M_p/ 4\pi
+\rho_d)^{1/3}$ is the linear dimension of the fragment."
+ The specifie angular momentum of a rocky planet at birth is that of the material that made it. ancl thus the rotation rate of the rocky planet is As μυαυμήκος We arrive at the following estimate of the natal rotationperiod of rocky planets: in hours. where p;=10p& g em”.," The specific angular momentum of a rocky planet at birth is that of the material that made it, and thus the rotation rate of the rocky planet is As $T_{\rm min,0} < T_{\rm rot,0} < T_{\rm max,0}$, we arrive at the following estimate of the natal rotationperiod of rocky planets: in hours, where $\rho_d = 10^{-8} \rho_{-8} $ g $^{-3}$."
+" Alternatively. wrltingZio= $54- Now the shortest possible rotation period for a planet of density. ον is actually for p,=5 gem ""."," Alternatively, writing$T_{\rm rot,0} = T_{\rm min,0} \; \xi_{\rm rot}^{-1}$ , Now the shortest possible rotation period for a planet of density $\rho_p$ is actually for $\rho_p = 5$ g $^{-3}$ ."
+ This is longer than the minimum, This is longer than the minimum
+"Πιτ, the submillimeter flux is approximately elven by which is normalized to the measured GGITz (δ0θμτιι) flux value for Arp 220 from RRiegopoulou. Lawrence. Rowan-Robiuson (1996); hero a,=3|23.","limit, the submillimeter flux is approximately given by which is normalized to the measured GHz $\mu$ m) flux value for Arp 220 from \markcite{rig96}R Rigopoulou, Lawrence, Rowan-Robinson (1996); here $\alpha_{s}=\beta+2=3$."
+ The approxination iu Eq., The approximation in Eq.
+ 3 is reasonably accurate for i<3: at higher zit extrapolates above the result of Eq. 2.., \ref{smmlong} is reasonably accurate for $z\lesssim3$; at higher $z$ it extrapolates above the result of Eq. \ref{smmexact}.
+ The radio flux is giveu by which is normalized to the measured GGIIz fiux for Απρ 220 from CCondon ct (1991)., The radio flux is given by which is normalized to the measured GHz flux for Arp 220 from \markcite{condon91}C Condon et (1991).
+" Since both S, and S, depend linearly ou the star formation rate. the scale factor f£ is the same."," Since both $S_s$ and $S_r$ depend linearly on the star formation rate, the scale factor $f$ is the same."
+ In local star forming ealaxies there is an observed spectral flattening at CGGIIz due to frec-free absorption (CC'ondou 1992)., In local star forming galaxies there is an observed spectral flattening at GHz due to free-free absorption \markcite{condon92}C Condon 1992).
+" Since CGIIz iieasureineuts at ligher redshifts sample higher frequencies that are uot sensitive to these absorption effects. we adopt the standard value a,=OLS from CCondon (1992) to extrapolate local GGIIz fluxes to GGIIz fuses. in order to be consistent with high redshift observations."," Since GHz measurements at higher redshifts sample higher frequencies that are not sensitive to these absorption effects, we adopt the standard value $\alpha_{r}=-0.8$ from \markcite{condon92}C Condon (1992) to extrapolate local GHz fluxes to GHz fluxes, in order to be consistent with high redshift observations."
+ The ratio of Eqs., The ratio of Eqs.
+" 3. and tL with 7,=1? Iv gives A relation of this type was obtained by CY99 based ou the SED of M52.", \ref{smmlong} and \ref{radio} with $T_d=47\ $ K gives A relation of this type was obtained by CY99 based on the SED of M82.
+ In their relation. the factor of 1.1 in Eq.," In their relation, the factor of 1.1 in Eq."
+" 5 As replaced by 0.25 and the power-law index «4.A,=3.8 is replaced by. L3.", \ref{ratio} is replaced by 0.25 and the power-law index $\alpha_s - \alpha_r=3.8$ is replaced by 4.3.
+ However. M82. with more tha an order of magnitude lower hnuninositv than Arp 220 (Lpjg=Lis«1010ho; Lo: UIIughes. Dunlop. Rawlings 1997). is not a ULIC. and so we would argue that our result is more appropriate for the distant submillimeter sources.," However, M82, with more than an order of magnitude lower luminosity than Arp 220 $L_{FIR}=1.78\times 10^{10}\ h_{65}\ {\rm L_\odot}$ ; \markcite{hughes97}H Hughes, Dunlop, Rawlings 1997), is not a ULIG, and so we would argue that our result is more appropriate for the distant submillimeter sources."
+ Iu addition to their analvtie result. CY99. eive. empirical results for the flux ratio versus redshift for oth Arp 220 and MS2.," In addition to their analytic result, CY99 give empirical results for the flux ratio versus redshift for both Arp 220 and M82."
+ Other authors have used the span of the four CY99 curves iu their Fie., Other authors have used the span of the four CY99 curves in their Fig.
+ 1 to estimate redshift ranges: however. the use of M82 again iutroduces au inappropriate nucertaimty in obtaining redshifts.," 1 to estimate redshift ranges; however, the use of M82 again introduces an inappropriate uncertainty in obtaining redshifts."
+ Iuvertiug Eq. 5.. ," Inverting Eq. \ref{ratio}, ,"
+the value of : can be related to the fiux ratio by which holds for :<3., the value of $z$ can be related to the flux ratio by which holds for $z\lesssim3$.
+ For higher redshitts the ratio of Eqs., For higher redshifts the ratio of Eqs.
+ 2 and { iust be used., \ref{smmexact} and \ref{radio} must be used.
+" The flax ratio method for estimatiue redshifts has the advantage that the Iuninositv distance drops out. aud thus there is no dependence on IL, or on the cosinologv."," The flux ratio method for estimating redshifts has the advantage that the luminosity distance drops out, and thus there is no dependence on ${\rm H_o}$ or on the cosmology."
+ We make the assumption that unevolved local ULICs aro representativo of the distant ULIC population., We make the assumption that unevolved local ULIGs are representative of the distant ULIG population.
+ We show in the following that the Arp 220 SED provides a good represcutation of the subinillaeter-to-radio flux ratios of the cnsemble of local ULICs placed at appropriate redshifts., We show in the following that the Arp 220 SED provides a good representation of the submillimeter-to-radio flux ratios of the ensemble of local ULIGs placed at appropriate redshifts.
+ We use a compilation of subnillineter measurements at 3000. 857. 667. and GGIIz from RRieopoulou et ((1996) to infer CGGIIz fluxes at redshifts of7.5. LLL. 0.9. and 0.06.," We use a compilation of submillimeter measurements at 3000, 857, 667, and GHz from \markcite{rig96}R Rigopoulou et (1996) to infer GHz fluxes at redshifts of7.5, 1.44, 0.9, and 0.06."
+ Fluxes at ιν were extrapolated from the, Fluxes at GHz were extrapolated from the
+In case of our observations of PSR BO3294+54 we were able to ascertain the impact of interstellar scintillations on our measurement pretty well.,In case of our observations of PSR B0329+54 we were able to ascertain the impact of interstellar scintillations on our measurement pretty well.
+ Both our values: 18.4 mJy from the simple arithmetic average and 11.7 mJy from the RISS weighted average agree well with the simple power-law spectra for this source. with spectral index of -2.2 (see Maron et al. 2000..," Both our values: 18.4 mJy from the simple arithmetic average and 11.7 mJy from the RISS weighted average agree well with the simple power-law spectra for this source, with spectral index of -2.2 (see Maron et al. \cite{maro00},"
+ and the authors on-line spectra database at html)., and the authors on-line spectra database at ).
+ Pulsar B03294-54 ts one of the strongest radio pulsars known so far. and scintillations of its radiation were analysed at various frequencies by many different authors. especially over the course of last 20 years.," Pulsar B0329+54 is one of the strongest radio pulsars known so far, and scintillations of its radiation were analysed at various frequencies by many different authors, especially over the course of last 20 years."
+ Table 3 and Figure 4. summarize the results we were able to find., Table \ref{tab3} and Figure \ref{fig4} summarize the results we were able to find.
+ Table 3. shows the list of all scintillation parameters for PSR B0329-54 at different frequencies ranging from 74 MHz to 4.8 GHz., Table \ref{tab3} shows the list of all scintillation parameters for PSR B0329+54 at different frequencies ranging from 74 MHz to 4.8 GHz.
+ At 4.8 GHz. at which we conducted the observations as well. we only found one measurement Of fpijss In Malofeev et al. (1996)).," At 4.8 GHz, at which we conducted the observations as well, we only found one measurement of $t_{\rm DISS}$ in Malofeev et al. \cite{malo}) ),"
+ which is 21.4 minutes. almost exactly two times lower than our value from the general average structure function.," which is 21.4 minutes, almost exactly two times lower than our value from the general average structure function."
+ This discrepancy may arrise because Malofeev et al. (1996)), This discrepancy may arrise because Malofeev et al. \cite{malo}) )
+ observations were very short: 130 minutes only. re. three diffractive cycles. according to our results.," observations were very short: 130 minutes only, i.e. three diffractive cycles, according to our results."
+ Additionally. one has to note that even during our observations we found that for individual sessions the measured value of fpjss can go as low as 19.0 minutes (see Table 2)).," Additionally, one has to note that even during our observations we found that for individual sessions the measured value of $t_{\rm DISS}$ can go as low as 19.0 minutes (see Table \ref{sf_table}) )."
+ Malofeev et al. (1996)), Malofeev et al. \cite{malo}) )
+ also estimated a lower limit for the diffractive modulation index Gnpiss> 0.5)., also estimated a lower limit for the diffractive modulation index $m_{\rm DISS} > 0.5$ ).
+ Figure 4 shows the frequency dependence for both diffractive and refractive. scintillation timescales (top right and top left. respectively).," Figure \ref{fig4} shows the frequency dependence for both diffractive and refractive scintillation timescales (top right and top left, respectively)."
+ Data taken from the literature are represented by full circles (for reference see Table 3)). our measurements are the open circles.," Data taken from the literature are represented by full circles (for reference see Table \ref{tab3}) ), our measurements are the open circles."
+ Evidently. our measurements for both fpiss and fgjss roughly agree with the previous. lower frequency estimates in a sense that they support the theory that the frequency dependence of the scintillation timescales is really governed by a simple power-law.," Evidently, our measurements for both $t_{\rm DISS}$ and $t_{\rm RISS}$ roughly agree with the previous, lower frequency estimates in a sense that they support the theory that the frequency dependence of the scintillation timescales is really governed by a simple power-law."
+ Adding our points to these plots extends the frequency range significantly. which in turn allows for better fit of the power-law to the data.," Adding our points to these plots extends the frequency range significantly, which in turn allows for better fit of the power-law to the data."
+ This is especially important for the RISS timescale. because it is very long at lower frequencies and the measurements often suffer because they are either poorly sampled or have only short time spans: rie. do not cover a sufficiently large number of observed refractive cycles.," This is especially important for the RISS timescale, because it is very long at lower frequencies and the measurements often suffer because they are either poorly sampled or have only short time spans; i.e. do not cover a sufficiently large number of observed refractive cycles."
+ These result in a large spread of ος measurements at lower frequencies., These result in a large spread of $t_{\rm RISS}$ measurements at lower frequencies.
+ On the other hand. our measurements were obtained from the continuous observations that usually covered at least 4-5 refractive cycles (in a normal one-day session) and up to over 20 cycles (for the longest 5-day session). which makes them very reliable.," On the other hand, our measurements were obtained from the continuous observations that usually covered at least 4-5 refractive cycles (in a normal one-day session) and up to over 20 cycles (for the longest 5-day session), which makes them very reliable."
+ Using the data summarized in Table 3 ο performed power-law fits. which yielded fpiss=9.3ΕΙ (minutes) and triss=3.9fyi? (days).," Using the data summarized in Table \ref{tab3} we performed power-law fits, which yielded $t_{\rm DISS} = 9.3 \ f_{\rm obs}^{~1.01}$ (minutes) and $t_{\rm RISS} = 3.9 \ f_{\rm obs}^{~-1.76}$ (days)."
+ Similar fits were performed for the RISS modulation index (see bottom left plot in Fig. 4)), Similar fits were performed for the RISS modulation index (see bottom left plot in Fig. \ref{fig4}) )
+ and the decorrelation bandwidth Έρως. although the latter is not a direct measurement but only an estimate (see Section 2.4)).," and the decorrelation bandwidth $B_{\mbox{\tiny DISS}}$, although the latter is not a direct measurement but only an estimate (see Section \ref{sect_sf}) )."
+ Nevertheless. both measurements roughly agree with expectations based on lower frequency data. and adding our points to the pool allowed us to obtain the frequency dependencies às ziss=0.38Fa and ," Nevertheless, both measurements roughly agree with expectations based on lower frequency data, and adding our points to the pool allowed us to obtain the frequency dependencies as $m_{\rm RISS} = 0.38 \ f_{\rm obs}^{~0.32}$ and $B_{\rm DISS} = 2.6 \ f_{\rm obs}^{~3.68}$ ."
+We have to note that for the purpose of two of these fits. namely fpjss and zipjss. we decided to omit the values obtained by Wang et al. (2008)).," We have to note that for the purpose of two of these fits, namely $t_{\rm RISS}$ and $m_{\rm RISS}$, we decided to omit the values obtained by Wang et al. \cite{wang08}) ),"
+ which are represented by triangles in their respective plots., which are represented by triangles in their respective plots.
+ These points stray by much from the rest of the data. and contradict the earlier estimates at the same frequency (1540 MHz) given by Wang et al. (2005)).," These points stray by much from the rest of the data, and contradict the earlier estimates at the same frequency (1540 MHz) given by Wang et al. \cite{wang}) )."
+ Adding our values to the data pool clearly shows that their earlier results (obtained from a different set of observations) seem to fit better into the general picture., Adding our values to the data pool clearly shows that their earlier results (obtained from a different set of observations) seem to fit better into the general picture.
+ We believe that the 2008 results are flawed: problem may lie in the method that was used for the purpose of that paper., We believe that the 2008 results are flawed; problem may lie in the method that was used for the purpose of that paper.
+ Based on other data. the actual refractive timescale at 1540 MHz should be of the order of 2 days (similar value as that given by their earlier paper. which is 2.5 days).," Based on other data, the actual refractive timescale at 1540 MHz should be of the order of 2 days (similar value as that given by their earlier paper, which is 2.5 days)."
+ Their observations were conducted almost continuously for 20 days in March. 2004. with the flux measurements based on 90-minute integrations.," Their observations were conducted almost continuously for 20 days in March 2004, with the flux measurements based on 90-minute integrations."
+ At the frequency of 1540 MHz this means that they observed only a few diffractive cycles during their single flux-integration. which may have affected. the measurements.," At the frequency of 1540 MHz this means that they observed only a few diffractive cycles during their single flux-integration, which may have affected the measurements."
+ The authors calculated /jiss based on that data., The authors calculated $m_{\rm RISS}$ based on that data.
+ However. their structure function for the flux measurements did not saturate.," However, their structure function for the flux measurements did not saturate."
+ Without the saturation level it 1s impossible to estimate the timescale from the structure function (see section 2.4)). and they decided to obtain the saturation level as Dy=ne/2 (see section 2.5)).," Without the saturation level it is impossible to estimate the timescale from the structure function (see section \ref{sect_sf}) ), and they decided to obtain the saturation level as $D_{\rm sat} = m^2/2$ (see section \ref{modul_sf}) )."
+ Using the value obtained this way. and applying it to the structure function (see their Fig.," Using the value obtained this way, and applying it to the structure function (see their Fig."
+ 5). they inferred the value of the RISS timescale of 8 hours. 6 times lower than expected and almost 8 times lower than in their previous paper.," 5), they inferred the value of the RISS timescale of 8 hours, 6 times lower than expected and almost 8 times lower than in their previous paper."
+ Probably the flaw lies in the underestimation of the RISS modulatior index (which clearly shows in our Fig., Probably the flaw lies in the underestimation of the RISS modulation index (which clearly shows in our Fig.
+ 4. as well) which resulted in a lowering of the used saturation level., \ref{fig4} as well) which resulted in a lowering of the used saturation level.
+ An inspection of the structure function they obtained (see their Fig 5. Wang et al. 2008))," An inspection of the structure function they obtained (see their Fig 5, Wang et al. \cite{wang08}) )"
+ leads to the conclusion that even a slight change of piss. and 1n tum the saturation level Di. could lead to timescale ever à half an order of magnitude higher than the 8 hours they quote. and we think that ts exactly what happened.," leads to the conclusion that even a slight change of $m_{\rm RISS}$, and in turn the saturation level $D_{\rm sat}$, could lead to timescale even a half an order of magnitude higher than the 8 hours they quote, and we think that is exactly what happened."
+ In Wang et al. (2008)), In Wang et al. \cite{wang08}) )
+ defense we may note that we had one session in our data that is coincident with their March 2004 observations. aswe performed a 38-hour observation between March 15 and 17 of 2004. and our structure function did not saturate as well. which may suggest that the fluctuations of the flux density were untypical during that period.," defense we may note that we had one session in our data that is coincident with their March 2004 observations, aswe performed a 38-hour observation between March 15 and 17 of 2004, and our structure function did not saturate as well, which may suggest that the fluctuations of the flux density were untypical during that period."
+ Knowledge of frequency dependence of both tpjss and feiss gives us an opportunity to calculate the frequency. 1.9. the observing frequency at which both timescales would be equal. which means a switch from strong to weak scintillation regimes.," Knowledge of frequency dependence of both $t_{\rm DISS}$ and $t_{\rm RISS}$ gives us an opportunity to calculate the , i.e. the observing frequency at which both timescales would be equal, which means a switch from strong to weak scintillation regimes."
+ Using the empirical formulae provided above (see, Using the empirical formulae provided above (see
+SN svstenmis.,SN systems.
+ A anultiply-imaged SN would appear as successive SNe. closely clustered on the sky (15%).," A multiply-imaged SN would appear as successive SNe, closely clustered on the sky $\la 15\arcsec$ )."
+" As gravitational lensing is achromatic. the light curves of the ""different? SNe would be identical in all bands. modulo an overall amplification shift between the different The probability of having a second SN appear iu a eivon ealaxy within one vear is roughly equivaleut to the inultiple-diuaeimes rate. and therefore well-samplec photometry of the light curves (to enable comparison of rise/fall parauueters). or searches for lenses aloug the line of sight. will be required to positively ideutifv leusiug cases."," As gravitational lensing is achromatic, the light curves of the “different” SNe would be identical in all bands, modulo an overall amplification shift between the different The probability of having a second SN appear in a given galaxy within one year is roughly equivalent to the multiple-imaging rate, and therefore well-sampled photometry of the light curves (to enable comparison of rise/fall parameters), or searches for lenses along the line of sight, will be required to positively identify lensing cases."
+ Characteristics of the SNAP mission impose a uuniber of selection constraints on observable πρίνπμασσα SNe systems., Characteristics of the SNAP mission impose a number of selection constraints on observable multiply-imaged SNe systems.
+ The distribution of time cdelavs between SN nuages can be expected to münüc those iu nultiplv-iuaged radio source aud quasar surveys, The distribution of time delays between SN images can be expected to mimic those in multiply-imaged radio source and quasar surveys.
+ From the theoretical distributions of time celavs obtained via the SUM. calculations. we find that 784 of stroug leusine events have time celavs such that they'd be seen iu one vears observation. rising to 93% with three wears of observation.," From the theoretical distributions of time delays obtained via the SUM calculations, we find that $78\%$ of strong lensing events have time delays such that they'd be seen in one year's observation, rising to $93\%$ with three years of observation."
+ The angular resolution of SNAP is to be 0/1. which allows it to directly resolve the vast majority of nage separations.," The angular resolution of SNAP is to be $0\farcs1$, which allows it to directly resolve the vast majority of image separations."
+ Au important possible systematic error for eround-based. optical surveys of multipli-inaged svstenis is that mages appearing ucar the ceuter of the lens can be obscured by the leus (Porciani&Madau2000)., An important possible systematic error for ground-based optical surveys of multiply-imaged systems is that images appearing near the center of the lens can be obscured by the lens \citep{pm00}.
+. Because of SNAP’s excellent augulu resolution. auc because baseline nuages of the galaxy without the supernovae will be available (and heuce cau be subtracted off). a SN image would be identifiable all the way into the core of a galaxy.," Because of SNAP's excellent angular resolution, and because baseline images of the galaxy without the supernovae will be available (and hence can be subtracted off), a SN image would be identifiable all the way into the core of a galaxy."
+ Even if the augular separation of SN nuages is well below the resolution limit of the SNAP telescope (which is exceedingly uulikelv for macrolensing cases). because the SNe are transitory we may nouetheless be able to see and distiuguisli images.," Even if the angular separation of SN images is well below the resolution limit of the SNAP telescope (which is exceedingly unlikely for macrolensing cases), because the SNe are transitory we may nonetheless be able to see and distinguish images."
+ As long as the light curves do not sienificautly overlap iu time. proximity (and bleudiug) of nuages is no mipedinenut to observation.," As long as the light curves do not significantly overlap in time, proximity (and blending) of images is no impediment to observation."
+" Iu addition. it is straightforward to associate even very widely spaced iniges, as we expect to see all SNe within SNAP larec (1 deg?) field of view (limited to systems with time sdelays which fall within the SNAP observational period)."," In addition, it is straightforward to associate even very widely spaced images, as we expect to see all SNe within SNAP's large $1\,\deg^2$ ) field of view (limited to systems with time delays which fall within the SNAP observational period)."
+" The likelihood results are quite sensitive to the formu, of the dark iamatter.", The likelihood results are quite sensitive to the form of the dark matter.
+ Thus far we have been considering smoothlv distributed isothermal halos filled with microscopic dark matter (e.g. axious. WIAIPS).," Thus far we have been considering smoothly distributed isothermal halos filled with microscopic dark matter (e.g. axions, WIMPS)."
+ If a significant (2 10%) fraction of halo matter is in macroscopic form (c.g. ATACTIOs. black holes). the lensing probability increases dramatically due to frequeut mucroleusing (Schucider Waeoucr 1987: Rauch 1991).," If a significant $\ga\!\!10\%$ ) fraction of halo matter is in macroscopic form (e.g. MACHOs, black holes), the lensing probability increases dramatically due to frequent microlensing (Schneider Wagoner 1987; Rauch 1991)."
+ At cosmological distances. lensing by compact objects with masses below 108AL. produces image separations of less than a few iülliareseconds aud time delays of less than a minute. indicating that multiple πιασας would no longer be directly discernible.," At cosmological distances, lensing by compact objects with masses below $10^6\,M_\odot$ produces image separations of less than a few milliarcseconds and time delays of less than a minute, indicating that multiple imaging would no longer be directly discernible."
+ The overall amplification of the combined image would still be noticeable down to lens imass seales of ~10Αι. below which the SN can no longer be treated as a point source (linear extent ~105cin corresponds to an aneular size ou the order of the Einstein angle of the lens). aud lensing ceases to have an effect.," The overall amplification of the combined image would still be noticeable down to lens mass scales of $\sim10^{-4}\,M_\odot$, below which the SN can no longer be treated as a point source (linear extent $\sim 10^{15}\,\mbox{cm}$ corresponds to an angular size on the order of the Einstein angle of the lens), and lensing ceases to have an effect."
+ In addition. if there is relative motion between the leus aud the line of sight to the source. the amplification becomes time dependoeut.," In addition, if there is relative motion between the lens and the line of sight to the source, the amplification becomes time dependent."
+ The lensing timescale at cosmological distances goes as bt~LSdavs(10?cursee!ή10PAL). with e the expansion velocity of theον SN. and a the mass of the compact leus.," The lensing timescale at cosmological distances goes as $\delta t\sim18\
+\mbox{days}\,(10^9\mbox{
+cm}\,\mbox{sec}^{-1}/v)\sqrt(M/10^{-3}\,M_\odot)$, with $v$ the expansion velocity of the SN, and $m$ the mass of the compact lens."
+ Therefore leases in the mass range 10FAL.SomXPM can produce a MACIIO-tvpe sjenal superposed ou the SN helt curve.," Therefore lenses in the mass range $10^{-4}\,M_\odot\la m \la
+10^{-3}\,M_\odot$ can produce a MACHO-type signal superposed on the SN light curve."
+ Microleusiug by colupact objects with m>10.7AL. will uauitest itself as anoverall amplification shift of the (standard candle) light curve. inducing further scatter iu the SN Ποιο diagram (Miuty. IHeaveus. Tawkius 2001).," Microlensing by compact objects with $m\ga10^{-3}\,M_\odot$ will manifest itself as an overall amplification shift of the (standard candle) light curve, inducing further scatter in the SN Hubble diagram (Minty, Heavens, Hawkins 2001)."
+ Over the course of a few vears. SNAP would provide a wutormly selected survey of many thousands of SNe. offering the possibility of a varicty of statistical tests of cosmolosv.," Over the course of a few years, SNAP would provide a uniformly selected survey of many thousands of SNe, offering the possibility of a variety of statistical tests of cosmology."
+ For example. the different curves in Fig.," For example, the different curves in Fig."
+ 1 should be distinguishable observationally—three. vearsjz data would differentiate between O94=2/3 aud O4=0 at better than 3o., \ref{opt_depth1} should be distinguishable observationally—three years' data would differentiate between $\Omega_\Lambda=2/3$ and $\Omega_\Lambda=0$ at better than $3\sigma$.
+ That the SNe are standard cauclles is completely irelevant to determining cosmology via lensing ikelihood. and thus strong lensing provides a powerful independent probe. and consistency check. for the SN ucasureineut of the distance-redshitt relation (the primary nission of SNAP).," That the SNe are standard candles is completely irrelevant to determining cosmology via lensing likelihood, and thus strong lensing provides a powerful independent probe, and consistency check, for the SN measurement of the distance-redshift relation (the primary mission of SNAP)."
+ Iu addition. the stroug lensing likelibood depends upon he chuupiness of the dark matter. and thereby allows he distinction between microscopic and. macroscopic dark matter (Linder et al.," In addition, the strong lensing likelihood depends upon the clumpiness of the dark matter, and thereby allows the distinction between microscopic and macroscopic dark matter (Linder et al."
+ 1988)., 1988).
+ Furthermore. Wvithe et al. (," Furthermore, Wyithe et al. ("
+2001) have suggested that the optical depth o iuultiple imaging could be used to distinguish )etween interacting aud non-dnteractiug dark matter.,2001) have suggested that the optical depth to multiple imaging could be used to distinguish between interacting and non-interacting dark matter.
+ A complementary cosmological measurement is offered bv he distribution of image separations aud leusiug nuage uorphologies (IKeeton.INochauck.&Sehak1997:Rusin&Ma2001:ItochanckWhite 2001).," A complementary cosmological measurement is offered by the distribution of image separations and lensing image morphologies \citep{kks97,rm01,kochwhite01}."
+. These approaches are less seusitive to the cosimological parameters. but quite sensitive to the mass distribution within halos.," These approaches are less sensitive to the cosmological parameters, but quite sensitive to the mass distribution within halos."
+ ANultiply-imaged systems allow for the determination of the IIubble coustaut. fy. through the measurement of time delays between mages (Bofsdal1961).," Multiply-imaged systems allow for the determination of the Hubble constant, $H_0$, through the measurement of time delays between images \citep{refsdal64}."
+.. This method has been developed for over a decade in multiply-inaged quasar svstenis; with varvius success (Ikoopnuiuis&Fass-nacht 1999).," This method has been developed for over a decade in multiply-imaged quasar systems, with varying success \citep{kf99}."
+. One of the limitations of such systems is that the time delays can often be quite difficult to determine (I&kundicetal.1997)., One of the limitations of such systems is that the time delays can often be quite difficult to determine \citep{kundic97}.
+. Tn the case of multiplinuaged SNe the determination of the time delay will be completely straightforward. as the time of peak luuinosity of a given SN will be determined to an accuracy of better than a dav as an inteeral part of the SNAP," In the case of multiply-imaged SNe the determination of the time delay will be completely straightforward, as the time of peak luminosity of a given SN will be determined to an accuracy of better than a day as an integral part of the SNAP"
+at a finite distance. making them subject to gauge artifacts. some propagating outward from the central DII/NS region frou t=0.,"at a finite distance, making them subject to gauge artifacts, some propagating outward from the central BH/NS region from $t=0$."
+ For f£.<200AL. an extraction sphere of 300A/ is free of propagating artifacts. whence AJ (CJ) is conserved to better than 0.3 for all cases at 1nediun resolution.," For $t<200M$, an extraction sphere of $300M$ is free of propagating artifacts, whence $M$ $J$ ) is conserved to better than $0.3$ for all cases at medium resolution."
+" Based on the above of rims. some guidauce frou perturbative results (Peters&Mathews1963:Turner1977:Berry&Cair 2010).. and a zoommwhirl ecodesic analogue of the two body problem (Pretorius&Khurana 2007).. we conjecture the following qualitative behavior as a function of r, for initial encounters resulting ina )ound svsteui."," Based on the above of runs, some guidance from perturbative results \citep{pm63,turner77,bg2010}, and a zoom-whirl geodesic analogue of the two body problem \citep{Pretorius:2007jn}, we conjecture the following qualitative behavior as a function of $r_p$ for initial encounters resulting in a bound system."
+ Consider s. the non-negative. integer nunber of oeriapsis passages before disuption/pluuge. aud ry the oeriapsis distance on the jth encouuter for dxiXon. v5I! to," Consider $n$, the non-negative, integer number of periapsis passages before disruption/plunge, and $r^i_p$, the periapsis distance on the $i^{th}$ encounter for $1\le i \le n$."
+" be aneffective periapsis distance for the final close eucouuter (νου, the corresponding r, before he dizuption/pluuge part of the final encounter)."," $r_p^{n+1}$ to be an periapsis distance for the final close encounter (i.e., the corresponding $r_p$ before the disruption/plunge part of the final encounter)."
+ Group (1) above has ο=0. exoup (2) 5)=1. aud eroup (3) DIol.," Group (1) above has $n=0$, group (2) $n=1$, and group (3) $n\geq1$."
+" The behavior of a close encounter will depeud sensitively on the distance à,=rirs between 5» and a radius ps2 of an effectiveο orbit with thesee energy aud aueular momentum.", The behavior of a close encounter will depend sensitively on the distance $\delta r^i_p=r^i_p-r_c$ between $r^i_p$ and a radius $r_c$ of an effective orbit with the energy and angular momentum.
+ If or) ls sufficicutly small (relative to AV). the orbit will exhibita whirl phase. where it asviuptotes to a nearly circular orbit.," If $\delta r^i_p$ is sufficiently small (relative to $M$ ), the orbit will exhibita whirl phase, where it asymptotes to a nearly circular orbit."
+" The smaller or, the longer the duration of the whirl with a maxinuun when ory=0 that equals the time required for the binary to lose itserccess orbital kinetic energy. either via GW omission or tidal transfer of cnerey to the NS material."," The smaller $\delta r^i_p$, the longer the duration of the whirl, with a maximum when $\delta r^i_p=0$ that equals the time required for the binary to lose its orbital kinetic energy, either via GW emission or tidal transfer of energy to the NS material."
+ The excess eucrgy is the difference between the total energv of the binary euteriug the whirl phase aud a putative binary on aquasi-circular Iuspiral at (p., The excess energy is the difference between the total energy of the binary entering the whirl phase and a putative binary on a inspiral at $r_c$.
+ Because of the requirement that ort be sinall. we ouly expect the possibility of significant whirling near the ultimate or penultimate eucounter.," Because of the requirement that $\delta r^i_p$ be small, we only expect the possibility of significant whirling near the ultimate or penultimate encounter."
+ If Is negative. the whirl will directly transition to a plunge.," If $\delta r^i_p$ is negative, the whirl will directly transition to a plunge."
+"or If positive and sinall. there will be a separation following the whirl: however the effective +. for the next encounter while rj,decreases due to CW emission. and since this is quite sizeable for our LL mass ratio system. pluuse.à,atl will likely be negative. resulting in a subsequenut"," If positive and small, there will be a separation following the whirl; however the effective $r_c$ for the next encounter while $r^i_p$ due to GW emission, and since this is quite sizeable for our 4:1 mass ratio system, $\delta r^{i+1}_p$ will likely be negative, resulting in a subsequent plunge."
+ Furthermore. the separation where the NS starts to be idally disrupted is within the radii of uustable whit orbits. hence the NS will not survive auy prolonged whir phase.," Furthermore, the separation where the NS starts to be tidally disrupted is within the radii of unstable whirl orbits, hence the NS will not survive any prolonged whirl phase."
+ Prior encounters (for larger i cases) will no exhibit significant whirling. aud while ort is laree the orbital evolution could better be described as a series of xecessiug ellipses with decreasing ecceutriieity aud semi-Major axis.," Prior encounters (for larger $n$ cases) will not exhibit significant whirling, and while $\delta r^i_p$ is large the orbital evolution could better be described as a series of precessing ellipses with decreasing eccentricity and semi-major axis."
+" Considering the number of orbits 7 as a function of heniab ry is mouotouically creasing. au he values of r, n(ry)where oue could see a notable fina whirl would be near the steps iu a(n)."," Considering the number of orbits $n$ as a function of the $r_p$, $n(r_p)$ is monotonically increasing, and the values of $r_p$ where one could see a notable final whirl would be near the steps in $n(r_p)$ ."
+" The mos »onoumaiced svhirl behavior will occur for small as Πρ) tucreases, the amount of excess kinetic cnerey (Γρ):lof over once the fnal orbit is reached decreases. and for sufficiently large η the late stages will cxscutially follow a quasi-cireular spiral."," The most pronounced whirl behavior will occur for small $n(r_p)$ ; as $n(r_p)$ increases, the amount of excess kinetic energy left over once the final orbit is reached decreases, and for sufficiently large $n$ the late stages will essentially follow a quasi-circular inspiral."
+" Figure l1 illustrates some of the varied phenomena eucouutered car the transition between n(r,)=0 aud n(ry)=1l. occmrving at à,26,58+0.08."," Figure \ref{snapshots} illustrates some of the varied phenomena encountered near the transition between $n(r_p)=0$ and $n(r_p)=1$, occurring at $r_p\approx 6.88\pm0.08$."
+" Most striking is the amount of rest mass remaining after the merecr as a fuuctiou of r, (Table 1)).", Most striking is the amount of rest mass remaining after the merger as a function of $r_p$ (Table \ref{master_table}) ).
+" For cases iu which tle NS plunges without siguificaut disruption. such as +,=5 or ry=7.5. less than of the initial mass is available to forni au accretion disk."," For cases in which the NS plunges without significant disruption, such as $r_p=5$ or $r_p=7.5$, less than of the initial mass is available to form an accretion disk."
+" IHowever. for ry, closer to the first transition (rj,=6.67 aud 6.51. sce Figure 1)). the NS is stretched into a long tidal tail. aud a sizeable amount of bound material is left to form an acerction of the initial NS rest mass for +,=6.51."," However, for $r_p$ closer to the first transition $r_p=6.67$ and 6.81, see Figure \ref{snapshots}) ), the NS is stretched into a long tidal tail, and a sizeable amount of bound material is left to form an accretion of the initial NS rest mass for $r_p=6.81$."
+" Figure 2. shows the approximate rate of fallback as a function of time for r,=Gael. 6.95. 7.22. and 7.4."," Figure \ref{fbplot} shows the approximate rate of fallback as a function of time for $r_p=6.81$, 6.95, 7.22, and 7.5."
+ This is the rate at which material on elliptical. orbits is expected to return to the accretion disk (seeBosswog2007).. (, This is the rate at which material on elliptical orbits is expected to return to the accretion disk \citep[see][]{rosswog07}. (
+These accretion rates are likely upper Inuits since they do not account for nuclear burning. see Moetzeeretal2010.,"These accretion rates are likely upper limits since they do not account for nuclear burning, see \citealt{metzger2}. .)"
+" The fallback rate is larger for cases with larger disk masses. such as p,=6.51. but all cases exhibit an approximate f°°? falloff."," The fallback rate is larger for cases with larger disk masses, such as $r_p=6.81$ , but all cases exhibit an approximate $t^{-5/3}$ falloff."
+ This time dependence was predicted. for stellar disruptions around supenuassive DII bv Rees(19858).., This time dependence was predicted for stellar disruptions around supermassive BHs by \citet{rees88}. .
+ It appears unlikely that DII-NS morecrs with fallback rates as in Fig., It appears unlikely that BH-NS mergers with fallback rates as in Fig.
+ 2. will be able to explain sGRDs with extended emission (see.e.g.Norris&Bounell2006) if this emission is due to feeding of the accretion disk at late times.," \ref{fbplot} will be able to explain sGRBs with extended emission \citep[see, e.g.,][]{nb06} if this emission is due to feeding of the accretion disk at late times."
+" For example. by f£.2:100 s. the huniuositv for tho r,=Gael case would be oulv LoMe?~2«10"" cre/s (assimnuiug :i ΟΠΟΙΟ jj= 0.1)."," For example, by $t\approx 100$ s, the luminosity for the $r_p=6.81$ case would be only $L\sim\eta\dot{M}c^2 \sim 2\times 10^{42}$ erg/s (assuming an efficiency $\eta=0.1$ )."
+" Table 1/ also showsthe total energv and angular moment lost to CVs for rj, between 5 and 12.5.", Table \ref{master_table} also showsthe total energy and angular momentum lost to GWs for $r_p$ between 5 and 12.5.
+ For the cases that we followed through merger. we find of the total mass lost to CAV οποιον. and estimate the final spins of the BUs to be (0.19x0.01.0.15.0.37.0.17.0.50.0.50). for δν=(5.00.6.67.6.81.6.95.7.22.7.50) respectively.," For the cases that we followed through merger, we find of the total mass lost to GW energy, and estimate the final spins of the BHs to be $(0.49\pm 0.01,0.45,0.37,0.47,0.50,0.50)$ for $r_p=(5.00,6.67,6.81,6.95,7.22,7.50)$ respectively."
+" The energy loss is largest for the transitional case (7,=6.95). which has a large pulse from the whirling first passage aud a second burst from the merecr (Fie. 3))."," The energy loss is largest for the transitional case $r_p=6.95$ ), which has a large pulse from the whirling first passage and a second burst from the merger (Fig. \ref{gwave_forms}) )."
+ Table 1 also shows the CAV losses for the initial encounter i cases where the NS survives the periapsis passage (columus 5 and 6)., Table \ref{master_table} also shows the GW losses for the initial encounter in cases where the NS survives the periapsis passage (columns 5 and 6).
+ These fv-by pulses can be compared with the prediction of Turner(1977) (Fig. 3..," These fly-by pulses can be compared with the prediction of \citet{turner77} (Fig. \ref{gwave_forms},"
+ lower panels. aud Fie. D).," lower panels, and Fig. \ref{gw_fit}) ),"
+ who used the Newtouian orbit together with quadrupole plivsics for the CAV enission. which we will call the (NOAÀ).," who used the Newtonian orbit together with quadrupole physics for the GW emission, which we will call the (NQA)."
+" Our waveforms show roughly the same pulse shape as he NQA prediction but have larecr amplitudes for the smaller rs, cases", Our waveforms show roughly the same pulse shape as the NQA prediction but have larger amplitudes for the smaller $r_p$ cases.
+" At r,=15 we find the gravitational waveform from the initial flv-by to be indistinguishable von the NQA prediction at our resolution (E1054).", At $r_p=15$ we find the gravitational waveform from the initial fly-by to be indistinguishable from the NQA prediction at our resolution $\pm 10$ ).
+ The cuhancement in CV energy losses for close encounters nav be due (n part) το zoom-whirl-like ychavior., The enhancement in GW energy losses for close encounters may be due (in part) to zoom-whirl-like behavior.
+" Figure [. showsthe CAV onergv loss as a ""unction of ry. along with the NQA prediction aud a fit consistent with zoom-whirl αναος (Pretorius&IKlu- iu the regime(7,= LO) where we start to see significant departures from the NQA approximation."," Figure \ref{gw_fit} showsthe GW energy loss as a function of $r_p$ , along with the NQA prediction and a fit consistent with zoom-whirl dynamics \citep{Pretorius:2007jn} in the regime$r_p \lesssim 10$ ) where we start to see significant departures from the NQA approximation."
+An interesting resultof this ecucral relativistic study that is qualitativelyconsistent with previous Newtomian studies (e.g.Leeetal. 2010).. is the ereat variability of the outcome as a function of impact parameter.,"An interesting resultof this general relativistic study that is qualitativelyconsistent with previous Newtonian studies \citep[e.g.][]{lee2010}, , is the great variability of the outcome as a function of impact parameter."
+ For example. the remuaut disk masses following merger rauge," For example, the remnant disk masses following merger range"
+"First. the stress is S—(u,u,)(b,b,).","First, the stress is $\mathcal{S} = \langle u_x u_y \rangle - \langle b_x b_y \rangle$."
+" This total stress consisting of Reynolds stress (i,t, and Maxwell stress (5,5,) gives a turbulent. viscosity 74 in Navier-Stokes equation for large-scale flows. which enhances the molecular viscosity to zi4-p.For the assumed shear flow C=—.Ar. the turbulent viscosity is given by S=v4A."," This total stress consisting of Reynolds stress $\langle u_x u_y \rangle$ and Maxwell stress $\langle b_x b_y \rangle$ gives a turbulent viscosity $\nu^T$ in Navier-Stokes equation for large-scale flows, which enhances the molecular viscosity to $\nu + \nu^T$.For the assumed shear flow $U_0 = - \A x$, the turbulent viscosity is given by $\mathcal{S} = \nu^T \A$."
+" Second. for the magnetic field considered here. the electromotive force reduces to: reltyjb, ατα LT48) ο-y= usb, abu Bo. L- usb, οί IB."," Second, for the magnetic field considered here, the electromotive force reduces to: _x = u_y b_z - u_z b_y = B_0, _y = u_z b_x - u_x b_z = B_0, _z = u_x b_y - u_y b_x = - ."
+ Note here that only three coefficients αγ. ανν and ο are non-vanishing in our configuration.," Note here that only three coefficients $\alpha_{yy}$, $\alpha_{zy}$ and $\beta$ are non-vanishing in our configuration."
+ In particular. phenomena such as the $2xJ (2). and shear current effects (?).. which have been advocated (ο generale magnetic field [or non-helical turbulence subject to rotation ancl shear as noted previously. are absent here (incontrastsimulationsof ?)..," In particular, phenomena such as the ${\bf \Omega \times J}$ \citep{Radler06} and shear current effects \citep{Rogachevskii03}, which have been advocated to generate magnetic field for non-helical turbulence subject to rotation and shear as noted previously, are absent here \citep[in contrast with the simulations of][]{Yousef08}."
+ Note that a shear-current effect could be studied by using a similar analvsis bul assuming the large-scale magnetic field to depend on z rather than vr. which will be addressed in a future contribution.," Note that a shear-current effect could be studied by using a similar analysis but assuming the large-scale magnetic field to depend on $z$ rather than $x$, which will be addressed in a future contribution."
+ To calculate the correlation Dunictions involved in the transport coellicients. we consider an incompressible forcing which is spatially homogeneous and temporally short correlated with the correlation time τε.," To calculate the correlation functions involved in the transport coefficients, we consider an incompressible forcing which is spatially homogeneous and temporally short correlated with the correlation time $\tau_f$."
+" Specifically. in Fourier space. the correlation function of the foreing is taken as: = T, snae, Ke) 0(04—lao, (K2) (4) where thetilde denotes a Fourier-transform wilh respect to (hie spatial variable."," Specifically, in Fourier space, the correlation function of the forcing is taken as: = _f )^3 ) (t_1-t_2) ) , where thetilde denotes a Fourier-transform with respect to the spatial variable."
+These data correspond to a solar portion including an isolated and round-shaped pore observed by on 1st June 2007.,These data correspond to a solar portion including an isolated and round-shaped pore observed by on 1st June 2007.
+ Images were acquired in G-band with a cadence of 30 seconds in a solar portion close to the disk center (u=0.87)., Images were acquired in G-band with a cadence of 30 seconds in a solar portion close to the disk center $\mu$ =0.87).
+ After subsonic filtering we obtained two time series; Table 1 (1 Jun) summarizes the parameters of both series in detail., After subsonic filtering we obtained two time series; Table \ref{poroseries} (1 Jun) summarizes the parameters of both series in detail.
+ The treatment performed on data does not include any restoration process because they were not degraded by atmospheric The data from on 30 September 2007 correspond to the coordinated observations described in Sect. 2.1.., The treatment performed on data does not include any restoration process because they were not degraded by atmospheric The data from on 30 September 2007 correspond to the coordinated observations described in Sect. \ref{S:obsSST}.
+ observed in G-band the emerging flux region NOAA 10971 from 00:14 to 17:59 UT with a few brief interruptions for calibrations., observed in G-band the emerging flux region NOAA 10971 from 00:14 to 17:59 UT with a few brief interruptions for calibrations.
+ Table 1 (30 Sep) summarizes the parameters of the time series in more detail., Table \ref{poroseries} (30 Sep) summarizes the parameters of the time series in more detail.
+ A significant part of the FOV in the SST images is covered by the data often presents misalignments due to tiny tracking flaws and also due to temporal interruptions., A significant part of the FOV in the SST images is covered by the data often presents misalignments due to tiny tracking flaws and also due to temporal interruptions.
+ We proceeded by aligning all of the 1030 images (18 hours) at a sub-pixel level., We proceeded by aligning all of the 1030 images (18 hours) at a sub-pixel level.
+ Subsonic filtering was applied to eliminate residual jittering in the time, Subsonic filtering was applied to eliminate residual jittering in the time
+the two particles: This makes explicit that a measurement on one particle does influence the other one. (,the two particles: This makes explicit that a measurement on one particle does influence the other one. (
+Le. the operator 1 acts trivially.),I.e. the operator $1$ acts trivially.)
+" The only effect a measurement has, is changing probabilities of other measurements into conditional probabilities, as explained just above."," The only effect a measurement has, is changing probabilities of other measurements into conditional probabilities, as explained just above."
+" More important, these conditional probabilities hold regardless of the moment at which you perform the measurement on the other particle."," More important, these conditional probabilities hold regardless of the moment at which you perform the measurement on the other particle."
+" Whether it occurs later, earlier or at the same time - that doesn’t matter at all."," Whether it occurs later, earlier or at the same time - that doesn't matter at all."
+" This forces us to abandon the (popular, but incorrect) view on the wave function collapse as an event stretching out along a space-like slice."," This forces us to abandon the (popular, but incorrect) view on the wave function collapse as an event stretching out along a space-like slice."
+" Even though this view is appealing, it creates a wrong intuition about the physics We can now solve the problem we started with in the introduction."," Even though this view is appealing, it creates a wrong intuition about the physics We can now solve the problem we started with in the introduction."
+" If the measurement is nothing but an isolated event in space time, there is no point whatsoever in trying to associate a spatial slice to it."," If the measurement is nothing but an isolated event in space time, there is no point whatsoever in trying to associate a spatial slice to it."
+ So the horizontal and tilted lines in Figure 2] actually have no meaning at all!, So the horizontal and tilted lines in Figure \ref{fig:EPR_boost} actually have no meaning at all!
+" Nothing happens these slices - the only place where something physical happens is the place of the measurement, and the implications on conditional probabilities hold for other measurements throughout the entire spacetime,past.."," Nothing happens these slices - the only place where something physical happens is the place of the measurement, and the implications on conditional probabilities hold for other measurements throughout the entire spacetime,."
+" Having answered our initial question, we now turn to the main goal of this note: understanding the delayed choice quantum eraser experiment in a more intuitive way."," Having answered our initial question, we now turn to the main goal of this note: understanding the delayed choice quantum eraser experiment in a more intuitive way."
+" We can write a schematic time evolution, as follows."," We can write a schematic time evolution, as follows."
+" The initial photon state |I> (or ‘photon wave function’ [4])) falls through the slit and then consists of an upper and a lower part: |L)))2) Each of Jsthe two components of the wave function will then be split into an entangled pair of photons: one traveling upward (7) and one downward () (IU AUN) IE AIL N2))03) The idler photon (\UNj) and |Z\,)) meets some beamsplitters, and is directed to the four detectors."," The initial photon state $|I>$ (or `photon wave function' ) falls through the slit and then consists of an upper and a lower part: + ) Each of the two components of the wave function will then be split into an entangled pair of photons: one traveling upward $\nearrow$ ) and one downward $\searrow$ ) ( |U |U + |L |L ) The idler photon $|U \searrow \rangle$ and $|L \searrow \rangle$ ) meets some beamsplitters, and is directed to the four detectors."
+" Denoting the components heading to detectors D; to D4 by their respective numbers, we get the state The prefactors can be checked by looking at Figure B]."," Denoting the components heading to detectors $D_1$ to $D_4$ by their respective numbers, we get the state The prefactors can be checked by looking at Figure \ref{fig:eraser}."
+" For simplicity, we have not taken into account polarization issues and the sign changes due to reflections."," For simplicity, we have not taken into account polarization issues and the sign changes due to reflections."
+ We now use what we have learned above., We now use what we have learned above.
+ The time ordering of the detections of the idler and signal photon does not influence at all the probability of a certain outcome., The time ordering of the detections of the idler and signal photon does not influence at all the probability of a certain outcome.
+" So the experimental outcome (encoded in the combined measurement outcomes) is bound to be the same even if we would measure the idler photon earlier, i.e. the signal photon by shortening the optical path length of the downwards configuration."," So the experimental outcome (encoded in the combined measurement outcomes) is bound to be the same even if we would measure the idler photon earlier, i.e. the signal photon by shortening the optical path length of the downwards configuration."
+" Then, if the idler is detected at D4 for example, the above state ‘collapses’ onto —4U"," Then, if the idler is detected at $D_4$ for example, the above state `collapses' onto |U"
+0.137 s per timestep.,0.137 s per timestep.
+ For 10. particles 100 timesteps ran with 57s per timestep and similar error., For $10^5$ particles 100 timesteps ran with 57s per timestep and similar error.
+ Ehe energy fraction error was 10.7 throughout the entire simulation and crifts downward svstematically with time such that the svstem becomes increasingly bound., The energy fraction error was $10^{-3}$ throughout the entire simulation and drifts downward systematically with time such that the system becomes increasingly bound.
+ The systematic loss of energy and associated: racial inward. drift we have determined. is caused by the single point lloating precision for calculations done during the Ixepler advances., The systematic loss of energy and associated radial inward drift we have determined is caused by the single point floating precision for calculations done during the Kepler advances.
+ In this paper we primarily discuss simulations with 2 planets and an exterior annulus of planetesimals., In this paper we primarily discuss simulations with 2 planets and an exterior annulus of planetesimals.
+ The initial planctesimal disk has a Uat distribution in semi-major axis so that the surface density is proportional to L/r for radius re, The initial planetesimal disk has a flat distribution in semi-major axis so that the surface density is proportional to $1/r$ for radius $r$.
+ An isotropic disk is generated with initial eccentricity and inclination dispersion such that £67)=2 and initial inclination dispersion set to 0.01., An isotropic disk is generated with initial eccentricity and inclination dispersion such that $\langle e^2\rangle = 2 \langle i^2 \rangle$ and initial inclination dispersion set to 0.01.
+ The smoothing length was typically set to. be slightly. smaller than the mean interparticle spacing., The smoothing length was typically set to be slightly smaller than the mean interparticle spacing.
+ “Pypical mean spacing values for 107 particles in units of the inner planct’s semi-major axis was around 0.01. with our smoothing length set corresponcdinglv to 0.01.," Typical mean spacing values for $10^4$ particles in units of the inner planet's semi-major axis was around 0.01, with our smoothing length set correspondingly to 0.01."
+ The smoothing length spacing was set the same for all simulations., The smoothing length spacing was set the same for all simulations.
+ The timestep was set at 0.1 where the orbital period of the inner planet is 23., The timestep was set at 0.1 where the orbital period of the inner planet is $2 \pi$.
+ A timestep of 0.1 is convenient because it ensures accuracy while still allowing of order. 10° orbital periods to be completed. within 24 hours for 105. particles., A timestep of 0.1 is convenient because it ensures accuracy while still allowing of order $10^3$ orbital periods to be completed within 24 hours for $10^4$ particles.
+ Comparisons between simulations with timesteps of 0.2. 0.1 and 0.05 vielded only insignificant dillerences in planetary migration rates. indicating that our simulations are not strongly sensitive to the chosen time step.," Comparisons between simulations with timesteps of 0.2, 0.1 and 0.05 yielded only insignificant differences in planetary migration rates, indicating that our simulations are not strongly sensitive to the chosen time step."
+ All other parameters are listed in table 1.., All other parameters are listed in table \ref{tab:tab1}.
+ Phe initial semi-major axes of the planetesimals ranges from 1.5 to 3.0., The initial semi-major axes of the planetesimals ranges from 1.5 to 3.0.
+ We work in units of the outermost planet's initial semi-major axis and with CM.=1 so that the initial period of the outermost planet is 2z., We work in units of the outermost planet's initial semi-major axis and with $GM_* = 1$ so that the initial period of the outermost planet is $2\pi$.
+ We discuss most. timescales in units of the outermost planets initial rotation period., We discuss most timescales in units of the outermost planet's initial rotation period.
+ The initial inclination ancl eccentricity of the two planets were set to zero., The initial inclination and eccentricity of the two planets were set to zero.
+ However. we have found that planet eccentricities and inclinations rapidly settled toward: zero for the planetesimal disks considered. here.," However, we have found that planet eccentricities and inclinations rapidly settled toward zero for the planetesimal disks considered here."
+ ALL simulations were run for approximately 4000 orbits., All simulations were run for approximately 4000 orbits.
+ Our code can compute all interparticle forces. however most λαοί migration simulations have neglected forces between Alanetesipials.," Our code can compute all interparticle forces, however most planet migration simulations have neglected forces between planetesimals."
+ We desire a direct. comparison between simulations that take into account all inter particles forces and those that neglect forces between planctosimals., We desire a direct comparison between simulations that take into account all inter particles forces and those that neglect forces between planetesimals.
+ To do his we run two versions of our code that are identical except hat interparticle forces are not computed i£ both particle masses are below that of a planet., To do this we run two versions of our code that are identical except that interparticle forces are not computed if both particle masses are below that of a planet.
+ In all other respects the wo sets of simulations are identical., In all other respects the two sets of simulations are identical.
+ We begin our study of migration with the simplest svsten that allows migration of the outer planet into the disk: that of two planets just inside an initially cole (low velocity dispersion) annulus of planetesimals., We begin our study of migration with the simplest system that allows migration of the outer planet into the disk; that of two planets just inside an initially cold (low velocity dispersion) annulus of planetesimals.
+ Our first set of simulation has a inner planet with mass 10* of that of the central star (similar to Jupiter) and an outer planet. with a mass ratio of 5.10 (similar to Uranus)., Our first set of simulation has a inner planet with mass $10^{-3}$ of that of the central star (similar to Jupiter) and an outer planet with a mass ratio of $5 \times 10^{-5}$ (similar to Uranus).
+ A low mass was initially chosen for the outer planet so that outer disk is not completely disrupted during migration., A low mass was initially chosen for the outer planet so that outer disk is not completely disrupted during migration.
+ A mass similar to that of Uranus was chosen simply because in our test runs it showed the most interesting variation in its migration rate., A mass similar to that of Uranus was chosen simply because in our test runs it showed the most interesting variation in its migration rate.
+ Larger planet masses showed increasing disk disruption while lower masses showed greater. degrees. of stochastic motion., Larger planet masses showed increasing disk disruption while lower masses showed greater degrees of stochastic motion.
+ A high mass inner planet was initially chosen to ensure that particles can be efficiently scattered by the outer planet and subsequently ejected. by the inner planet., A high mass inner planet was initially chosen to ensure that particles can be efficiently scattered by the outer planet and subsequently ejected by the inner planet.
+ While this mechanism can function for an inner planet that has mass ratio less than 10.7. the migration of the outer planet is then more sensitive to the distance between the planets.," While this mechanism can function for an inner planet that has mass ratio less than $10^{-3}$, the migration of the outer planet is then more sensitive to the distance between the planets."
+ Our choice of a large inner planet reduces the dependence on the distance between the planets., Our choice of a large inner planet reduces the dependence on the distance between the planets.
+ HE the distance between the two planets is large then it takes longer for the outer planet to scatter particles suflicicnth that they cross the inner planet's orbit., If the distance between the two planets is large then it takes longer for the outer planet to scatter particles sufficiently that they cross the inner planet's orbit.
+ The parameters of our simulations are listed in table 1.., The parameters of our simulations are listed in table \ref{tab:tab1}.
+ Fhree sets of simulations are run., Three sets of simulations are run.
+ The first ten. denoted 1-10. include all pair-wise forces and contain a range of planetesimal disk masses and number of planetesimals.," The first ten, denoted C1-C10, include all pair-wise forces and contain a range of planetesimal disk masses and number of planetesimals."
+ In these simulations the planetesimal mass is fixed at 10 of the stellar mass., In these simulations the planetesimal mass is fixed at $10^{-7}$ of the stellar mass.
+ The most massive disks. simulated (simulations C1. NI. and VI-V10) have a mass ratio of 10.7.," The most massive disks simulated (simulations C1, N1, and V1-V10) have a mass ratio of $10^{-3}$."
+ This is a factor of a [ew Larger than the more massive 200 Earth mass disks simulated by previous studies (Hahn&Malhotra1999:Gomesetal. 2004).," This is a factor of a few larger than the more massive 200 Earth mass disks simulated by previous studies \citep{hahn99,gomes04}."
+. Simulations. N1-N10 are identical to C1-CT10 except forces between planetesimals are neglected., Simulations N1-N10 are identical to C1-C10 except forces between planetesimals are neglected.
+ Simulations V1-V10 compute all force pairs. have the same total planetesimal disk mass. but vary the number of planetesimals and the planetesimal mass.," Simulations V1-V10 compute all force pairs, have the same total planetesimal disk mass, but vary the number of planetesimals and the planetesimal mass."
+ We first discuss the semi-major axis and cecentricity distributions curing the simulations., We first discuss the semi-major axis and eccentricity distributions during the simulations.
+ We compare simulations with and without planetesimal self-egravity., We compare simulations with and without planetesimal self-gravity.
+ We then discuss dillerences seen in the migration rates of the outer planet., We then discuss differences seen in the migration rates of the outer planet.
+ Last. we consider the effects of increasing the number of particles in à simulation while maintaining a constant mass clisk.," Last, we consider the effects of increasing the number of particles in a simulation while maintaining a constant mass disk."
+ Figure 1 shows the semi-major axis vs eccentricity distribution at dillerent. times during the Cl simulation computing all force pairs for 10000 planetesimals., Figure \ref{fig:10000} shows the semi-major axis vs eccentricity distribution at different times during the C1 simulation computing all force pairs for 10000 planetesimals.
+ Lach rame in the figure is separated by 200 orbital times where an orbital time corresponds to the initial orbital period of the outer planet., Each frame in the figure is separated by 200 orbital times where an orbital time corresponds to the initial orbital period of the outer planet.
+ As can be seen from the position of he outer planet in cach frame it migrates outward during he simulation., As can be seen from the position of the outer planet in each frame it migrates outward during the simulation.
+ The inner planet drifts inwards but not significantly so because it is more massive than the outer Xdanet., The inner planet drifts inwards but not significantly so because it is more massive than the outer planet.
+ The particles with semi-major axis interior to the outer planet are higher after the planet has passed than they were before hand. leaving a possible signature of a migrating lancet (e... Lufkinetal. 2006)).," The particles with semi-major axis interior to the outer planet are higher after the planet has passed than they were before hand, leaving a possible signature of a migrating planet (e.g., \citealt{lufkin06}) )."
+ We also notice an increase in the eccentricity. dispersion of particles in the outer disk with time during the first half of the simulation., We also notice an increase in the eccentricity dispersion of particles in the outer disk with time during the first half of the simulation.
+ We can compare this simulation to an identical one.," We can compare this simulation to an identical one,"
+"SHTs with |s|>0 take roughly twice the time of corresponding scalar ones; here the s=2 case is slightly more expensive than s=1, because computing the 5F7, from Am requires more operations than computing the ;77.","SHTs with $|s|>0$ take roughly twice the time of corresponding scalar ones; here the $s=2$ case is slightly more expensive than $s=1$, because computing the ${}_2F^\pm_{lm}$ from $\lambda_{lm}$ requires more operations than computing the ${}_1F^\pm_{lm}$."
+" The timings also indicate the scaling of the SHT cost withΝΟ., which is practically equal for HEALPix and ECP grids, but half as high for the Gaussian grid."," The timings also indicate the scaling of the SHT cost with, which is practically equal for HEALPix and ECP grids, but half as high for the Gaussian grid."
+" These timings, as well as all other timing results for iin this paper, were obtained using SSE2-accelerated code."," These timings, as well as all other timing results for in this paper, were obtained using SSE2-accelerated code."
+ Switching off this feature will result in significantly longer execution times; on the benchmark computer the factor lies in the range of 1.6 to 1.8., Switching off this feature will result in significantly longer execution times; on the benchmark computer the factor lies in the range of 1.6 to 1.8.
+" SHTs of different spins and a wide range of band limits were performed on a single CPU core, in order to assess the correlation between CPU time and/max;; the results are shown in Fig. Al."," SHTs of different spins and a wide range of band limits were performed on a single CPU core, in order to assess the correlation between CPU time and; the results are shown in Fig. \ref{l_scaling}."
+" As expected, the time consumption scales almost with /3.,.; deviations from this behaviour only become apparent for lower band limits, where SHT components of lower overall complexity (like the FFTs, computation of the Απ and By, coefficients and initialisation of data structures) are no longer negligible."," As expected, the time consumption scales almost with $\lmax^3$; deviations from this behaviour only become apparent for lower band limits, where SHT components of lower overall complexity (like the FFTs, computation of the $A_{lm}$ and $B_{lm}$ coefficients and initialisation of data structures) are no longer negligible."
+" Switching to a normalised representation of the CPU time, as was done in Fig. B],"," Switching to a normalised representation of the CPU time, as was done in Fig. \ref{norm_l_scaling},"
+" makes this effect much more evident; even for transforms of very high band limit the curves are not entirely horizontal, which would indicate a pure algorithm."," makes this effect much more evident; even for transforms of very high band limit the curves are not entirely horizontal, which would indicate a pure algorithm."
+" refhpcomp shows the ratio of CPU times for equivalent SHTs carried out using the facility of Fortran HEALPix and libpsht,, using identical hardware resources."," \\ref{hpcomp} shows the ratio of CPU times for equivalent SHTs carried out using the facility of Fortran HEALPix and , using identical hardware resources."
+" It allows several deductions: As illustrated in Fig. [i],"," It allows several deductions: As illustrated in Fig. \ref{loopstruct},"
+ all CPU-intensive parts of libpsht’ss transform algorithm are instrumented for parallel execution on multicore computers using OpenMP directives., all CPU-intensive parts of s transform algorithm are instrumented for parallel execution on multicore computers using OpenMP directives.
+ Fig., Fig.
+ shows the wall clock times measured for a single SHT that was run using a varying number ofOpenMP threads., \ref{ompscale} shows the wall clock times measured for a single SHT that was run using a varying number ofOpenMP threads.
+ The, The
+or a double exponential.,or a double exponential.
+ Therefore we are quite confident. that our results are not influenced by the particular choice of the analytical potential.," Therefore we are quite confident, that our results are not influenced by the particular choice of the analytical potential."
+ The results show no difference with respect to the method of gas- (mass-jloss we applied., The results show no difference with respect to the method of gas- (mass-)loss we applied.
+ The final configurations are similar irrespective if we remove the gas instantaneously or more smoothly during a erossing-time of the star cluster., The final configurations are similar irrespective if we remove the gas instantaneously or more smoothly during a crossing-time of the star cluster.
+ But we see immediately a strong dependency on the initial mass of the star cluster., But we see immediately a strong dependency on the initial mass of the star cluster.
+ Low mass clusters (2.2. 104M. and 2.2.10° . ) disperse their stars only close to the Galactic plane and therefore can only contribute to the build-up of the thin dise., Low mass clusters $2.2 \times 10^{4}$ $_{\odot}$ and $2.2 \times 10^{5}$ $_{\odot}$ ) disperse their stars only close to the Galactic plane and therefore can only contribute to the build-up of the thin disc.
+ This ehanges as soon as we look at the high mass clusters M. and 2.2.10° AZ. 1.," This changes as soon as we look at the high mass clusters $2.2 \times 10^{6}$ $_{\odot}$ and $2.2 \times
+10^{7}$ $M_{\odot}$ )."
+ These clusters are able to distribute their stars out to 1—£ kpe in z jeight and therefore could be the building blocks of the thick disc., These clusters are able to distribute their stars out to $1$ $4$ kpc in $z$ -height and therefore could be the building blocks of the thick disc.
+ Finally the SFE has only a second order effect on the final distribution., Finally the SFE has only a second order effect on the final distribution.
+ The lower the SFE the wider is the spread of the final particle distribution., The lower the SFE the wider is the spread of the final particle distribution.
+ The reason is quite simple as a lower SFE means that the initially binding potential of the embedded cluster is deeper such that the stars. which amount to a mass ÀA/ say. move more rapidly within it.," The reason is quite simple as a lower SFE means that the initially binding potential of the embedded cluster is deeper such that the stars, which amount to a mass $M$ say, move more rapidly within it."
+ The same stellar mass. A7. but contined by a less-massive gas potential (being equivalent to a higher star-formation efficiency). has a smaller velocity dispersion before the gas Is removed.," The same stellar mass, $M$, but confined by a less-massive gas potential (being equivalent to a higher star-formation efficiency), has a smaller velocity dispersion before the gas is removed."
+ One also has to note that even though a SFE of 0.4 should in principal lead to a surviving remnant star cluster. for all masses except our highest mass models. these remnants were not massive enough to survive the whole LO Gyr of simulation time.," One also has to note that even though a SFE of $0.4$ should in principal lead to a surviving remnant star cluster, for all masses except our highest mass models, these remnants were not massive enough to survive the whole $10$ Gyr of simulation time."
+ With a SFE of 0.4 the cluster still “pops” but instead of dissolving completely a small fraction of stars re-virialises into a bound star cluster., With a SFE of 0.4 the cluster still 'pops' but instead of dissolving completely a small fraction of stars re-virialises into a bound star cluster.
+" Kroupaetal.(2001). performed high-precision Nbody modelling of this process explaining that a pre-Orion Nebula Cluster evolves into a Pleiades type cluster whereby 2/3rds of the population ""pop away""."," \citet{Kro2001} performed high-precision Nbody modelling of this process explaining that a pre-Orion Nebula Cluster evolves into a Pleiades type cluster whereby $2/3$ rds of the population ""pop away""."
+ The majority of stars thus gets distributed as unbound stars inside the MW potential., The majority of stars thus gets distributed as unbound stars inside the MW potential.
+ The fraction of stars which remains bound is a function of the SPE., The fraction of stars which remains bound is a function of the SFE.
+ Por the low-mass cases this bound remnant does not survive the tidal forces of the MW for very long., For the low-mass cases this bound remnant does not survive the tidal forces of the MW for very long.
+ Only the more massive remnants stemming from massive initial clusters are able to survive the full 10 Gyr., Only the more massive remnants stemming from massive initial clusters are able to survive the full $10$ Gyr.
+ The future fate of these remnants 1s of no concerns in our study. as our code is not able to describe the internal dynamics of star clusters over long time intervals correctly.," The future fate of these remnants is of no concerns in our study, as our code is not able to describe the internal dynamics of star clusters over long time intervals correctly."
+ As stuted above we then fitted exponential functions to the obtained star distributions in z-direction., As stated above we then fitted exponential functions to the obtained star distributions in $z$ -direction.
+ For all models we tried to fit a single exponential of the form: where {τμ represents the z-scale height of the thin dise and compared the results with the ones obtained if we used a double exponential function of the form: where fyick is the z-seale height of the thick disc.," For all models we tried to fit a single exponential of the form: where $h_{\rm z,thin}$ represents the z-scale height of the thin disc and compared the results with the ones obtained if we used a double exponential function of the form: where $h_{\rm z,thick}$ is the z-scale height of the thick disc."
+" The fitted values of P,ni and fyµας are shown in Tab. 2.."," The fitted values of $h_{\rm z,thin}$ and $h_{\rm z,thick}$ are shown in Tab. \ref{tab:fit}."
+ For the low mass clusters the best fit was always the single exponential tit. while for the high mass clusters the distribution was better fitted with a double exponential.," For the low mass clusters the best fit was always the single exponential fit, while for the high mass clusters the distribution was better fitted with a double exponential."
+ In the table we only give the values for the best fit of the two fitting functions adopted., In the table we only give the values for the best fit of the two fitting functions adopted.
+ Again it is clear that the method of removing the gas has no influence on the tinal distribution., Again it is clear that the method of removing the gas has no influence on the final distribution.
+ But we see a clear trend towards larger scale-lengths by increasing the cluster mass and by lowering the SFE., But we see a clear trend towards larger scale-lengths by increasing the cluster mass and by lowering the SFE.
+" For an example. in the case M4=2.2«10"" M. and a SFEs of 40% or 20%. the z scale-height of the thin dise is about 14 pe."," For an example, in the case $M_{\rm ecl} = 2.2 \times 10^{6}$ $_{\odot}$ and a SFEs of $40\%$ or $20\%$, the z scale-height of the thin disc is about $14$ pc."
+ Using a SFE of only 1 per cent gives a z-scale height of 16—15.2 pe., Using a SFE of only $1$ per cent gives a z-scale height of $16$ $18.2$ pc.
+ The same type of behavior can be observed for the thick disces., The same type of behavior can be observed for the thick discs.
+(2002) and (9~0.5) by Seharl (2001).,(2002) and $\beta \sim 0.5$ ) by Scharf (2004).
+ However. this model is unable to reproduce the shape of the surface brightuess distribution iu the edge region (1 kpc <rede).," However, this model is unable to reproduce the shape of the surface brightness distribution in the edge region $4$ kpc $< r < r_{\rm edge}$ )."
+ Both cleusity models I?LICMI (solid liue) aud PL-ICMN2 (daslied line) reproduce the surface brightness profile across the edge well., Both density models PL+ICM1 (solid line) and PL+ICM2 (dashed line) reproduce the surface brightness profile across the edge well.
+ The fit parameters for these models are summarized in Table 2.., The fit parameters for these models are summarized in Table \ref{tab:fits}.
+ The electron deusity we derive by fitting the surface brightuess profile directly for gas inside the edge (3.910an [for the PL+ICK] model and 1.3x10.em. ?for the PL-ICM? model) are in excellent agreement with the beta model result (~Lx107 5)) result by Scharf. (2001). althoughour value for the cluster gas is 20% higher.," The electron density we derive by fitting the surface brightness profile directly for gas inside the edge $3.9 \times 10^{-3}$ for the PL+ICM1 model and $4.3 \times 10^{-3}$ for the PL+ICM2 model) are in excellent agreement with the beta model result $\sim 4 \times 10^{-3}$ ) result by Scharf (2004), althoughour value for the cluster gas is $\sim 20\%$ higher."
+ Following Vikhlinin (2001) we use the densities aud the temperatures [rom our surface brightness aud spectral fits to caleulate the gas pressures in tlie [ree stream region and in the galaxy eas just inside the edge. which 1s iu pressure equilibrium with gas at the stagnation point (where the relative cluster gas velocity vanishes).," Following Vikhlinin (2001) we use the densities and the temperatures from our surface brightness and spectral fits to calculate the gas pressures in the free stream region and in the galaxy gas just inside the edge, which is in pressure equilibrium with gas at the stagnation point (where the relative cluster gas velocity vanishes)."
+" The difference between tle two pressures is interpreted as the rai presswe of the intracluster medium on the leadiug edge of the galaxy (tje cold front): while the ratio of the pressures allows us to determitue the lach number Af,=v/e, (where ej is the speec of souud in the cluster [ree stream region or the cold gas cloud moving tlrough the hot ICM.", The difference between the two pressures is interpreted as the ram pressure of the intracluster medium on the leading edge of the galaxy (the cold front); while the ratio of the pressures allows us to determine the Mach number $M_1=v/c_1$ (where $c_1$ is the speed of sound in the cluster free stream region) for the cold gas cloud moving through the hot ICM.
+ Forcompleteness. the pressure-jump versus ach number relation. plotted i Figure 7.. is given by :," Forcompleteness, the pressure-jump versus Mach number relation, plotted in Figure \ref{fig:mach}, , is given by :"
+understood. and so these initial conditions. although reasonable. are not rigorously justified.,"understood, and so these initial conditions, although reasonable, are not rigorously justified."
+ In particular. we note that planetesimals can acquire moderate (~0.05) eccentricities if they orbit within a protoplanetary dise that has density fluctuations induced by its gravity (?) or by the MRI (?)..," In particular, we note that planetesimals can acquire moderate $\sim 0.05$ ) eccentricities if they orbit within a protoplanetary disc that has density fluctuations induced by its self-gravity \citep*{2008MNRAS.385.1067B} or by the MRI \citep{2005A&A...443.1067N}."
+" Starting at different eccentricities can be dealt with by changing the proper eccentricity c, in Equation C149). so long as ορ is a single-valued function of semi- axis."," Starting at different eccentricities can be dealt with by changing the proper eccentricity $e_{\mathrm{p}}$ in Equation \ref{eq:tcross}) ), so long as $e_{\mathrm{p}}$ is a single-valued function of semi-major axis."
+ We have also introduced the planet instantaneously. assuming that it forms at time /=0.," We have also introduced the planet instantaneously, assuming that it forms at time $t=0$."
+ Statements about the stirring time-scale should therefore be qualitied by including the time taken for the planet to form. during which time the planetesimals will themselves be growing.," Statements about the stirring time-scale should therefore be qualified by including the time taken for the planet to form, during which time the planetesimals will themselves be growing."
+ This will affect the maximum radius of bodies which exist in the dise when the orbits begin to cross., This will affect the maximum radius of bodies which exist in the disc when the orbits begin to cross.
+ Dynamically. we have neglected any non-secular dynamical effects of the planet on the disc.," Dynamically, we have neglected any non-secular dynamical effects of the planet on the disc."
+ While this is valid for the planetesimals far from the planet. when they are close they begin to experience resonant interactions as well as secular (see Figure 29).," While this is valid for the planetesimals far from the planet, when they are close they begin to experience resonant interactions as well as secular (see Figure \ref{fig:tsec-a}) )."
+ Furthermore. the asymptotic expression for crossing time-scale (Equation 153) is no longer valid.," Furthermore, the asymptotic expression for crossing time-scale (Equation \ref{eq:tcrossint}) ) is no longer valid."
+" In fact. this approximation over- the time-scale because. as edy. We find oloox while e,ep C2). and ο {ωνQO. whereas extrapolation of the asymptotic expression predicts a finite. value."," In fact, this approximation over-estimates the time-scale because, as $a \to a_{\mathrm{pl}}$, we find $A\to\infty$ while $e_{\mathrm{f}}\to e_{\mathrm{pl}}$ \citep{1999ssd..book.....M}, and so $t_{\mathrm{cross}}\to 0$, whereas extrapolation of the asymptotic expression predicts a finite value."
+ So both simplifying the secular interactions and neglecting non-secular interactions will tend to overestimate the crossing time-scale., So both simplifying the secular interactions and neglecting non-secular interactions will tend to overestimate the crossing time-scale.
+ Thus the simplified dynamics underestimates the ability of planets to stir discs located close to them., Thus the simplified dynamics underestimates the ability of planets to stir discs located close to them.
+ This may be particularly relevant for such systems as Fomalhaut. where several strong resonances lie in the disc (2)..," This may be particularly relevant for such systems as Fomalhaut, where several strong resonances lie in the disc \citep{2009ApJ...693..734C}."
+ We have neglected. any damping in the disc., We have neglected any damping in the disc.
+ Collisional damping has been invoked to explain the low proper eccentricities in Fomalhaut's dise (22)..," Collisional damping has been invoked to explain the low proper eccentricities in Fomalhaut's disc \citep{2006MNRAS.373.1245Q,2009ApJ...693..734C}."
+ Reducing proper eccentricity increases both the time-scale for orbit crossing and the relative velocities. as we have already described.," Reducing proper eccentricity increases both the time-scale for orbit crossing and the relative velocities, as we have already described."
+ Finally. while we have focussed on planetary secular perturbations as a dise stirring mechanism. and compared them to planet formation. it is important to realise that there may be other causes of dynamical excitation.," Finally, while we have focussed on planetary secular perturbations as a disc stirring mechanism, and compared them to planet formation, it is important to realise that there may be other causes of dynamical excitation."
+ ? investigated the effects of a stellar flyby. but found that the perturbations were rapidly damped by collisions.," \cite{2002AJ....123.1757K} investigated the effects of a stellar flyby, but found that the perturbations were rapidly damped by collisions."
+" Other mechanisms which have not been so thoroughly investigated in the context of debris disc evolution include planet formation proceeding more rapidly when a opening giant planet has formed €??).. or an outer disc being stirred by planetesimals that have been scattered out from the inner system (2),"," Other mechanisms which have not been so thoroughly investigated in the context of debris disc evolution include planet formation proceeding more rapidly when a gap-opening giant planet has formed \citep{2000ApJ...540.1091B,2005ApJ...626.1033T}, or an outer disc being stirred by planetesimals that have been scattered out from the inner system \citep{2004ApJ...614..497G}."
+ Our main conclusion is that a planetesimal belt at several 10's of AU can be stirred by an eccentric giant planet at only a few AU., Our main conclusion is that a planetesimal belt at several 10's of AU can be stirred by an eccentric giant planet at only a few AU.
+ Debris dises do not require any bodies larger than a few km in size beyond a few AU from the star to explain the observed dust production., Debris discs do not require any bodies larger than a few km in size beyond a few AU from the star to explain the observed dust production.
+" To reach this conclusion, we investigated the effects of secular perturbations from an eccentric planet on a dynamically cold disc. to assess whether this might be the origin of debris dise stirring."," To reach this conclusion, we investigated the effects of secular perturbations from an eccentric planet on a dynamically cold disc, to assess whether this might be the origin of debris disc stirring."
+ Over time. neighbouring orbits acquire sufficiently different eccentricities and longitudes of periapse that they begin ο cross.," Over time, neighbouring orbits acquire sufficiently different eccentricities and longitudes of periapse that they begin to cross."
+ We derived an analytical expression for the time for... or this to occur (Equation 162). which agrees well with previously sublished N-body simulations.," We derived an analytical expression for the time $t_{\mathrm{cross}}$ for this to occur (Equation \ref{eq:tcrossext}) ), which agrees well with previously published N-body simulations."
+ After this time. the planet's secular xrturbations quickly impose a mean relative velocity leeges.," After this time, the planet's secular perturbations quickly impose a mean relative velocity $\left<v_{rel}\right> \approx 1.4e_{\mathrm{f}}v_{\mathrm{kep}}$ ."
+ This is similar to that for a planetesimal swarm without external perturbers. despite the high degree of apsidal alignment orced on the orbits by secular perturbations.," This is similar to that for a planetesimal swarm without external perturbers, despite the high degree of apsidal alignment forced on the orbits by secular perturbations."
+ When the relative velocities increase. the dise may brighten as a result of increased dust production. and further growth of jxanetesimals may be inhibited or halted.," When the relative velocities increase, the disc may brighten as a result of increased dust production, and further growth of planetesimals may be inhibited or halted."
+ We derived an expression or «. the maximum range of semi-major axes over which a jyanet's perturbations can destroy planetesimals (Equation 269).," We derived an expression for $a^*$, the maximum range of semi-major axes over which a planet's perturbations can destroy planetesimals (Equation \ref{eq:a*weak}) )."
+ The range « increases with planetary eccentricity and strongly increases with planetary semi-major axis., The range $a^*$ increases with planetary eccentricity and strongly increases with planetary semi-major axis.
+ It is also a function of jxanetesimal size., It is also a function of planetesimal size.
+ For the weakest planetesimals (~100 mm in radius). we found typical values of a? of several hundred AU.," For the weakest planetesimals $\sim 100$ m in radius), we found typical values of $a^*$ of several hundred AU."
+ We then compared the time-scales for planet stirring with the ime-scales for self stirring from the models of ?.., We then compared the time-scales for planet stirring with the time-scales for self stirring from the models of \cite{2008ApJS..179..451K}.
+ Because for... ws a stronger dependence on e than does /;4. assuming that only hese two proceeses are operating. we found that typically the dise closer to the planet will be planet-stirred. and the dise further away will be self-stirred. and we identified another parameter Φ which demareates the outer reach of the planet's perturbations. beyond which the self-stirring time-scale is shorter.," Because $t_{\mathrm{cross}}$ has a stronger dependence on $a$ than does $t_{\mathrm{Pl}}$, assuming that only these two proceeses are operating, we found that typically the disc closer to the planet will be planet-stirred, and the disc further away will be self-stirred, and we identified another parameter $\Phi$ which demarcates the outer reach of the planet's perturbations, beyond which the self-stirring time-scale is shorter."
+ This parameter «P is typically much smaller than e. and whether a dise can be planet- before it is self-stirred depends on the dise density as well as planet parameters.," This parameter $\Phi$ is typically much smaller than $a^*$, and whether a disc can be planet-stirred before it is self-stirred depends on the disc density as well as planet parameters."
+ For a sample of RV planet hosts observed by we find no correlation between the magnitude of or a and the presence of a disc., For a sample of RV planet hosts observed by we find no correlation between the magnitude of $\Phi$ or $a^*$ and the presence of a disc.
+ While this may suggest that planet stirring is not ubiquitous. the degree of correlation we would expect is not clear.," While this may suggest that planet stirring is not ubiquitous, the degree of correlation we would expect is not clear."
+ However. for some individual systems it appears that a known planet will have stirred the dise on a time-scale shorter than the system age and/or before it it likely to self-stir.," However, for some individual systems it appears that a known planet will have stirred the disc on a time-scale shorter than the system age and/or before it it likely to self-stir."
+ In particular. we identify ο Eridani. Fomalhaut. and HD 38529 as being in this category.," In particular, we identify $\epsilon$ Eridani, Fomalhaut, and HD 38529 as being in this category."
+ Future studies of such dises should take into account the effects of planetary perturbations on the dise’s evolution. collisional as well as dynamical.," Future studies of such discs should take into account the effects of planetary perturbations on the disc's evolution, collisional as well as dynamical."
+ We also speculate that the dises of HD 181327 and HD 202917 may have been stirred by as yet undiscovered planets. since the only other stirring mechanism proposed would only work for discs with high surface densities LOMMMSN).," We also speculate that the discs of HD 181327 and HD 202917 may have been stirred by as yet undiscovered planets, since the only other stirring mechanism proposed would only work for discs with high surface densities $\ge 10$ MMSN)."
+ Future work in this area should further investigate the effects of planetary perturbations on the collisional evolution of a disc. with a view to clarifying the extent to which the perturbations can inhibit the further growth of planetesimals and determining the evolution of dust production and hence IR. luminosity.," Future work in this area should further investigate the effects of planetary perturbations on the collisional evolution of a disc, with a view to clarifying the extent to which the perturbations can inhibit the further growth of planetesimals and determining the evolution of dust production and hence IR luminosity."
+ The latter in particular will enable valuable observational tests of planet-stirring and its role other stirring mechanisms., The latter in particular will enable valuable observational tests of planet-stirring and its role other stirring mechanisms.
+ We also note that. while we have focused on internal perturbers in our discussion. our formulae for external perturbers will be relevant for investigations into the effects of binary companions on dises.," We also note that, while we have focused on internal perturbers in our discussion, our formulae for external perturbers will be relevant for investigations into the effects of binary companions on discs."
+ AJM is grateful for the support of an STFC studentship., AJM is grateful for the support of an STFC studentship.
+ The authors wish to thank Z. Leinhardt for useful discussions on collisions. and the reviewer John Chambers for suggesting a number of improvements to the paper.," The authors wish to thank Z. Leinhardt for useful discussions on collisions, and the reviewer John Chambers for suggesting a number of improvements to the paper."
+of confirmed that the CE channel becomes very inefficient after the first few Gyrs. and the WD-WD merging should be the predominant mechanism for the formation of EHB stars in old stellar systems.,"of confirmed that the CE channel becomes very inefficient after the first few Gyrs, and the WD-WD merging should be the predominant mechanism for the formation of EHB stars in old stellar systems."
+ As a consequence. cluster EHBs should be principally single-star products of merging. and close systems should be rare.," As a consequence, cluster EHBs should be principally single-star products of merging, and close systems should be rare."
+ However. measured a higher f. in and and25%.. respectively). and argued that the models proposed by cannot simultaneously account for the low f. measured in and these much higher values found in clusters only 1-2 Gyr younger.," However, measured a higher $f_\mathrm{c}$ in and and, respectively), and argued that the models proposed by cannot simultaneously account for the low $f_\mathrm{c}$ measured in and these much higher values found in clusters only 1-2 Gyr younger."
+ Unfortunately. their results suffered too large uncertainties to be conclusive.," Unfortunately, their results suffered too large uncertainties to be conclusive."
+ In summary. the binary scenario has not been disproved. but its ability to reproduce all the observations in GCs still has to be demonstrated.," In summary, the binary scenario has not been disproved, but its ability to reproduce all the observations in GCs still has to be demonstrated."
+ An alternative model of EHB star formation has received great attention in recent years: one incorporating the primordial helium enhancement., An alternative model of EHB star formation has received great attention in recent years: one incorporating the primordial helium enhancement.
+ A super-solar surface helium abundance was recognized early on as a possible cause of the heavy-mass loss underlying the formation of EHB stars. but non-canonical mixing phenomena had to be invoked," A super-solar surface helium abundance was recognized early on as a possible cause of the heavy-mass loss underlying the formation of EHB stars, but non-canonical mixing phenomena had to be invoked."
+ The discovery of multiple stellar populations in GCs opened a new frontier. because the observational results of apparently imply that the bluer main-sequence (MS) in is most probably populated by He-enriched stars.," The discovery of multiple stellar populations in GCs opened a new frontier, because the observational results of apparently imply that the bluer main-sequence (MS) in is most probably populated by He-enriched stars."
+ The currently preferred scenario 1s that these objects constitute a second stellar generation that formed from material polluted by either intermediate-mass AGB stars(??).. or rapidly rotating massive MS stars(??)..," The currently preferred scenario is that these objects constitute a second stellar generation that formed from material polluted by either intermediate-mass AGB stars, or rapidly rotating massive MS stars."
+ The models of multiple populations with different helium abundances successfully reproduce both the MS splitting and the multimodal HB morphology of both and(?)., The models of multiple populations with different helium abundances successfully reproduce both the MS splitting and the multimodal HB morphology of both and.
+. The He-enhancement thus represents a promising model. and alternative to the binary scenario. for the formation of EHB stars in GCs.," The He-enhancement thus represents a promising model, and alternative to the binary scenario, for the formation of EHB stars in GCs."
+ In this paper. we present our results of a search for EHB binaries in2808.," In this paper, we present our results of a search for EHB binaries in."
+. Preliminary results. pointing to a close binary fraction higher than in6752.. were presented by and?.," Preliminary results, pointing to a close binary fraction higher than in, were presented by and."
+. In this context. this cluster 1s a key object because 1t is noticeably younger than (?).. allowing the study of the f.-age relation foreseen by the binary scenario.," In this context, this cluster is a key object because it is noticeably younger than , allowing the study of the $f_\mathrm{c}$ -age relation foreseen by the binary scenario."
+ Moreover. this cluster represents one of the greatest successes of the He-enhancement scenario. because and were able to model both its multimodal HB and the multiple MS with three stellar populations of increasing primordial heltum abundance.," Moreover, this cluster represents one of the greatest successes of the He-enhancement scenario, because and were able to model both its multimodal HB and the multiple MS with three stellar populations of increasing primordial helium abundance."
+ We selected 83 hot HB stars in from the photometric catalog of?.. of magnitude between V=17.4. corresponding approximatively to the Grundahl jump(2).. and the limiting magnitude V219.5 imposed by program feasibility.," We selected 83 hot HB stars in from the photometric catalog of, of magnitude between $V$ =17.4, corresponding approximatively to the Grundahl jump, and the limiting magnitude $V$ =19.5 imposed by program feasibility."
+ In Table 4.. we give the IDs. the coordinates. and the photometric data of the targets from?.. and their location in the cluster CMDs shown in Figure I..," In Table \ref{t_datatarg}, we give the IDs, the coordinates, and the photometric data of the targets from, and their location in the cluster CMDs shown in Figure \ref{f_cmd}."
+ Twenty-four single spectra of target stars were collected in fourteen epochs between January |] and March 24. 2007. with the GIRAFFE spectrograph at the VLT-UT2 telescope. in both visitor and service mode.," Twenty-four single spectra of target stars were collected in fourteen epochs between January 11 and March 24, 2007, with the GIRAFFE spectrograph at the VLT-UT2 telescope, in both visitor and service mode."
+ The instrument setup H7A provided high-resolution 0000) spectra centered on the Hy line., The instrument setup H7A provided high-resolution 000) spectra centered on the $_\beta$ line.
+ The temporal sampling was carefullyplanned to maximize the detection probability of binaries with any period, The temporal sampling was carefullyplanned to maximize the detection probability of binaries with any period
+Clusters of ealaxies contaiu thermal electrous with a temperature of several keV in the IC.,Clusters of galaxies contain thermal electrons with a temperature of several keV in the ICM.
+ Several clusters also contain clectrous with cuergics of several GeV. as demonstrated bv the presence of cluster radio halos (Feretti&Ciovanuimd.1996:Feretti.1999.forobserva-tional reviews)...," Several clusters also contain electrons with energies of several GeV, as demonstrated by the presence of cluster radio halos \cite[for observational
+reviews]{1996IAUS..175..333F,1999dtrp.conf....3F}."
+ The origin of these high energy. electrons is still a iuvsterv., The origin of these high energy electrons is still a mystery.
+ Several theories have been proposed (Euflin.1999b.forareview) such as electrou iujectiou by radio galaxies (Jaffe.1977:Rephach.1979:Valtaoja. 1981).. intu electron (re-Jacceleration by plasima- or shock-waves (Jaffe.1977:Schlickeiseretal. 2000). and secondary article injection after hadromic interactions of cosmic ταν protons with the backeround gas (Dennison.1980:Vestraud.1982:Blasi&Colafrancesco.1999:Dolag&Endlin. 2000)..," Several theories have been proposed \cite[for a review]{Pune99} such as electron injection by radio galaxies \cite{1977ApJ...212....1J,1977ApJ...212..608R,1979ApJ...227..364R,1984A&A...135..141V}, in-situ electron (re-)acceleration by plasma- or shock-waves \cite{1977ApJ...212....1J,1987A&A...182...21S,1993MNRAS.263...31T,1993ApJ...406..399G,1993ApJ...407L..53R,1996SSRv...75..279V,1997A&A...321...55D,1999ApJ...518..594R,1999ApJ...520..529S,2000ApJ...544..686L,2001MNRAS.320..365B,2000A&A...362..886B,2000ApJ...535..586T}, and secondary particle injection after hadronic interactions of cosmic ray protons with the background gas \cite{1980ApJ...239L..93D,1982AJ.....87.1266V,1999APh....12..169B,2000A&A...362..151D}."
+ The different SCCLALIOS lead to VOYV iffereut expectations for the electron προςσπα in the unobserved intermediate CLCOYSOV rouigo., The different scenarios lead to very different expectations for the electron spectrum in the unobserved intermediate energy range.
+ Tn the situ acceleration scenario. the thermal aud ultra-relativistic electrou »pulatious have to be connected by a siupra-thermal »pulatiou.," In the in-situ acceleration scenario, the thermal and ultra-relativistic electron populations have to be connected by a supra-thermal population."
+ The exact shape of this population dependy. ou the details of he acceleration mechanisu aud the high Coulomb euergy losses in this energy regime (Dogicl.2000:Blasi.2000:Petrosian. 2001)..," The exact shape of this population depends on the details of the acceleration mechanism and the high Coulomb energy losses in this energy regime \cite{2000A&A...357...66D,2000ApJ...532L...9B,2001ApJ...557..560P}."
+ In Fig., In Fig.
+ 1 oue possible realization of such a supra-tlermal population is sketched., \ref{fig:clsketch} one possible realization of such a supra-thermal population is sketched.
+ But if the electrons are injected with already high cucreics iuto the ICAL e.g. by escape from radio galaxies. or by radronic secondary particle production. then the lower eud of the intermediate enerev range should be eiptied bv the rapid Coulomb euergv losses. also sketched iu Fie. L..," But if the electrons are injected with already high energies into the ICM, e.g. by escape from radio galaxies, or by hadronic secondary particle production, then the lower end of the intermediate energy range should be emptied by the rapid Coulomb energy losses, also sketched in Fig. \ref{fig:clsketch}."
+ Further. additional spectral features can exist iu this euergv range. since the loug electron cooling time of ICAL electrons at 100. MeV. allows the conservation of the imprints of earlier electron injection events for cosmological tines (Sarazin&Lieu.1998:Sarazin.1999)..," Further, additional spectral features can exist in this energy range, since the long electron cooling time of ICM electrons at 100 MeV allows the conservation of the imprints of earlier electron injection events for cosmological times \cite{1998ApJ...494L.177S,1999ApJ...520..529S}."
+ The activity phases of individual radio galaxies. as they appear to our instruments. are short-lived phenomena compared to cosmological time-scales.," The activity phases of individual radio galaxies, as they appear to our instruments, are short-lived phenomena compared to cosmological time-scales."
+ After a few teu wullion vears the ceutral engine in the ealactic ceuter ceases to power the radio cocoon (radio lobe’ iu the radio astronomical jargon). aud the observable radio cussion vanishes rapidly (Alexander&Leal.LOST:IKonüssamrov&Cubanov.1991:Veuturietal..L998:SleeRov. 1998)..," After a few ten million years the central engine in the galactic center ceases to power the radio cocoon (`radio lobe' in the radio astronomical jargon), and the observable radio emission vanishes rapidly \cite{1987MNRAS.225....1A,1994A&A...285...27K,1998MNRAS.298.1113V,1998MNRAS.297L..86S}."
+" The lifetinie of very low frequency cussion (0.01-10 MITz) roni radio galaxies should be much longer than at higher requencies, due to the stroug energv-depeudenuce of the radiative lifetimes of ultra-rclativistic electrous."," The lifetime of very low frequency emission (0.01-10 MHz) from radio galaxies should be much longer than at higher frequencies, due to the strong energy-dependence of the radiative lifetimes of ultra-relativistic electrons."
+ The electron. energv range below 100 MeV is mactically unexplored vet., The electron energy range below 100 MeV is practically unexplored yet.
+ These electron have very long ifetimes. aud they can be present a long time after ie initial radio galaxy has stopped its activity (Eubliu&Iaiser. 2000).. since both them IC radiation and Very low frequency svuchrotron enussion do not cool rem efficieutlv.," These electron have very long lifetimes, and they can be present a long time after the initial radio galaxy has stopped its activity \cite{2000A&A...360..417E}, since both their IC radiation and very low frequency synchrotron emission do not cool them efficiently."
+" Coulomb losses are also expected to be «λα, because of the very low deusitv a possible thermal component ust have m order to be consisteut with je absence of the Faraday depolarization effect in radio cocoous (Feigelsonetal... 1995)."," Coulomb losses are also expected to be small, because of the very low density a possible thermal component must have in order to be consistent with the absence of the Faraday depolarization effect in radio cocoons \cite{1995ApJ...449L.149F}. ."
+1115-N66. or in short N66 (Henize.1956).. is the largest and the most luminous rregion in the Small Magellanic Cloud (SMC).,"115-N66, or in short N66 \citep[][]{Henize56}, is the largest and the most luminous region in the Small Magellanic Cloud (SMC)."
+ It is also known as DEM $1103 (Davies et al., It is also known as DEM 103 (Davies et al.
+ 1976) and 3346. the latter referring to the bright OB association located at its center.," 1976) and 346, the latter referring to the bright OB association located at its center."
+ We refer the reader to the magnificent images of the region taken with the Advanced Camera for Surveys (ACS) onboard the and theJHK composite obtained with ISAAC on the ESO VLT telescope (Notaetal..2006:Gouliermisetal.. 2009).," We refer the reader to the magnificent images of the region taken with the Advanced Camera for Surveys (ACS) onboard the and the composite obtained with ISAAC on the ESO VLT telescope \citep[][]{Nota06,Gouliermis09}."
+. N66 is considered to be the scaled- counterpart of the Large Magellanic Cloud starburst 30 Doradus., N66 is considered to be the scaled-down counterpart of the Large Magellanic Cloud starburst 30 Doradus.
+ It indeed hosts the largest sample of young. massive stars in the whole SMC with 33 O-type stars among which 11 are of type O6.5 or earlier (Masseyetal..1989:Walborn2000:Evansetal.. 2006).," It indeed hosts the largest sample of young, massive stars in the whole SMC with 33 O-type stars among which 11 are of type O6.5 or earlier \citep[][]{Massey89,Walborn00,Evans06}."
+. It contains at least one W-R star in the massive binary or maybe triple system 55980., It contains at least one W-R star in the massive binary or maybe triple system 5980.
+ An age of .33 Myr has been estimated for 3346 from evolutionary models in the H-R diagram (Bouretetal..2003)., An age of 3 Myr has been estimated for 346 from evolutionary models in the H-R diagram \citep[][]{Bouret03}.
+. 55980 lies behind à SNR (B0057-724.Reidetal..2006:Nazé2002;Danforthetal..2003:Nazé 2004).. which has no known optical counterpart.," 5980 lies behind a SNR \citep[B0057-724,][]{Reid06,Naze02,Danforth03,Naze04}, which has no known optical counterpart."
+ Compared to the Orion Nebula. N66 has an luminosity almost 60 times higher (Kennicutt.1984).," Compared to the Orion Nebula, N66 has an luminosity almost 60 times higher \citep[][]{Kennicutt84}."
+.. This radiant flux is also reminiscent of those of giant rregions in distant metal-poor galaxies. such as regions Al and A2 in IC 4662 lying 2.44 Mpe away (Heydari-Malayertetal..1990:Crowther&Bibby.2009.andreferences therein).," This radiant flux is also reminiscent of those of giant regions in distant metal-poor galaxies, such as regions A1 and A2 in IC 4662 lying 2.44 Mpc away \citep[][and references therein]{MHM90,Crowther09}."
+" Therefore. N66 offers a valuable template for studying these kinds of distant galaxies with high resolution,"," Therefore, N66 offers a valuable template for studying these kinds of distant galaxies with high resolution."
+ Apart from its recently formed massive starpopulation mentioned above. 3346 also has a large population of low-mass. pre-main-sequence stars (Gouliermisetal... covering a mass range down to the subsolar regime.," Apart from its recently formed massive starpopulation mentioned above, 346 also has a large population of low-mass, pre-main-sequence stars \citep[][]{Gouliermis06,Nota06,Sabbi07,Hennekemper08} covering a mass range down to the subsolar regime."
+ The PMS population is found to be mainly concentrated in à number of subclusters away from the massive star association (Sabbietal..2007;Schmejaet 2000)..," The PMS population is found to be mainly concentrated in a number of subclusters away from the massive star association \citep[][]{Sabbi07,Schmeja09}."
+ The typical ages of the PMS population derived from models appear to suggest that star formation events occurred at two different epochs about 4 and 10 Myr ago (Hennekemperetal..2008)., The typical ages of the PMS population derived from models appear to suggest that star formation events occurred at two different epochs about 4 and 10 Myr ago \citep[][]{Hennekemper08}.
+. Being such an extraordinary object. the N66 complex has been extensively studied from various viewpoints using almost the whole electromagnetic range. from radio wavelengths to X-rays.," Being such an extraordinary object, the N66 complex has been extensively studied from various viewpoints using almost the whole electromagnetic range, from radio wavelengths to X-rays."
+ More specifically. the studies have used the radio continuum (Yeetal..1991;Reid2006) and the 21 em atomic line (Staveley-Smithetal...1997;Stanimirovié.al..1999).. the CO line with the SEST telescope (Rubioet 2000). the infrared with both the Infrared Space Observatory (Contursietal.2000) and the Spitzer Space Telescope (Bolattoetal..2007;Simon2007).. optical spectroscopy (Masseyetal..1989:Walbornetal.. 2006).. ultraviolet and optical spectra (Walbornetal.. 2000).. with data from theHST spectrograph STIS. AAT. and ESO 3.6 m telescope. imaging with ACS and WFC on board (Notaetal..2006).. (e.g..Kennicutt.1984;LeCoareretal.1993;Smithal. 2000).. in the UV with the International Ultraviolet Explorer (deBoer&Savage.1980) and the Far-Ultraviolet Spectroscopic Explorer (Danforthetal...2003).. and in the X-rays with XMM-Newton and Chandra (Nazéetal..2002. 2004).," More specifically, the studies have used the radio continuum \citep[][]{Ye91,Reid06} and the 21 cm atomic line \citep[][]{Staveley-Smith97,Stanimirovic99}, the CO line with the SEST telescope \citep{Rubio96,Rubio00}, the infrared with both the Infrared Space Observatory \citep{Contursi00} and the Spitzer Space Telescope \citep{Bolatto07, Simon07}, optical spectroscopy \citep[][]{Massey89,Walborn00,Evans06}, , ultraviolet and optical spectra \citep{Walborn00}, with data from the spectrograph STIS, AAT, and ESO 3.6 m telescope, imaging with ACS and WFC on board \citep[][]{Nota06}, \citep[e.g.,][]{Kennicutt84,LeCoarer93,Smith00}, in the UV with the International Ultraviolet Explorer \citep[][]{deBoer80} and the Far-Ultraviolet Spectroscopic Explorer \citep[][]{Danforth03}, and in the X-rays with XMM-Newton and Chandra \citep[][]{Naze02,Naze04}."
+ The present study is concerned with massive star formation in the N66 complex., The present study is concerned with massive star formation in the N66 complex.
+ Clustered mainly in 3346. as mentioned above. massive stars dominate the central part of the whole rregion with their strong UV radiation field.," Clustered mainly in 346, as mentioned above, massive stars dominate the central part of the whole region with their strong UV radiation field."
+ Twenty-two of the above-mentioned 33 O stars are contained in the central cluster., Twenty-two of the above-mentioned 33 O stars are contained in the central cluster.
+ The hottest star. W3. is reclassified as O2 IIK£*) (Walbornetal.. 2002a.b).," The hottest star, W3, is reclassified as O2 III(f*) \citep[][]{Walborn02a,Walborn02b}."
+. The most massive star. WI. of the central cluster. classified O4 ΠΠ (Walborn&Blades.1986).. has multiple components and the mass of the brightest component is at most 85 (Heydari-Malayert&Hutsemékers.1991).," The most massive star, W1, of the central cluster, classified O4 III(n)(f) \citep[][]{Walborn86}, has multiple components and the mass of the brightest component is at most 85 \citep[][]{MHM91}."
+. The cluster has disrupted the bulk of the natal molecular cloud. and therefore not much CO emission is detected towardsN66. except for two positions which are mapped in the (1-0) and (2-1) transitions (Rubioetal..1996. 2000)..," The cluster has disrupted the bulk of the natal molecular cloud, and therefore not much CO emission is detected towardsN66, except for two positions which are mapped in the (1-0) and (2-1) transitions \citep{Rubio96,Rubio00}. ."
+ One of these CO peaks Is associated with a remarkable feature of the whole landscape. a compact," One of these CO peaks is associated with a remarkable feature of the whole landscape, a compact"
+the code. time is expressed in units of the Alfvénn travel time. Le. Tap=L/Vap3.74 min.,"the code, time is expressed in units of the Alfvénn travel time, i.e., $\ta = L / \vap \approx 3.74$ min."
+ In addition. in all the following computations we have used p/p.=200 and po/p.=I.," In addition, in all the following computations we have used $\rp / \rc = 200$ and $\re / \rc = 1$."
+" First. we use the eigenfunction of the fundamental kink mode as the initial condition for v, at f=0."," First, we use the eigenfunction of the fundamental kink mode as the initial condition for $v_r$ at $t=0$."
+ The eigenfunction is obtained by solving the dispersion relation of the normal mode problem and computing the spatial distribution of the corresponding perturbation (seedetailsinDymova&Ruder-man2005;Soleretal.2010;Diaz 2010)..," The eigenfunction is obtained by solving the dispersion relation of the normal mode problem and computing the spatial distribution of the corresponding perturbation \citep[see details in][]{dymovaruderman, solerstatic, diazperiods}."
+ Hence. we make sure that. after the initial excitation. the magnetic tube mainly oscillates in its fundamental mode.," Hence, we make sure that, after the initial excitation, the magnetic tube mainly oscillates in its fundamental mode."
+ As a check of the numerical code. we consider the static case and put the thread at the center of the magnetic tube. 1.e.. vo= Oand zo=0.," As a check of the numerical code, we consider the static case and put the thread at the center of the magnetic tube, i.e., $v_0 = 0$ and $z_0=0$."
+" In this test simulation. we take Z,/L=0.1."," In this test simulation, we take $\lp / L = 0.1$."
+ By looking at the time-dependent evolution of νε. we check that the magnetic tube oscillates as a whole.," By looking at the time-dependent evolution of $v_r$, we check that the magnetic tube oscillates as a whole."
+" A plot of v, at 2=0 as a function of time (not displayed here for the sake of simplicity) shows that the amplitude of the oscillation ts constant during the whole duration of the simulation. meaning that numerical dissipation is negligible in the simulation."," A plot of $v_r$ at $z=0$ as a function of time (not displayed here for the sake of simplicity) shows that the amplitude of the oscillation is constant during the whole duration of the simulation, meaning that numerical dissipation is negligible in the simulation."
+" Later. we perform a power spectrum of v, at z=O (see.e.g..Car-bonell&Ballester1991) and find a large peak centered around the fundamental normal mode frequency."," Later, we perform a power spectrum of $v_r$ at $z=0$ \citep[see, e.g.,][]{carball} and find a large peak centered around the fundamental normal mode frequency."
+ The maximum of the power spectrum agrees very well with the approximate frequency of the normal mode given by Equation (16))., The maximum of the power spectrum agrees very well with the approximate frequency of the normal mode given by Equation \ref{eq:freq2}) ).
+ For the set of parameters used in this numerical test. the period in dimensional units is P.s2.6 min.," For the set of parameters used in this numerical test, the period in dimensional units is $P \approx 2.6$ min."
+" We have performed similar simulations but using different values of L,/L and co.", We have performed similar simulations but using different values of $\lp / L$ and $z_0$.
+ Equivalent results to those commented before have been obtained in all cases., Equivalent results to those commented before have been obtained in all cases.
+ This indicates. on the one hand. that the normal mode interpretation 1s a very good representation for the time-dependent evolution of the oscillation in the static case and. on the other hand. that the numerical code works properly.," This indicates, on the one hand, that the normal mode interpretation is a very good representation for the time-dependent evolution of the oscillation in the static case and, on the other hand, that the numerical code works properly."
+ Hereafter. we incorporate the effect of the flow.," Hereafter, we incorporate the effect of the flow."
+" First. we fix L,/L=0.1 and zo=0. and consider vo/va,=0.05."," First, we fix $\lp / L = 0.1$ and $z_0 = 0$, and consider $v_0 / \vap = 0.05$."
+ In this situation. the thread is initially located at the center of the magnetic tube.," In this situation, the thread is initially located at the center of the magnetic tube."
+ We make sure that the simulation stops before the threads reaches the photospheric wall., We make sure that the simulation stops before the threads reaches the photospheric wall.
+ We plot in Figure S5((a) the radial velocity perturbation at ;=O as a function of time., We plot in Figure \ref{fig:flow1}( (a) the radial velocity perturbation at $z=0$ as a function of time.
+ Contrary to the static case. now the amplitude of the oscillation decreases in time as the thread moves from the center towards the end of the magnetic tube.," Contrary to the static case, now the amplitude of the oscillation decreases in time as the thread moves from the center towards the end of the magnetic tube."
+ This behavior is qualitatively described by Equation (22)) obtained in the WKB approximation (see dashed line in Fig. 5((, This behavior is qualitatively described by Equation \ref{eq:qtempdepsol}) ) obtained in the WKB approximation (see dashed line in Fig. \ref{fig:flow1}( (
+a)). although Equation (22)) underestimates the actual decrease of the amplitude.,"a)), although Equation \ref{eq:qtempdepsol}) ) underestimates the actual decrease of the amplitude."
+ Inspired by Equation (22)). we propose the," Inspired by Equation \ref{eq:qtempdepsol}) ), we propose the"
+ , 
+"total beam averaged column density can be calculated [rom (Lucas where Q, is the rotational partition finetion. Ej is the lower state energy. Sy? is the product of the line strength and dipole moment. ancl the opacity. 7. is defined below giving where {ων is the main beam brightness temperature of the line. 7(7)) is defined as (Rolls&Wilson2000) ⋅∕⋖⋡∁⊽∐↕⊔⋟∖⊽⊔∐↲≺∢∪⋟∖⊽∐↓↕≺∢∐↓↕≺∢↕⋅∪∖∖⊽≀↧↴∖⇁≼↲∣↽≻≀↧↴≺∢↳↽↖⊂↽↔↴↕⋅∪⋯∐⇂⋖↜∃⋅⊺⋮⋅↽⊰⊔⊹⊳∶∣⋝∣⋟⋅≀↧↴∐≺⇂↕⇁⋖⊲↕⋟∖⊽⊔∐↲≺∢∪∐∐∐∏∏∐↓ brightness temperature.","total beam averaged column density can be calculated from \citep{ll93,greaves96b}
+ where $Q_r$ is the rotational partition function, $E_l$ is the lower state energy, $S\mu^2$ is the product of the line strength and dipole moment, and the opacity, $\tau$ , is defined below giving where $T_{MB}$ is the main beam brightness temperature of the line, $J(T)$ is defined as \citep{rw00}
+ $J(CB)$ is the cosmic microwave background $2.73(1+z)$ ), and $T_C$ is the continuum brightness temperature."
+" For observations not done in a temperature scale one can convert [rom Jv ! to IX using (Rohlfs&Wilson2000) where /, the is peak intensity of the line in Jv | and B is the beam filling [actor (incorporating both the continuum and absorption source factors)."," For observations not done in a temperature scale one can convert from Jy $^{-1}$ to K using \citep{rw00}
+ where $I_0$ the is peak intensity of the line in Jy $^{-1}$ and $B$ is the beam filling factor (incorporating both the continuum and absorption source factors)."
+ La all equations (he excitation temperature {ες is assumed to be 10 kx. From these equations we get an opacity (7) of 1.7 and a total column density of [ον “Lill. This value is similar to thatof (Combes&Wiklind1997)., In all equations the excitation temperature $T_{ex}$ is assumed to be 10 K. From these equations we get an opacity $\tau$ ) of 1.7 and a total column density of for $^7$ LiH. This value is similar to thatof \citep{combes97}.
+. Calculated Hs column densities for this source range from 5x107! - 5x10) 7 (Menten 1995).," Calculated $_2$ column densities for this source range from $5\times10^{21}$ - $5\times10^{23}$ $^{-2}$ \citep{menten96,combes95}."
+. From these values we caleulate a Lill/IIs ratio of 2.9x10 i 29x10P., From these values we calculate a $_2$ ratio of $2.9\times10^{-10}$ - $2.9\times10^{-12}$.
+ If we assume that the gas is cold and dark. as did Combes&Wiklind(1993).. we have [0 ) = (0.5.," If we assume that the gas is cold and dark, as did \citet{combes98}, we have $_2$ ) = 0.5."
+ TheLi/Il ratio lor a source with the redshift of ccan be estimated to be ~LO° (Combes&Wiklind 1998).., TheLi/H ratio for a source with the redshift of can be estimated to be $\sim10^{-9}$ \citep{combes98}. .
+ From (hese values we calculate a Lill/Li ratio of 0.287 - 0.003., From these values we calculate a LiH/Li ratio of 0.287 - 0.003.
+accretion svstems ancl has strong observational evidence.,accretion systems and has strong observational evidence.
+ One of the best examples comes from Ser A*. whose center harbors a supermassive black hole surrounded: by an accretion How that is likely to be in the form of the acdvection-dominated aceretion flow. (ADA. Naravan&Yi 1994)).," One of the best examples comes from Sgr A*, whose center harbors a supermassive black hole surrounded by an accretion flow that is likely to be in the form of the advection-dominated accretion flow (ADAF, \citealp{Narayan94}) )."
+ Raclio polarization observations constrain the accretion rate in the innermost region is nearly two orders of magnitude lower than that measured. at. the Bondi radius (c.g. Alarroneetal. 2006)). which indicates that intense outflows may present in this svstem.," Radio polarization observations constrain the accretion rate in the innermost region is nearly two orders of magnitude lower than that measured at the Bondi radius (e.g., \citealp{Marrone06}) ), which indicates that intense outflows may present in this system."
+ Desides. the absorption lines from highly ionized elements. which have been detected in the X-ray spectrum of some microquasars such as GRO J1655-40 (Wecdaetal.1998: Yamaoka 2001:: Milleretal. 2006)). GRS 1915|105 (Ixotaniοἱal. 2000:: Leeetal. 2002)) and Atoll sources (see e.g. the review by DíazTrigoetal.2006— and references therein). also indicate. the existence of outllows.," Besides, the absorption lines from highly ionized elements, which have been detected in the X-ray spectrum of some microquasars such as GRO J1655-40 \citealp{Ueda98}; \citealp{Yamaoka01}; \citealp{Miller06}) ), GRS 1915+105 \citealp{Kotani00}; \citealp{Lee02}) ) and Atoll sources (see e.g. the review by \citealp{Diaz Trigo06} and references therein), also indicate the existence of outflows."
+ On the other hand. Jiao&Wu(2011) found. that. outflows eencrally exist in accretion discs no matter that the Dow is advection-dominated. such as the slim cise GXbramowicz 1988)) and the ADA. or is raciation-dominated such as the standard thin disc (Shakura&Sunvaey 1973)).," On the other hand, \citet{Jiao11} found that outflows generally exist in accretion discs no matter that the flow is advection-dominated such as the slim disc \citealp{Abramowicz88}) ) and the ADAF, or is radiation-dominated such as the standard thin disc \citealp{Shakura73}) )."
+ In particular. for the three tvpes of advection-dominated lows: ADAFLEs (gas internal encrey dominant) slim disces (trapped. photon energy dominant). anc hyper-accretion disces (trapped: neutrino energy. dominant)). outflows mav be significantly strong due to positive Bernoulli parameters (ο-σ.. Naravanetal.1997: Liuetal. 2011)) or the Large radiation pressure (e.g. Gu&Lu 2007)).," In particular, for the three types of advection-dominated flows: ADAFs (gas internal energy dominant), slim discs (trapped photon energy dominant), and hyper-accretion discs (trapped neutrino energy dominant), outflows may be significantly strong due to positive Bernoulli parameters (e.g., \citealp{Narayan97}; \citealp{Liu11}) ) or the large radiation pressure (e.g., \citealp{Gu07}) )."
+ Furthermore. outllows have generally. been found. in many simulation works (e.g. Ohsuga& Mineshige.2011.— ancl references therein)," Furthermore, outflows have generally been found in many simulation works (e.g., \citealp{Ohsuga11} and references therein)."
+ In the Lwubarskiis scheme. the power spectrum. of luminosity [luctuations is sensitive to the varving mass accretion rate. thus we expect that outflows may have significant collects on the power-law index.," In the Lyubarskii's scheme, the power spectrum of luminosity fluctuations is sensitive to the varying mass accretion rate, thus we expect that outflows may have significant effects on the power-law index."
+ The paper is organized as follows., The paper is organized as follows.
+ In Section 2. we investigate the Huctuation. power spectrum under a racdius-dependent accretion rate following the method. of Lyubarski.," In Section 2, we investigate the fluctuation power spectrum under a radius-dependent accretion rate following the method of Lyubarskii."
+ The potential application of our results to observations is discussed in Section 3., The potential application of our results to observations is discussed in Section 3.
+ In the present work. following the Lyubarskii’s gencral scheme. we revisit the power spectral density of Luminosity. Iluctuations by considering a racius-dependent accretion rate.," In the present work, following the Lyubarskii's general scheme, we revisit the power spectral density of luminosity fluctuations by considering a radius-dependent accretion rate."
+ “Phe relevant. processes are the conservation of mass and angular momentum., The relevant processes are the conservation of mass and angular momentum.
+" With a radius-dependent accretion rate. these two processes are described as follows: where X. M. and T, are the surface density. accretion rate. angular velocity and re component of the stress tensor at the radius r. respectively,"," With a radius-dependent accretion rate, these two processes are described as follows: where $\Sigma $ , $\dot {M}$ , $\Omega$ and ${\rm T}_{r\varphi} $ are the surface density, accretion rate, angular velocity and $r\varphi$ component of the stress tensor at the radius $r$ , respectively."
+ 9 is the change of accretion rate over r in the stationary accretion state. Lc. 9—01ὃν. and fy is the angular momentum of d.," ${\Phi} $ is the change of accretion rate over $r$ in the stationary accretion state, i.e. ${\Phi}={\partial
+\dot {M}}/{\partial r}$, and $l_{\Phi} $ is the angular momentum of ${\Phi}$."
+" Wiüh£Q7725h, B=T2027)=Oat the inner radius Uu. we can obtain With the following definitions: the above cleseription of f° can be simplified as For standard thin dises and XADAES. QoxOy=VOAl rand thus fox(017."," With ${\Omega}r^2 = l_{\rm in} $, $F(=-{\rm T}_{r\varphi} r^2)= 0$ at the inner radius $r_{\rm in} $, we can obtain With the following definitions: the above description of $F$ can be simplified as For standard thin discs and ADAFs, $\Omega \propto \Omega_{\rm K} = \sqrt {GM / r^3}$ , and thus $h 
+\propto r^{1 / 2}$ ."
+ We assume hb=bM in this paper. where b is a constant.," We assume $h=b\sqrt {GMr}$ in this paper, where $b$ is a constant."
+ LE the accretion rate has weak dependence on the radius and faxQ1?vr. the value of Q would remain nearly constant.," If the accretion rate has weak dependence on the radius and $l_{\Phi} \propto {\Omega}r^2 \propto \sqrt{r}$, the value of $Q$ would remain nearly constant."
+ In order to simplify the problem. we assume (Q to be constant in our analysis.," In order to simplify the problem, we assume $Q$ to be constant in our analysis."
+ The validity of this assumption is cliscussed in Appendix A. With à=(10/23)[oleOn?\dr. equation (2)) becomes Substituting the above equation into the equation (1). we have The relationship between X and. £ can be deduced with a-prescription and the local balance between the heating and the cooling in the disc.," The validity of this assumption is discussed in Appendix A. With $\chi =(1/2\pi)\int {{\Phi} (l_{\Phi} - \Omega r^2)dr}$, equation \ref{Angular_conservation}) ) becomes Substituting the above equation into the equation (1), we have The relationship between $\Sigma$ and $F$ can be deduced with $\alpha $ -prescription and the local balance between the heating and the cooling in the disc."
+ The formula is similar to that with a constant accretion. rate. (FilipovLG.1984: Lyubarskit 1987: Naravan&Yi 1994)). Le. where the exponents m and n are determined by the disc model. which are discussed in 822.2. and 2 is a function of o.," The formula is similar to that with a constant accretion rate \citealp{Filipov84}; \citealp{Lyubarskii87}; \citealp{Narayan94}) ), i.e. where the exponents $m$ and $n$ are determined by the disc model, which are discussed in 2.2, and $D$ is a function of $\alpha$."
+ It should be noted thatthe analysis presented here is not strict [orADAEs., It should be noted thatthe analysis presented here is not strict forADAFs.
+X stringent analysis for the ADAESmay refer to Appendix D. Substituting equation (7)) into equation (6)). we have By assuming à=avd| 3(.r)]. where 30.rte1) is de-correlated at cdillerent racial scales and its correlation time-scale is of the order of the local viscous time-scale. the disturbed. quantities 2 anc f° are,"A stringent analysis for the ADAFsmay refer to Appendix B. Substituting equation \ref{relationship04})) into equation \ref{perturbation03}) ), we have By assuming $\alpha = \alpha _0 [1 + \overline{\beta} (t,r)]$ , where $\overline{\beta}(t,r) (\ll 1)$ is de-correlated at different radial scales and its correlation time-scale is of the order of the local viscous time-scale, the disturbed quantities $D$ and $F$ are"
+2009).,.
+. Additionally. theoretical ρω and Av can be calculated using equations (Kjeldsen&Bedding1995) and Thus. we can obtain the theoretical distributions of οι and Av by simulating the distributions of stellar mass. radius anc effective temperature.," Additionally, theoretical $\nu_{max}$ and $\Delta\nu$ can be calculated using equations \citep{kjel95}
+ and Thus, we can obtain the theoretical distributions of $\nu_{max}$ and $\Delta\nu$ by simulating the distributions of stellar mass, radius and effective temperature."
+ Using the stellar population. synthesis code TRILEGAL. Miglioetal.2009). identified that the oscillating giants observed by CoRoT are primarily red-clump stars. i.e. post-flash core-He-burning (CHeB) stars. and found that theoretical distributions of Wine and Av are good global agreement with observed ones.," Using the stellar population synthesis code TRILEGAL, \citet{migl09}
+ identified that the oscillating giants observed by CoRoT are primarily red-clump stars, i.e. post-flash core-He-burning (CHeB) stars, and found that theoretical distributions of $\nu_{max}$ and $\Delta\nu$ are good global agreement with observed ones."
+ Equations | and 2. show that the theoretical distributions of Vinee and Av are relative to the mass and radius distributions of population., Equations \ref{eq1} and \ref{eq2} show that the theoretical distributions of $\nu_{max}$ and $\Delta\nu$ are relative to the mass and radius distributions of population.
+ Miglioetal.(2009) suggested that mass- rate during the red-giant branch (RGB) and star formation rate may affect the distributions of the ρω and Av.," \cite{migl09}
+ suggested that mass-loss rate during the red-giant branch (RGB) and star formation rate may affect the distributions of the $\nu_{max}$ and $\Delta\nu$."
+ In this paper. we focused the effects of mass loss. binary interactions. the mixing-length parameter. etc.," In this paper, we focused the effects of mass loss, binary interactions, the mixing-length parameter, etc."
+ onstars., on.
+. We used the Hurley rapid single and binary evolution codes (Hurley.Pols.&Tout2000:Hurley.Tout&Pols2002) to construct stellar models.," We used the Hurley rapid single and binary evolution codes \citep{hurl00,
+hurl02} to construct stellar models."
+ In Hurlev's codes. mass-loss efficiency is an adjustable parameter.," In Hurley's codes, mass-loss efficiency is an adjustable parameter."
+ They can.thus.. be applied to studying the effect of mass loss ondistributions.," They can, be applied to studying the effect of mass loss on."
+. The paper is organized asfollows., The paper is organized as.
+. We simply show our stellar population svnthesis method in section 2., We simply show our stellar population synthesis method in section 2.
+ The results are represented in section 3., The results are represented in section 3.
+ In section 4+. we discuss and summarize the results.," In section 4, we discuss and summarize the results."
+" To simulate the stellar population observed 5, and Av (Hekkeretal. 2009). we calculated single-star stellar population (SSP) and binary-star stellar. population. (BSP). respectively."," To simulate the stellar population observed $\nu_{max}$ and $\Delta\nu$ \citep{hekk09}, we calculated single-star stellar population (SSP) and binary-star stellar population (BSP), respectively."
+ Stellar samples are generated by the Monte Carlo simulation., Stellar samples are generated by the Monte Carlo simulation.
+ The basic assumptions for the simulations are as follows. (, The basic assumptions for the simulations are as follows. (
+1) Star formation rate (SPR) is assumed to be a constant over the past 15 Gyr. (,i) Star formation rate (SFR) is assumed to be a constant over the past 15 Gyr. (
+i) The age-metallicity relation is taken from Rocha-Pintoetal.(20002).. (,ii) The age-metallicity relation is taken from \cite{roch00a}. (
+ii) The lognormal initial mass function IMF) of Chabrier(2001) is adopted.,iii) The lognormal initial mass function (IMF) of \cite{chab01} is adopted.
+ Additionally. for BSP. we generate the mass of the primary. M. according to the IMF.," Additionally, for BSP, we generate the mass of the primary, $M_{1}$, according to the IMF."
+ The ratio (q) of the mass of the secondary to that of the primary is assumed to be an uniform distribution within 0-1 for simplicity., The ratio $q$ ) of the mass of the secondary to that of the primary is assumed to be an uniform distribution within 0-1 for simplicity.
+ Then the mass of the secondary star is given by qM., Then the mass of the secondary star is given by $qM_{1}$ .
+" Weassume that all stars are members of binary systems and that the distribution of separations is constant in loge for wide binaries and falls off smoothly at close separation: where Qi,&0.070. a=αθήνα —5.7510"".0.13pc and mse1.2."," Weassume that all stars are members of binary systems and that the distribution of separations is constant in $\log a$ for wide binaries and falls off smoothly at close separation: where $\alpha_{\rm sep}\approx0.070$, $a_{\rm 0}=10R_{\odot}$, $a_{\rm 1}=5.75\times 10^{\rm 6}R_{\odot}=0.13{\rm pc}$ and $m\approx1.2$."
+ This distribution implies that the of wide binary system per logarithmic interval equal. and that approximately of the stellar systems are binary systems with orbital periods less than 100 vr (Han.Podsiadlowski.&Eggleton1995).. a," This distribution implies that the of wide binary system per logarithmic interval equal, and that approximately of the stellar systems are binary systems with orbital periods less than 100 yr \citep{HAN95}. ,"
+ssumptions.. we calculated the evolutions of 10° stars with an initial mass between 0.8 and 5.8 AJ. and mixing- parameter à = 2.0., we calculated the evolutions of $10^{6}$ stars with an initial mass between 0.8 and 5.8 $M_{\odot}$ and mixing-length parameter $\alpha$ = 2.0.
+ The theoretical ως and Av of single and binary stars are calculated using equations (1)) and (2)., The theoretical $\nu_{max}$ and $\Delta\nu$ of single and binary stars are calculated using equations \ref{eq1}) ) and \ref{eq2}) ).
+ Fig., Fig.
+ 2 shows the distributions of mass and radius of CHeB stars of simulated SSP with different Reimers (Reimers1975) mass-loss efficiency (7)., \ref{fmrd} shows the distributions of mass and radius of CHeB stars of simulated SSP with different Reimers \citep{reim75} mass-loss efficiency $\eta$ ).
+ The distributions of initial mass of progenitors of the CHeB stars also are shown in Fig. 2.., The distributions of initial mass of progenitors of the CHeB stars also are shown in Fig. \ref{fmrd}.
+ Reimers mass loss mainly affects the stars with initial mass less than 2 ΑΠ., Reimers mass loss mainly affects the stars with initial mass less than 2 $M_{\odot}$.
+ With the increase in the #7. the mass lost by a star during the RGB increases. which leads to many low-mass stars can not evolve into the CHeB stage. and the lower limit of initial mass of stars being able to evolve into the CHeB stage increases too.," With the increase in the $\eta$, the mass lost by a star during the RGB increases, which leads to many low-mass stars can not evolve into the CHeB stage, and the lower limit of initial mass of stars being able to evolve into the CHeB stage increases too."
+ For example. when the value of 7) increases from 1.0 to 2.0. the lower limit increases from about 0.95 A7. to 1.25 A7...," For example, when the value of $\eta$ increases from 1.0 to 2.0, the lower limit increases from about 0.95 $M_{\odot}$ to 1.25 $M_{\odot}$ ."
+ Thus the fraction of low-mass CHeB stars and the sample size of the simulated CHeB stars decreases with increasing η., Thus the fraction of low-mass CHeB stars and the sample size of the simulated CHeB stars decreases with increasing $\eta$.
+ For 7) = 2.0. the mass distribution is almost uniform.," For $\eta$ = 2.0, the mass distribution is almost uniform."
+ Moreover. the distribution of logarithmic radius has a fixed peak location between 1.05 and 1.15. which is almost not affected by the mass loss.," Moreover, the distribution of logarithmic radius has a fixed peak location between 1.05 and 1.15, which is almost not affected by the mass loss."
+" The distributions of v,,,.., and Av of the simulated CHeB stars are represented in Fig. 3..", The distributions of $\nu_{max}$ and $\Delta\nu$ of the simulated CHeB stars are represented in Fig. \ref{fmud}.
+ Both distributions have a dominant peak., Both distributions have a dominant peak.
+ But the peak locations are slightly lower thanones., But the peak locations are slightly lower than.
+". With the increase in 7j. thenumber of the CHeB stars with Mio, aand Av decreases."," With the increase in $\eta$ , thenumber of the CHeB stars with $\nu_{max}$ and $\Delta\nu$ decreases."
+ However. the dominant peak for Av locates between the effect of the mass loss on the Av is not enough to change the peak location.," However, the dominant peak for $\Delta\nu$ locates between the effect of the mass loss on the $\Delta\nu$ is not enough to change the peak location."
+ For ;jx |. the dominant peak of ρω locates//;Hz.," For $\eta \leq$ 1, the dominant peak of $\nu_{max}$ locates."
+. However. when the 7) increases from 1.0 to 2.0. the peak location of ως moves to Hz.," However, when the $\eta$ increases from 1.0 to 2.0, the peak location of $\nu_{max}$ moves to ."
+.. The changes in the distributions of ρα and Av in Fig., The changes in the distributions of $\nu_{max}$ and $\Delta\nu$ in Fig.
+ 3 should not be due to the typical variations in Monte Carlo simulations., \ref{fmud} should not be due to the typical variations in Monte Carlo simulations.
+" In order to show the effect of sample size on the distributions of ας and Av. we calculated the evolutions of 10"" stars which produced 3973 CHeB stars and the evolutions of LO” stars which given 793 CHeBstars."," In order to show the effect of sample size on the distributions of $\nu_{max}$ and $\Delta\nu$, we calculated the evolutions of $10^{6}$ stars which produced 3973 CHeB stars and the evolutions of $\times10^{5}$ stars which given 793 CHeBstars."
+ Both samples have thesame evolutionary parameters., Both samples have thesame evolutionary parameters.
+ Figure 4. shows that the distributions of Minos and Av of the twosamples are very similar., Figure \ref{fmtcl} shows that the distributions of $\nu_{max}$ and $\Delta\nu$ of the twosamples are very similar.
+" A Kolmogorov- test shows that the discrepanciesbetween the two sample distributionsare not significant: the Kolmogorov-Smirnov statistic D= 0.030 for ρω and 0.033 for Av. and the significance level of the DD is 0.61 for /,,,, and 0.45 for Av."," A Kolmogorov-Smirnov test shows that the discrepanciesbetween the two sample distributionsare not significant: the Kolmogorov-Smirnov statistic $D \simeq$ 0.030 for $\nu_{max}$ and 0.033 for $\Delta\nu$ , and the significance level of the $D$ is 0.61 for $\nu_{max}$ and 0.45 for $\Delta\nu$ ."
+ In Fig. 5..," In Fig. \ref{faged}, ,"
+ we plotted the distributions of, we plotted the distributions of
+Accretion dises play an important role in many astrophysical situations. from protostellar systems through to disces around supermassive black holes in active galactic nuclei (AGN).,"Accretion discs play an important role in many astrophysical situations, from protostellar systems through to discs around supermassive black holes in active galactic nuclei (AGN)."
+ What is still very uncertain is the process through which angular momentum is transported. outwards in such dises., What is still very uncertain is the process through which angular momentum is transported outwards in such discs.
+ It is clear. from observations of accretion rates. that classical hydrodynamical viscosity is insufficient to play this role.," It is clear, from observations of accretion rates, that classical hydrodynamical viscosity is insufficient to play this role."
+ The typical solution is to assume an parametrisation of the viscosity. whose origin is not well understood.," The typical solution is to assume an parametrisation of the viscosity, whose origin is not well understood."
+ The archetype is the a@-parametrisation introduced by ?. in which the viscosity v is assumed to depend on the dise sound speed. c.. and thickness. //. through v=oc.4/. where a«1.," The archetype is the $\alpha$ -parametrisation introduced by \citet{Shakura_Sunyaev_73} in which the viscosity $\nu$ is assumed to depend on the disc sound speed, $c_s$, and thickness, $H$, through $\nu = \alpha c_s H$, where $\alpha << 1$."
+ This allows a number of different physical mechanisms o be considered as the origin of this viscosity., This allows a number of different physical mechanisms to be considered as the origin of this viscosity.
+ The most rTequently invoked is turbulent viscosity. shifting the problem to he origin of the turbulence.," The most frequently invoked is turbulent viscosity, shifting the problem to the origin of the turbulence."
+ If the dise is sufficiently ionised. the magnetorotational instability (222) can result in turbulence that can oovide the required viscosity.," If the disc is sufficiently ionised, the magnetorotational instability \citep{Balbus1991,Balbus1999, Papaloizou2003} can result in turbulence that can provide the required viscosity."
+ However. if the dise is very weakly ionised (as in the case of most protostellar disces at early times). another source must be sought.," However, if the disc is very weakly ionised (as in the case of most protostellar discs at early times), another source must be sought."
+ During the earliest stages of star ormation. when dise masses are likely to be high relative to the mass of the centralprotostar. gravitational instabilities (GI) may orovide the answer (22).," During the earliest stages of star formation, when disc masses are likely to be high relative to the mass of the centralprotostar, gravitational instabilities (GI) may provide the answer \citep{Lin1987,Laughlin1994}. ."
+use similar cross-correlation techniques to those employed here to search for time lags between the soft band KkeV) light curve and the hard band kKkeV) light curve.,use similar cross-correlation techniques to those employed here to search for time lags between the soft band keV) light curve and the hard band keV) light curve.
+ They found that the hard band lags the soft band by kksec., They found that the hard band lags the soft band by $1.6\pm 0.5$ ksec.
+ In this work. we find lags £; 1000ss between bands El kkeV) and E4 kkeV).," In this work, we find lags $\approxlt 1000$ s between bands E1 keV) and E4 keV)."
+ Noting that thermal Comptonization predicts a time lag that varies logarithmically with energy. our results are consistent with those of Nowak Chiang.," Noting that thermal Comptonization predicts a time lag that varies logarithmically with energy, our results are consistent with those of Nowak Chiang."
+ Reynolds (1999) uses an interpolation method to constrain trial transfer functions linking two given bands., Reynolds (1999) uses an interpolation method to constrain trial transfer functions linking two given bands.
+ He found a lag of 50-100ss between the 2—+kkeV and 8-ISkkeV bands. rather smaller than that found here.," He found a lag of s between the keV and keV bands, rather smaller than that found here."
+ To reconcile these results. ohe Must appreciate that these methods probe lags at different Fourier frequencies.," To reconcile these results, one must appreciate that these methods probe lags at different Fourier frequencies."
+ The CCF methods tend to probe lags across a broad spectrum of Fourier frequencies., The CCF methods tend to probe lags across a broad spectrum of Fourier frequencies.
+ Due to the red nature of the power spectrum. such methods are naturally biased towards the lower Fourier frequencies.," Due to the red nature of the power spectrum, such methods are naturally biased towards the lower Fourier frequencies."
+ The method of Reynolds (1999). instead. probes the higher Fourier frequencies since he uses fairly spiky trial transfer functions.," The method of Reynolds (1999), instead, probes the higher Fourier frequencies since he uses fairly spiky trial transfer functions."
+ Hence. to paraphrase these technical results. the rapid flickering seems to get transmitted up the observed energy spectrum with a smaller time lag (by an order of magnitude) than experienced by the slower variations.," Hence, to paraphrase these technical results, the rapid flickering seems to get transmitted up the observed energy spectrum with a smaller time lag (by an order of magnitude) than experienced by the slower variations."
+ While spectral studies of X-ray variability in time sequence may hold the key to understanding the underlying processes that are responsible for producing the observed dramatic flux changes. it is insufficient for constraining the size of the emitting region in the absence of a good understanding of the flare mechanisms.," While spectral studies of X-ray variability in time sequence may hold the key to understanding the underlying processes that are responsible for producing the observed dramatic flux changes, it is insufficient for constraining the size of the emitting region in the absence of a good understanding of the flare mechanisms."
+ The line profile obtained from observations suggest that the X-ray emission originates from ~ 10-20 gravitational radii of the black hole., The line profile obtained from observations suggest that the X-ray emission originates from $\sim$ 10-20 gravitational radii of the black hole.
+ This together with a periodic signal can constrain the size of the emitting region., This together with a periodic signal can constrain the size of the emitting region.
+ Alternatively. we can attempt to estimate the black hole mass by assessing where the break frequency in the PDS oceurs (assuming that the break frequency scales with mass).," Alternatively, we can attempt to estimate the black hole mass by assessing where the break frequency in the PDS occurs (assuming that the break frequency scales with mass)."
+ We calculate the power density spectrum using the Lomb-Seargle method (Lomb 1976: Scargle 1982: Press et al., We calculate the power density spectrum using the Lomb-Scargle method (Lomb 1976; Scargle 1982; Press et al.
+ 1992) appropriate for unevenly sampled data. and fit a power law slope independently tothe and data between LO7 and 10 Hz. (," 1992) appropriate for unevenly sampled data, and fit a power law slope independently tothe and data between $10^{-5}$ and $10^{-4}$ Hz. ("
+We ignore data above the latter in order to avoid contamination to the tit from the 96 minute orbital period of .),We ignore data above the latter in order to avoid contamination to the fit from the 96 minute orbital period of .)
+ Figs., Figs.
+" I8. show that respectively. fo? and f|"" is sufficiently representative of the and data down to ~45.10 ""Hz. where the break in the power spectrum may occur. ("," \ref{fig19-lomb} show that respectively, $f^{-1.3}$ and $f^{-1.5}$ is sufficiently representative of the and data down to $\sim$ $4-5 \times 10^{-6}$ Hz, where the break in the power spectrum may occur. ("
+"We note that this value is only given as a limit tofi, - while there appears to be no additional evidence for a break below 5.10 ""Hz.the observations are insufficiently long to claim a definitive determination.)","We note that this value is only given as a limit to$f_{br}$ - while there appears to be no additional evidence for a break below $4-5 \times 10^{-6}$ Hz, the observations are insufficiently long to claim a definitive determination.)"
+ This is consistent with the findings of Hayashida et al. (, This is consistent with the findings of Hayashida et al. (
+1998) for 6-30-15.,1998) for $-$ 6-30-15.
+ Since the count rate throughout these observations remain steadily between 6-24 cts ! for and 0.5-3 cts! for. the findings above woulc ραuggest that large scale power does not exist in abundance even ct10ugh much shorter time scale variability is highly evident.," Since the count rate throughout these observations remain steadily between 6-24 ct $\rm
+s^{-1}$ for and 0.5-3 ct $\rm s^{-1}$ for, the findings above would suggest that large scale power does not exist in abundance even though much shorter time scale variability is highly evident."
+" We note that Papadakis Lawrence (1993) caution agains standard Fourier analysis techniques for estimating power spectra. in the form of a bias of the periodogram due to a windowing effect. in addition to a possible ""red noise leak’."," We note that Papadakis Lawrence (1993) caution against standard Fourier analysis techniques for estimating power spectra, in the form of a bias of the periodogram due to a windowing effect, in addition to a possible `red noise leak'."
+ The former is tied to a concern that the sampling window function can alias power from frequencies above and below the central frequency (ο. thereby distorting the true shape of the power spectrum: the latter ‘red noise leak” effect is concerned with a transfer of power from low to high frequencies. (," The former is tied to a concern that the sampling window function can alias power from frequencies above and below the central frequency $\omega_p$, thereby distorting the true shape of the power spectrum; the latter `red noise leak' effect is concerned with a transfer of power from low to high frequencies. ("
+"The problem of the ""red noise leak’ was first noted by Deeter Boyton (1982) and Deeter (1984).)",The problem of the `red noise leak' was first noted by Deeter Boyton (1982) and Deeter (1984).)
+ However. we conclude according to subsequent equations (6) and (7) that the effects of this bias is negligible for the data sets in question.," However, we conclude according to subsequent equations (6) and (7) that the effects of this bias is negligible for the data sets in question."
+ The expected value of the periodogram is defined such that: where f'y(wwy) is theKernel (Priestley 1981) which assumes the shape of the s/nc function for frequencies  iy., The expected value of the periodogram is defined such that: where $F_N(\omega - \omega_p)$ is the (Priestley 1981) which assumes the shape of the $sine^2$ function for frequencies $\sim \omega_p$ .
+ Itfollows that the bias is : For lengthy time series (e.g. 400 ks). the mean value of the," Itfollows that the bias is : For lengthy time series (e.g. 400 ks), the mean value of the"
+wilh a prototvpe (ransilion-edge sensor detector.,with a prototype transition-edge sensor detector.
+ Llowever. it is not clear whether the timing calibration of the instrument was complete at the time.," However, it is not clear whether the timing calibration of the instrument was complete at the time."
+ Ulmeretal.(1994). presented results from OSSE observations (50-100 keV). indicating that the hard. X-ray. (o radio lag was «30-30 fis. Ht would appear that the estimate of their errors was too optimistic.," \citet{ulme1994} presented results from OSSE observations (50-100 keV), indicating that the hard X-ray to radio lag was $<$ $\pm$ 30 $\mu$ s. It would appear that the estimate of their errors was too optimistic."
+ Nolanetal.(1993). present pulse profiles but no absolute phases., \citet{nola1993} present pulse profiles but no absolute phases.
+ Kuiperetal.(2003) report on INTEGRAL data. covering 6-50 keV. and deriving atime delay (the radio trailing) of 2802-40 pis for a single epoch.," \citet{kuip2003} report on INTEGRAL data, covering 6-50 keV, and deriving atime delay (the radio trailing) of $\pm$ 40 $\mu$ s for a single epoch."
+ The precise “imine of the pulses in the different wavelength regimes has important repercussions for the understaiding of (he nature aud spatial origin of the emission processes that give rise to the pulses in different parts of the spectrum., The precise timing of the pulses in the different wavelength regimes has important repercussions for the understanding of the nature and spatial origin of the emission processes that give rise to the pulses in different parts of the spectrum.
+ Romani have suggested (hat. while the radio precursor comes [rom the polar cap. the pulse and interpulse originate in (he outer gap in the magnetosphere. with the higher energv pulses being generated al significantly greater height.," \citet{roma1995} have suggested that, while the radio precursor comes from the polar cap, the pulse and interpulse originate in the outer gap in the magnetosphere, with the higher energy pulses being generated at significantly greater height."
+ Thus. measuring the pulse shapes aud the absolute timing throughout the electromagnetic spectrum places important constraints on (he shape of the the outer gap and on the height in (he magnetosphere at which the radiation is generatec.," Thus, measuring the pulse shapes and the absolute timing throughout the electromagnetic spectrum places important constraints on the shape of the the outer gap and on the height in the magnetosphere at which the radiation is generated."
+ In (his paper we will present the results of the INTE monitoring campaign of the Crab Pulsar from the start of that mission., In this paper we will present the results of the RXTE monitoring campaign of the Crab Pulsar from the start of that mission.
+ We have adopted the radio nomenclature for the features in the pulse profile (main pulse. bridge. aud interpulse).," We have adopted the radio nomenclature for the features in the pulse profile (main pulse, bridge, and interpulse)."
+ The observations presented here were made as part of an on-going monitoring campaign of the Rossi X-rav. Timing Explorers (RATE)., The observations presented here were made as part of an on-going monitoring campaign of the Rossi X-ray Timing Explorer's (RXTE).
+ Observations of about. 1000 s in duration were initially made at weekly. later bi-weekly. intervals. with the exception of a period from mid-Max until mid-July when the Crab pulsar is too close to the sun.," Observations of about 1000 s in duration were initially made at weekly, later bi-weekly, intervals, with the exception of a period from mid-May until mid-July when the Crab pulsar is too close to the sun."
+ In (his paper we shall use the events in the 2-16 keV range. collected with the RATE’s Proportional Counter Array (PCA) in 177 observations between MJD. 50129 and. 52941.," In this paper we shall use the events in the 2-16 keV range, collected with the RXTE's Proportional Counter Array (PCA) in 177 observations between MJD 50129 and 52941."
+ For the first four ancl a half vears the observations were made in pulsar fold mode. with approximately 80 bins per period.," For the first four and a half years the observations were made in pulsar fold mode, with approximately 80 bins per period."
+ Halhwav through the filth vear this was changed (o an event mode with 250 jes time resolution., Halfway through the fifth year this was changed to an event mode with 250 $\mu$ s time resolution.
+ The accuracy of the RNTE clock in absolute time is about 8 prs for data taken before 29 April 1997 (MJD 50561): see Rotsetal.(1998b)., The accuracy of the RXTE clock in absolute time is about 8 $\mu$ s for data taken before 29 April 1997 (MJD 50567); see \citet{rots1998b}.
+. After that date. the error decreased to 2 js (Markwirdt 2003)..," After that date, the error decreased to 2 $\mu$ s \citep{mark2003}. ."
+ This accuracy can be achieved by applving the fine clock corrections, This accuracy can be achieved by applying the fine clock corrections
+where Oy is the observed. Lux in the wavelength. bin À. Ady is the moclel Hux in (1)). ancl wey is the inverse of the noise in the measurement Oy.,"where $O_{\lambda}$ is the observed flux in the wavelength bin $\lambda$, $M_{\lambda}$ is the model flux in \ref{cfmodel}) ), and $w_{\lambda}$ is the inverse of the noise in the measurement $O_{\lambda}$ ."
+ Theo summation is over the Ny wavelength bins in the observed. spectrum., The summation is over the $N_{\lambda}$ wavelength bins in the observed spectrum.
+ “Phe minimization of (2)) is performed over A|4+ parameters: wpety.AMxuiilescoss and ess where A. is the number of basis spectra in £.," The minimization of \ref{chisq}) ) is performed over $K+4$ parameters: $x_1,...,x_K, M_{\lambda_0}, A_V, v_*$, and $\sigma_*$, where $K$ is the number of basis spectra in $\mathbf{\xi}$."
+ The speed of the algorithini scales as A7., The speed of the algorithm scales as $K^2$.
+ For the analysis of. ce. 210° galaxy. spectra in the SDSS database. it is crucial to pick abasis with a small number of spectra.," For the analysis of, e.g. $\simeq 10^6$ galaxy spectra in the SDSS database, it is crucial to pick abasis with a small number of spectra."
+ In. for example. CidFernandesetal.(2004) and CidFernandesetal.(2005).. the STABRLIGLIT model— (1) and fitting algorithm are tested using simulations.," In, for example, \citet{CF2004} and \citet{CF2005}, the STARLIGHT model \ref{cfmodel}) ) and fitting algorithm are tested using simulations."
+ The simulated spectra are generated from. the model in. (1)) using dy=30 SSPs from Tremonti(2003) and A=45 SSPs from Bruzual&Charlot(2003)).. respectively.," The simulated spectra are generated from the model in \ref{cfmodel}) ) using $K=30$ SSPs from \citet{Tre2003} and $K=45$ SSPs from \citet{BC03}, respectively."
+ These authors lit the simulations using the same basis of SSPs that was used to generate the simulations., These authors fit the simulations using the same basis of SSPs that was used to generate the simulations.
+ From their analyses. they conclude that in the absence of noise. the algorithm accurately recovers the input. parameters.," From their analyses, they conclude that in the absence of noise, the algorithm accurately recovers the input parameters."
+ In the presence of noise. the population vector X is not recovered. ancl it is advised that à time-binned description of the SELL be adopted.," In the presence of noise, the population vector $\mathbf{x}$ is not recovered, and it is advised that a time-binned description of the SFH be adopted."
+ Although their use of the same basis for both generating and fitting simulated. spectra is an appropriate test that the algorithm works. it is not a fair assessment of the expecte performance of the methods for a database of real galaxy spectra.," Although their use of the same basis for both generating and fitting simulated spectra is an appropriate test that the algorithm works, it is not a fair assessment of the expected performance of the methods for a database of real galaxy spectra."
+ In. 82.1.2 we discuss our concerns with their simulation assessments., In \ref{models} we discuss our concerns with their simulation assessments.
+ In 82.2. we present a new suite of methods that we test with several sets of simulated galaxy spectra., In \ref{basis} we present a new suite of methods that we test with several sets of simulated galaxy spectra.
+ These simulations. which are described in 83.. are designed to be more representative of a large database of galaxy spectra. such as SDSS.," These simulations, which are described in \ref{simulations}, are designed to be more representative of a large database of galaxy spectra, such as SDSS."
+ In reality. the population of all observed. galaxies is not constrained to be mixtures of a small number of SSP's on a discrete. erid of age and metallicity.," In reality, the population of all observed galaxies is not constrained to be mixtures of a small number of SSPs on a discrete grid of age and metallicity."
+ A more physically accurate representation of galaxies is as mixtures of an infinitely large. basis of SSPs on a continuous grid. of age and metallicity and. depending on the complexity of the uncerlving physics. a possibly infinite &rid of prescriptions for initial mass function and evolutionary track.," A more physically accurate representation of galaxies is as mixtures of an infinitely large basis of SSPs on a continuous grid of age and metallicity and, depending on the complexity of the underlying physics, a possibly infinite grid of prescriptions for initial mass function and evolutionary track."
+ ln. our simulations in 83. we simulate galaxies from a Large database ol ~10° SSPs on a Line grid of age and metallicity to determine whether we can accurately describe. their SEL. metallicities. and kinematic parameters using an appropriately chosen. computationally tractable basis of fy=45 or 150 spectra.," In our simulations in \ref{simulations} we simulate galaxies from a large database of $\sim 10^3$ SSPs on a fine grid of age and metallicity to determine whether we can accurately describe their SFHs, metallicities, and kinematic parameters using an appropriately chosen, computationally tractable basis of $K=45$ or 150 spectra."
+ We adopt a SSP database containing 1278 spectra., We adopt a SSP database containing 1278 spectra.
+ This set of spectra is used to both choose an appropriate basis (82.2)) and generate simulated spectra 3)., This set of spectra is used to both choose an appropriate basis \ref{basis}) ) and generate simulated spectra \ref{simulations}) ).
+ The database is meant to represent the [large population of SSPs from which observed. galaxies can be composed., The database is meant to represent the large population of SSPs from which observed galaxies can be composed.
+" We use the spectra from the models of Bruzual (2003).. computed for a Chabrier(2003). ME. ""Padova 1994. evolutionary tracks (Aloneietal.1903:Bres-sanetal.1993:Fagotto1994a.b:Cirardi 1996)). and STELLID library (LeBorgneetal. 2003))."," We use the spectra from the models of \citet{BC03}, computed for a \citet{Cha2003} IMF, `Padova 1994' evolutionary tracks \citealt{Alo1993,Bre1993,Fag1994a,Fag1994b,Gir1996}) ), and STELIB library \citealt{LeB2003}) )."
+ Pheoretically. we could also vary the IME and evolutionary tracks across spectra and. use the same methods developed in 82.2.. but for simplicity decide to fix them in this study.," Theoretically, we could also vary the IMF and evolutionary tracks across spectra and use the same methods developed in \ref{basis}, but for simplicity decide to fix them in this study."
+ The SSPs in our database are generated on a grid of 213 approximately evenly log-spaced time bins from 0 to 18 Civrs and 6 dilferent metallicities: Z = 0.0001. 0.0004. 0.004. 0.008. 0.02 and 0.05. where Z is the fraction of the mass of a star composed. of metals (Z.=0.02).," The SSPs in our database are generated on a grid of 213 approximately evenly log-spaced time bins from 0 to 18 Gyrs and 6 different metallicities: $Z$ = 0.0001, 0.0004, 0.004, 0.008, 0.02 and 0.05, where $Z$ is the fraction of the mass of a star composed of metals $Z_{\odot}=0.02$ )."
+ This grid is designed to approximate a continuous variation across age and Z., This grid is designed to approximate a continuous variation across age and $Z$.
+ The coarseness in the current Z grid is due to limitations in the current stellar population models., The coarseness in the current $Z$ grid is due to limitations in the current stellar population models.
+ Again. a finer grid of age and Z can be implemented with the techniques that we develop in ," Again, a finer grid of age and $Z$ can be implemented with the techniques that we develop in \ref{basis}."
+In previous studies using the S'TATULIGIUPE code. researchers hancd-selected sets of SSP spectra to use in the mocel fitting.," In previous studies using the STARLIGHT code, researchers hand-selected sets of SSP spectra to use in the model fitting."
+" According to Cid.Fernandesctal.(2004).. ""the elements of the base should span the range of spectral properties observed. in. the sample galaxies ancl provide enough resolution in age and metallicity to aclress the desired scientific questions.”"," According to \citet{CF2004}, “the elements of the base should span the range of spectral properties observed in the sample galaxies and provide enough resolution in age and metallicity to address the desired scientific questions.”"
+ In this same paper. they claim that their basis of dv=30 SSPs adequately recovers parameters of age and metallicity in their simulations.," In this same paper, they claim that their basis of $K=30$ SSPs adequately recovers parameters of age and metallicity in their simulations."
+ However. as discussed in 82.1.. these simulations are unrealistic since thev employ the same basis to both generate and fit the simulated spectra.," However, as discussed in \ref{specsyn}, these simulations are unrealistic since they employ the same basis to both generate and fit the simulated spectra."
+ In this section we present methods to numerically determine a basis £ of à small number ofspechra.. £j. from a large set of SSP spectra.," In this section we present methods to numerically determine a basis $\mathbf{\xi}$ of a small number of, $\xi_j$, from a large set of SSP spectra."
+ “Phe resultant basis of prototype spectra will be designed to capture a large proportion of the variation of the set of SSPs. which is chosen to span the range of observed. spectral properties of a data set of galaxies at a high resolution in age and metallicity. such as the database of 1278 SSPs described in 2.1.2...," The resultant basis of prototype spectra will be designed to capture a large proportion of the variation of the set of SSPs, which is chosen to span the range of observed spectral properties of a data set of galaxies at a high resolution in age and metallicity, such as the database of 1278 SSPs described in \ref{models}."
+ The numerically chosen basis of prototype spectra. £. should recover physical parameters better than any hand-chosen basis of the same size because hand-selected bases ignore subsets of the parameter space and may cause the solution to be degenerate by. includingmultiple SSPs with nearly identical spectra.," The numerically chosen basis of prototype spectra, $\mathbf{\xi}$, should recover physical parameters better than any hand-chosen basis of the same size because hand-selected bases ignore subsets of the parameter space and may cause the solution to be degenerate by includingmultiple SSPs with nearly identical spectra."
+ In refsimulations we show that the bases computed from our numerical methods outperform hancl-sclected bases: used in the literature in. parameter estimation for simulated ealaxies., In \\ref{simulations} we show that the bases computed from our numerical methods outperform hand-selected bases used in the literature in parameter estimation for simulated galaxies.
+ Dillusion map is a popular new technique for data parametrization and dimensionality reduction that has been developed. by Coifman&Lafon(2006) and (2006).. ancl was recently used by Richardsetal.(2009) to analvse SDSS spectra and Freemanctal.(2000). to estimate photometric redshifts.," Diffusion map is a popular new technique for data parametrization and dimensionality reduction that has been developed by \citet{Coif2006} and \citet{LafonLee2006}, and was recently used by \citet{Rich2009} to analyse SDSS spectra and \citet{Free2009} to estimate photometric redshifts."
+ As shown in these latter works. dilfusion map has the powerful ability to uncover simple structure in complicated. high-dimensional data.," As shown in these latter works, diffusion map has the powerful ability to uncover simple structure in complicated, high-dimensional data."
+ These analyses demonstrate that thissimple structure can be used to make precise statistical inferences about the data., These analyses demonstrate that thissimple structure can be used to make precise statistical inferences about the data.
+ In the problem at hand. we begin with a set B of high- broac-banc SSP spectra b;./= 1... where AN can be very [large ancl each b; has llux measurementsN. in pXO wavelength bins.," In the problem at hand, we begin with a set $\mathbf{B}$ of high-resolution broad-band SSP spectra $\mathbf{b}_{i}, i=1,...,N$ , where $N$ can be very large and each $\mathbf{b}_{i}$ has flux measurements in $p \gtrsim
+10^3$ wavelength bins."
+ Prom this set. our goal is to derive a small basis of ἐν prototype spectra. £=[£1.£y P. cach," From this set, our goal is to derive a small basis of $K$ prototype spectra, $\mathbf{\xi} =
+\{\mathbf{\xi}_{1},...,\mathbf{\xi}_{K}\}$ , each"
+Simultaneous multi-frequency observations carrieck out at clilferent epochs are. necessary to monitor the spectral behaviour and variability of high-frequcney peaking racio sources.,Simultaneous multi-frequency observations carried out at different epochs are necessary to monitor the spectral behaviour and variability of high-frequency peaking radio sources.
+ Variations in the spectral properties. like peak frequeney. spectral shape and Bux density. are. strong indicators of the true nature of the source (Orientietal.2007:Tingav&delxool 2008).," Variations in the spectral properties, like peak frequency, spectral shape and flux density, are strong indicators of the true nature of the source \citep{mo07,
+ tingay03}."
+. ln voung radio sources the spectral properties are not. expected. to change. while spectral variability is a twpical characteristic of beamed racio To determine a possible variation of the spectral peak. for cach source we fitted the simultaneous radio spectrum at each epoch with a pure analytical function: where S ds the [flux density. £6 the [οσον and5 ande are numeric parameters without any direct. physical meaning.," In young radio sources the spectral properties are not expected to change, while spectral variability is a typical characteristic of beamed radio To determine a possible variation of the spectral peak, for each source we fitted the simultaneous radio spectrum at each epoch with a pure analytical function: where $S$ is the flux density, $\nu$ the frequency, and and are numeric parameters without any direct physical meaning."
+ We prefer to adopt. this function. instead. of hat used by Stanghellinietal.(2009) because it better represents the data. providing more accurate values for the »vak parameters.," We prefer to adopt this function instead of that used by \citet{cs09} because it better represents the data, providing more accurate values for the peak parameters."
+ The best fits to the spectra are shown in Fig. 5..," The best fits to the spectra are shown in Fig. \ref{radio_spettri},"
+ and the derived. peak frequencies at the various epochs are reported in Table 5.., and the derived peak frequencies at the various epochs are reported in Table \ref{variability}.
+ Statistical errors. derived rom the fit are not representative of the real uncertainty on the estimate of the peak frequency., Statistical errors derived from the fit are not representative of the real uncertainty on the estimate of the peak frequency.
+ For this reason we orefer to assume a conservative uncertainty on the peak requeney of10:4., For this reason we prefer to assume a conservative uncertainty on the peak frequency of.
+ The position of the spectral peak is well constrained when the peak occurs at a frequency well sampled. by the observations. becoming less accurate when he frequeney coverage is not as appropriate.," The position of the spectral peak is well constrained when the peak occurs at a frequency well sampled by the observations, becoming less accurate when the frequency coverage is not as appropriate."
+ For example. in the case à source has a spectrum peaking at the edges of the frequeney. coverage (ic. Lo or Ix/€O) bands). the fit provides parameters that are less constrained. than in the case of sources with the spectral peak occurring. around 5-10 6111. where both the optically thin and thick emission are properly sampled.," For example, in the case a source has a spectrum peaking at the edges of the frequency coverage (i.e. L or K/Q bands), the fit provides parameters that are less constrained than in the case of sources with the spectral peak occurring around 5-10 GHz, where both the optically thin and thick emission are properly sampled."
+ Spectra poorly constrained. are also those lacking observations at some frequencies. as in the case of J1044|4328. J1052|3355. and J1547|3518 where the lack of data either at 1.4/1.7 or 15.3 GlIz preclucled a reliable determination of their peak frequencies.," Spectra poorly constrained are also those lacking observations at some frequencies, as in the case of J1044+4328, J1052+3355, and J1547+3518 where the lack of data either at 1.4/1.7 or 15.3 GHz precluded a reliable determination of their peak frequencies."
+ ὃν comparing the clistribution of the peak frequency of all the sources at the various epochs. we do not. find remarkable dillerences: the changes are usually within the uncertainties.," By comparing the distribution of the peak frequency of all the sources at the various epochs, we do not find remarkable differences; the changes are usually within the uncertainties."
+" The median value of the peak frequency of the whole sample at cach epoch has notchanged significantly: wm,=5.8 Gllz in the 1998-1999 epoch (namelytheobser-selectingthefaintLIEDsample:Stanghellinietal. 2009):: fA,=6.2 Gllz in the 2003-2004 epoch: and ον=5.4 Giz in the 2006-2007 epoch."," The median value of the peak frequency of the whole sample at each epoch has notchanged significantly: $\nu_{\rm p} = 5.8$ GHz in the 1998-1999 epoch \citep[namely the observations of the HFP candidates that have been
+ used for selecting the faint HFP sample;][]{cs09}; $\nu_{\rm p} = 6.2$ GHz in the 2003-2004 epoch; and $\nu_{\rm p} = 5.4$ GHz in the 2006-2007 epoch."
+ Among the sources studied here. a few cases show significant variation: in 7 sources (1053|4610. J105813353. J1203]|4803. J1251|4317. J1258|2520. J1300|4352. ancl J1616|4632) the peak measured. in the most recent epoch has substantially shifted: towards either higher or lower frequencies.," Among the sources studied here, a few cases show significant variation: in 7 sources (J1053+4610, J1058+3353, J1203+4803, J1251+4317, J1258+2820, J1300+4352, and J1616+4632) the peak measured in the most recent epoch has substantially shifted towards either higher or lower frequencies."
+" This result indicates that these sources are part of the blazar population. since in voung racio sources the shift of the peak towards lower frequencies is not expected to be so remarkable at least on this short time scale (see Section We computed. the spectral index both below (e) and above (04) the peak frequency. by fitting a straight line in the optically-thick ancl -thin part of the spectrum. respectively. following the approach by ""Torniainenetal. ancl Orientietal. (2007).."," This result indicates that these sources are part of the blazar population, since in young radio sources the shift of the peak towards lower frequencies is not expected to be so remarkable at least on this short time scale (see Section We computed the spectral index both below $\alpha_{\rm b}$ ) and above $\alpha_{\rm a}$ ) the peak frequency, by fitting a straight line in the optically-thick and -thin part of the spectrum, respectively, following the approach by \citet{torniainen05} and \citet{mo07}. ."
+ We considered. αι, We considered “flat”
+The intrinsic alignments of galaxies have been searched observationally.,The intrinsic alignments of galaxies have been searched observationally.
+ The SuperCOSAIOS data on nearby galaxies (with median redshift 2~0.1) reveal a level of LO? to 10.+ on the correlation of intrinsic ellipticilies of galaxies over the angular scales of a lew tens ol areminutes (Brown et al., The SuperCOSMOS data on nearby galaxies (with median redshift $z\sim 0.1$ ) reveal a level of $10^{-5}$ to $10^{-4}$ on the correlation of intrinsic ellipticities of galaxies over the angular scales of a few tens of arcminutes (Brown et al.
+ 2002)., 2002).
+ Analyses on close pairs of galaxies from. COMDO-17 data (with z~ 0.6) find the intrinsic alignments to be consistent with zero but with uncertainties on the order of αfewx107 on scales of a few arcminutes (IIevians οἱ al., Analyses on close pairs of galaxies from COMBO-17 data (with $z\sim 0.6$ ) find the intrinsic alignments to be consistent with zero but with uncertainties on the order of $a\ few \times 10^{-4}$ on scales of a few arcminutes (Heymans et al.
+ 2004)., 2004).
+ Investigations on SDSS main sample with z0.1 and ils subsamples with z~0.07 lo zeQ.21 conclude that no significant intrinsic alignments are detected (Mandelbaium el al., Investigations on SDSS main sample with $z\sim 0.1$ and its subsamples with $z\sim 0.07$ to $z\sim 0.21$ conclude that no significant intrinsic alignments are detected (Mandelbaum et al.
+ 2006)., 2006).
+ The observational results on the intrinsic alignments are close to the lower limits given bv different. numerical simulations on dark matter halos (e.g.. IIevians οἱ al.," The observational results on the intrinsic alignments are close to the lower limits given by different numerical simulations on dark matter halos (e.g., Heymans et al."
+ 2006)., 2006).
+ It has been pointed out that the existence of misalionment between barvonic matter and dark matter can significantly reduce the intrinsic alignment. of background. galaxies in comparison with that of dark matter halos and may explain (he low observational results founcl in different survevs (Hevimans et al., It has been pointed out that the existence of misalignment between baryonic matter and dark matter can significantly reduce the intrinsic alignment of background galaxies in comparison with that of dark matter halos and may explain the low observational results found in different surveys (Heymans et al.
+ 2006)., 2006).
+ In (his paper. we study the effects of the intrinsic alienment on finding mass concentrations through weak lensing effects.," In this paper, we study the effects of the intrinsic alignment on finding mass concentrations through weak lensing effects."
+ Being the largest virialized objects in the universe. clusters of ealaxies are important cosmological probes because (their formation aud evolution depend sensitivelv on cosmologies (e.g.. Dorgani 2006: Fan Chiueth 2001).," Being the largest virialized objects in the universe, clusters of galaxies are important cosmological probes because their formation and evolution depend sensitively on cosmologies (e.g., Borgani 2006; Fan Chiueth 2001)."
+ However. large uncertainties exist in linking cluster observables. such as galaxy richness. X-ray brightuess and Zeldovich οδοί. to their mass. the important quantity in cosmological analvses (e.g.. Doce οἱ al.," However, large uncertainties exist in linking cluster observables, such as galaxy richness, X-ray brightness and Sunyaev-Zeldovich effect, to their mass, the important quantity in cosmological analyses (e.g., Bode et al."
+ 2006)., 2006).
+" On the other Πας, lensing effects are generated through gravitation. and depend on (he total mass distribution."," On the other hand, lensing effects are generated through gravitation, and depend on the total mass distribution."
+ Thus it is expected (that a cluster sample detected through weak lensing effects is better suited [or cosmological studies in comparison with those selected bv other probes (e.g.. White et al.," Thus it is expected that a cluster sample detected through weak lensing effects is better suited for cosmological studies in comparison with those selected by other probes (e.g., White et al."
+ 2002; Hamana et al., 2002; Hamana et al.
+ 2004: Tang Fan 2005; Fang llaiman 2006)., 2004; Tang Fan 2005; Fang Haiman 2006).
+ Without involving complicated gas physics. weak lensing cluster detections. however. have their own shortcomings.," Without involving complicated gas physics, weak lensing cluster detections, however, have their own shortcomings."
+ Besides observational errors. physical svstematics. such as projection effects ancl complex mass distributions of clusters of galaxies. affect. the selection function of weak lensing clusters considerably.," Besides observational errors, physical systematics, such as projection effects and complex mass distributions of clusters of galaxies, affect the selection function of weak lensing clusters considerably."
+ Thus weak lensing cluster samples are not (uly mass-selected. (Tang Fan 2005)., Thus weak lensing cluster samples are not truly mass-selected (Tang Fan 2005).
+ The intrinsic ellipticities of background ealaxies result false peaks in lensing maps ancl reduce the eflicieney of cluster detections significantlv (e.g.. White et al.," The intrinsic ellipticities of background galaxies result false peaks in lensing maps and reduce the efficiency of cluster detections significantly (e.g., White et al."
+ 2002)., 2002).
+ The false peaks could also be misinterpreted as clark clumps. which nueht lead (o a laulty conclusion regarding Che validity of a cosmological model.," The false peaks could also be misinterpreted as dark clumps, which might lead to a faulty conclusion regarding the validity of a cosmological model."
+ Ilere we explore how the existence of the intrinsic alignment of background. galaxies allects the number of false peaks in lensing convergence maps., Here we explore how the existence of the intrinsic alignment of background galaxies affects the number of false peaks in lensing convergence maps.
+ We further propose that the nunmber of false peaks can be used to constrain sensitively the level of the intrinsic alignment if one can separate true and false peaks observationally., We further propose that the number of false peaks can be used to constrain sensitively the level of the intrinsic alignment if one can separate true and false peaks observationally.
+ The rest of the paper is organized as Lollows., The rest of the paper is organized as follows.
+ In 822. we discuss the correlations of galaxy," In 2, we discuss the correlations of galaxy"
+"is strongly correlated with (rai,fe, fy,); stronger shocks are expected in regions very close to the central engine particularly around a rotating black hole.","is strongly correlated with $r_{\rm sh}, f_E,
+f_{\dot{M}}$ ); stronger shocks are expected in regions very close to the central engine particularly around a rotating black hole."
+ Strong shocks in principle are accompanied by high energy and mass loss fractions., Strong shocks in principle are accompanied by high energy and mass loss fractions.
+" To sum up, our results show that strongest stable shocks generally develop at the smallest radii (closer to the black hole), accompanied by the largest mass loss f,, (and the largest energy loss fg as well)."," To sum up, our results show that strongest stable shocks generally develop at the smallest radii (closer to the black hole), accompanied by the largest mass loss $f_{\dot{M}}$ (and the largest energy loss $f_E$ as well)."
+" The rotation of the black hole apparently amplifies the shock strength by more than a factor of two, also extending the outflowing site significantly inward (i.e., ra,/m2—3 when a/m=0.99 while ra,/m212 when a/m= 0)."," The rotation of the black hole apparently amplifies the shock strength by more than a factor of two, also extending the outflowing site significantly inward (i.e., $r_{\rm sh}/m \gtrsim 2-3$ when $a/m=0.99$ while $r_{\rm sh}/m
+\gtrsim 12$ when $a/m=0$ )."
+" We will make some implications of the obtained shock-outflow solutions in the last section, 84."," We will make some implications of the obtained shock-outflow solutions in the last section, 4."
+ Samples of the global run of the physical parameters of accretion flows that include shocks are given in Figures 9 (a/m= 0) and 10 (a/m= 0.99).," Samples of the global run of the physical parameters of accretion flows that include shocks are given in Figures \ref{fig:global-1}
+ $a/m=0$ ) and \ref{fig:global-2} $a/m=0.99$ )."
+ In Tables 3 and 4 we, In Tables \ref{tab:tbl-3} and \ref{tab:tbl-4} we
+seen in Panel b of Figure 4..,seen in Panel b of Figure \ref{fig:flux}.
+ Therelore. we suspect that the relative increase in the flux in (he observations compared with those of the other observatories may. be due to instrumental calibration issues.," Therefore, we suspect that the relative increase in the flux in the observations compared with those of the other observatories may be due to instrumental calibration issues."
+ The fluxes measured from the last [our observations taken byChandra... and Suzakucan be fitted with a constant {hax model (reduced 4?=1.9 for 3 degrees of freedom. corresponding to a probability of having occurred. by chance of 0.125).," The fluxes measured from the last four observations taken by, and can be fitted with a constant flux model (reduced $\chi^2 = 1.9$ for 3 degrees of freedom, corresponding to a probability of having occurred by chance of 0.125)."
+ Thus we conclude that the phase-averaged 410 keV [Iuxes of the last four observations were consistent with being constant. anc we put an upper limit of on long-term variability in this energv band.," Thus we conclude that the phase-averaged 4–10 keV fluxes of the last four observations were consistent with being constant, and we put an upper limit of on long-term variability in this energy band."
+ I[nportantly. the (vo observations were taken only 88 and 90 days alter (he first eliteh. and the observation was taken only 27 davs alter the third gliteh.," Importantly, the two observations were taken only 88 and 90 days after the first glitch, and the observation was taken only 27 days after the third glitch."
+ By contrast. ihe 410 keV flux of was higher than in quiescence 21 davs alter its2002 eliteh (?).. that of was a [actor of G higher 238 days alter its 2007 elitches (?).. and was higher [or ~53 davs alter its first 2005 glitch as inferred [rom ?Tools/w3pinmns.," By contrast, the 4–10 keV flux of was higher than in quiescence 21 days after its2002 glitch \citep{zkd+08}, that of was a factor of 6 higher $\ge$ 38 days after its 2007 glitches \citep{tgd+08}, , and was higher for $\sim$ 53 days after its first 2005 glitch as inferred from \citet{gri+07}."
+htmnl). Therefore. we conclude that unlike in22592-586.. LE —5937 and possibly.. there is no evidence for gliteh-correlated flix changes in ANP..," Therefore, we conclude that unlike in, 1E $-$ 5937 and possibly, there is no evidence for glitch-correlated flux changes in AXP."
+ One caveat of our study is that we were limited to the harder part of the neutron stars emission spectrum., One caveat of our study is that we were limited to the harder part of the neutron star's emission spectrum.
+ The flux from the blackbody component was not well constrained., The flux from the blackbody component was not well constrained.
+" ""Therefore. we cannot rule out changes in the neutron stars thermal radiation. only changes in the power-law component in the 410 keV band. which we note constitutes ~0.25 of the stellar flux (BB+POW) in the 110 keV band."," Therefore, we cannot rule out changes in the neutron star's thermal radiation, only changes in the power-law component in the 4–10 keV band, which we note constitutes $\sim$ 0.25 of the stellar flux (BB+POW) in the 1–10 keV band."
+ Thus our results argue against the hvpothesis that there exists a generic correlation between the X-raw flux variability and elitch epochs in ANDs. further supporting the argument that glitehes in ANPs can be either radiativelv loud or radiativelv silent (2)..," Thus our results argue against the hypothesis that there exists a generic correlation between the X-ray flux variability and glitch epochs in AXPs, further supporting the argument that glitches in AXPs can be either radiatively loud or radiatively silent \citep{dkg08}."
+ There is of course precedent for radiativelv silent elitches in neutron stars. in that no rotation-powered pulsar eliteh has ever been reported to be accompanied with anv radiative change. al(hough rapid X-ray follow-up has only been accomplished in one case (?)..," There is of course precedent for radiatively silent glitches in neutron stars, in that no rotation-powered pulsar glitch has ever been reported to be accompanied with any radiative change, although rapid X-ray follow-up has only been accomplished in one case \citep{hgh01}."
+ Any. physical model of maegnetar elitehes will have to explain the simultaneous existence of both tvpes., Any physical model of magnetar glitches will have to explain the simultaneous existence of both types.
+ This is true of even a single source. as there is evidence that ANP 22594-586has both. given that its most recent eliteh showed no pulsedflux change (2)..," This is true of even a single source, as there is evidence that AXP has both, given that its most recent glitch showed no pulsedflux change \citep{dkg08b}. ."
+ Recently ? have argued thal AXP, Recently \citet{es10} have argued that AXP
+time.,time.
+ For linear variation. the present result yields. [i/i5.6«10i* per year. more than an order of magnitude better than the best (model-dependent) laboratory constraint on changes in j/ (Rosenbandetal.2008).," For linear variation, the present result yields $\dot{\mu}/\mu < 5.6 \times 10^{-17}$ per year, more than an order of magnitude better than the best (model-dependent) laboratory constraint on changes in $\mu$ \citep{rosenband08alt}."
+". In conclusion. we have used the GBT to obtain high S/N spectra in the NH; (1.1). CS 1-0 and H2CO θυ- Io, transitions from πο τι0.685 absorber. towards B02184-357."," In conclusion, we have used the GBT to obtain high S/N spectra in the $_3$ (1,1), CS 1-0 and $_2$ CO $_{\rm 00}$ $_{\rm 01}$ transitions from the $z \sim 0.685$ absorber towards B0218+357."
+" A comparison between the redshifts of the inversion and rotational lines yields a velocity offset of AV.=|0,56+4O.1L0Cstat.) km/s. with the NH, lines blueshifted relative toO.068(syst.)| the rotational ones."," A comparison between the redshifts of the inversion and rotational lines yields a velocity offset of $\Delta V = [-0.36 \pm 0.10 (stat.) 
+\pm 0.068 (syst.)]$ km/s, with the $_3$ lines blueshifted relative to the rotational ones."
+ Two sources of systematic error remain to be quantified. arising from (1) velocity offsets between nitrogen-bearing and carbon-bearing species. and (2) time variability in the background source morphology. which might yield different sightlines at different epochs.," Two sources of systematic error remain to be quantified, arising from (1) velocity offsets between nitrogen-bearing and carbon-bearing species, and (2) time variability in the background source morphology, which might yield different sightlines at different epochs."
+ We find no statistically-significant (=3) evidence for changes in jr. with the 30 constraint |Api/ji]<<3.6<10* over a lookback time of 6.2 Gyrs.," We find no statistically-significant $\ge 3\sigma$ ) evidence for changes in $\mu$, with the $3\sigma$ constraint $\dmu < 3.6 \times 
+10^{-7}$ over a lookback time of 6.2 Gyrs."
+ This. ας the strongest present constraint on temporal changes in the proton-electron mass ratio., This is the strongest present constraint on temporal changes in the proton-electron mass ratio.
+ I thank Bob Carswell and Michael Murphy for much help with using for the Voigt-profile fitting. Carl Bignell and Bob Garwood for help in the scheduling of the GBT observations and the GBT data analysis. and Dave Meier for permission to use the H»CO line detected in our survey of BO218+357 for this project.," I thank Bob Carswell and Michael Murphy for much help with using for the Voigt-profile fitting, Carl Bignell and Bob Garwood for help in the scheduling of the GBT observations and the GBT data analysis, and Dave Meier for permission to use the $_2$ CO line detected in our survey of B0218+357 for this project."
+ I also acknowledge support from the Department of Science and Technology. India. via a Ramanujan Fellowship.," I also acknowledge support from the Department of Science and Technology, India, via a Ramanujan Fellowship."
+ The National Radio Astronomy Observatory is operated by Associated Universities. Inc. cπάει cooperative agreement with the NSF.," The National Radio Astronomy Observatory is operated by Associated Universities, Inc, under cooperative agreement with the NSF."
+DeCOMCS The functiou defined by equation (12)) is reduced to &=gf|29/76.geOl? where f=oufay is the specific angular moment as measured at infinity.,"becomes The function defined by equation \ref{eq:kk}) ) is reduced to $k = -g^{tt} +2g^{t\phi}\ell -g^{\phi\phi}\ell^2 $ , where $\ell\equiv -u_\phi/u_t$ is the specific angular momentum as measured at infinity."
+" The wuuerator (20)) and denominator (21)) become where the function &=S(r.0) is related to the configuration of a stream line: for example. for a radial stream line. Slr)=(2,2Wrooa3)/CrN)."," The numerator \ref{eq:numer}) ) and denominator \ref{eq:domom}) ) become where the function ${\cal S}={\cal S}(r,\theta)$ is related to the configuration of a stream line; for example, for a radial stream line, $ {\cal S}(r) = -(2r^2-3mr+a^2)/(r\Delta) $."
+ Tere. we sliold remember that. in the ideal MIID case. a stream line coincides with a magnetic field line.," Here, we should remember that, in the ideal MHD case, a stream line coincides with a magnetic field line."
+" So. even forDa weak magnetic feld. the factor S should be related to the magnetic field lines. aud im fact it must be 10]laced by S=(luB,,)’ for MIID flows."," So, even for a weak magnetic field, the factor ${\cal S}$ should be related to the magnetic field lines, and in fact it must be replaced by ${\cal S}=(\ln B_p)'$ for MHD flows."
+ Further. we should sav that. iu the weak-aenetic field liit. Org represeuts the augular frequency of the stream ine for a bydrodvuamical flow. aud it would be determined as the augular velocity of the injection point.," Further, we should say that, in the weak-magnetic field limit, $\Omega_F$ represents the angular frequency of the stream line for a hydrodynamical flow, and it would be determined as the angular velocity of the injection point."
+" ""Thus. the above expressions are reduced to the relativistic ls""drodvuanuical flow equations formulated by (1986)."," Thus, the above expressions are reduced to the relativistic hydrodynamical flow equations formulated by \citet{Lu86}."
+. Iu equations (21)). (25)) aud (26)). we can also check that the fast 1iagnetosonic wave speed equals he sound wave speed. while the Alfvéuu wave speed aud the slow magnuetosonic wave speed become zero.," In equations \ref{eq:aw}) ), \ref{eq:fw}) ) and \ref{eq:sw}) ), we can also check that the fast magnetosonic wave speed equals the sound wave speed, while the Alfvénn wave speed and the slow magnetosonic wave speed become zero."
+ ToJj the coutrary. forJ 7Pπα...Q(z). which correspouds to case (1). the hydrodynamical expression for he poloidal equation can not be obtained by a flow solution.," To the contrary, for $u_p^2 \sim u_{\rm AW}^2 \sim O(\varepsilon)$, which corresponds to case (i), the hydrodynamical expression for the poloidal equation can not be obtained by a flow solution."
+" At the Alfvén poiut. A3=ayO(z"") aud ix~Oe)."," At the Alfvénn point, $M^2_{\rm A}=\alpha_{\rm A}\sim O(\varepsilon^{0})$ and $u_{\rm A}^2 \sim O(\varepsilon)$."
+ At the fast maguctosonic point. ug~Ole).," At the fast magnetosonic point, $u_{\rm F}^2 \sim O(\varepsilon)$."
+" So. the velocity of traus-Alfvénnic ΑΠΟ accretion would also be the order of 212,"," So, the velocity of trans-Alfvénnic MHD accretion would also be the order of $\varepsilon^{1/2}$."
+ We have considered stationary aud axisvuunetric hot ideal MITD accretion alone a flux-tube connected frou a plasina source to the event horizon., We have considered stationary and axisymmetric hot ideal MHD accretion along a flux-tube connected from a plasma source to the event horizon.
+ To argue the details of the boundary condit1ος at the plasma source would take us bevoud the scope of this paper., To argue the details of the boundary conditions at the plasma source would take us beyond the scope of this paper.
+ Therefore. we have surveved the dependence of the trans-fast ΑΠΟ flows on a wide range of source parameters.," Therefore, we have surveyed the dependence of the trans-fast MHD flows on a wide range of source parameters."
+ We lave shown that. when t16 forbidden region is type TA or TTA. there exist two physically different accretion regimes: (à) MIID accretion aud Q1) MIID accretion.," We have shown that, when the forbidden region is type IA or IIIA, there exist two physically different accretion regimes: (i) MHD accretion and (ii) MHD accretion."
+ The magueto-lkeMIID accreion would be injected from jear the separation point aud passes through the inner Alfvéun point with a smaller 4., The magneto-likeMHD accretion would be injected from near the separation point and passes through the inner Alfvénn point with a smaller $\hat\eta$.
+ Ou the other haud. tje lvdro-like MIID accretion with a sufficiently large 5 would be injected. frou )etween the outer Alfvónu point and the outer ieht surface and passes throueh the outer Alfvéóun point.," On the other hand, the hydro-like MHD accretion with a sufficiently large $\hat\eta$ would be injected from between the outer Alfvénn point and the outer light surface and passes through the outer Alfvénn point."
+ Ibvdro-like accretion nav. also be initially 1iaguetosonic or super-Alfvéuuic., Hydro-like accretion may also be initially super-slow magnetosonic or super-Alfvénnic.
+ A hot ideal MIID plasna with larger 1) cannot acerete stationary onto he black hole after passing through the inner Alfvéun point., A hot ideal MHD plasma with larger $\hat\eta$ cannot accrete stationary onto the black hole after passing through the inner Alfvénn point.
+ Then. if the value of jj increases with a secular nnuescale. the maegneto-like MITD accretion solution should trausit to the lydro-like MIID accretion solution: he inverse process would be also possible.," Then, if the value of $\hat\eta$ increases with a secular timescale, the magneto-like MHD accretion solution should transit to the hydro-like MHD accretion solution; the inverse process would be also possible."
+ The criterion for cistineuishine between the two reeimes is based on the locatious of both the Alfvóun point and the fast maeuetosoue point. which change discoutiniously during the transition.," The criterion for distinguishing between the two regimes is based on the locations of both the Alfvénn point and the fast magnetosonic point, which change discontinuously during the transition."
+ For the magnueto-lke MIID accretion. we have found that the location of the N-type fast 1niagnetosonic point is not unique for fixed intrinsicparameters of the accreting plasma.," For the magneto-like MHD accretion, we have found that the location of the X-type fast magnetosonic point is not unique for fixed intrinsicparameters of the accreting plasma."
+For example. iu Figures baa. το and 6.. iu the range of rpxrgX r.a gj coustaut line crosses a solid curve of Gex:0.1 at three poiuts:,"For example, in Figures \ref{fig:eta-OM}a a, \ref{fig:eta-spin}b b and \ref{fig:eta-alf}, , in the range of $r_{\rm H}<r_{\rm F}<r_{\rm A}^{in}$ , a $\hat\eta=$ constant line crosses a solid curve of $\zeta_{\rm F}\leq 0.1$ at three points;"
+We thank Ix.V. Johnston lor providing the models used in (his paper and for useful discussion.,We thank K.V. Johnston for providing the models used in this paper and for useful discussion.
+ We also thank G. van de Venn. David Valls-Gabaud. Andrew Cooper and M. A. Gomez-Flechoso for useful discussions.," We also thank G. van de Venn, David Valls-Gabaud, Andrew Cooper and M. A. Gomez-Flechoso for useful discussions."
+ D. M-D acknowledges funding from the Spanish Ministry. of Education (Ramon ν Cajal fellowship aud research project AYA 2007-65090) and the Instituto de Astrofisica de Canarias (provect 3I. SRM appreciates support from NSF grant. AST-0807945., D. M-D acknowledges funding from the Spanish Ministry of Education (Ramon y Cajal fellowship and research project AYA 2007-65090) and the Instituto de Astrofisica de Canarias (proyect 3I. SRM appreciates support from NSF grant AST-0807945.
+Even if the interpretation of the broad component of the iron line as a blend of two ionized lines successfully fits the data. we decided to test the alternative hypothesis that this component is actually the result of reprocessing by an ionized disc.,"Even if the interpretation of the broad component of the iron line as a blend of two ionized lines successfully fits the data, we decided to test the alternative hypothesis that this component is actually the result of reprocessing by an ionized disc."
+ This interpretation was tried and rejected by ?.. on the basis of a model including andDISKLINE.," This interpretation was tried and rejected by \cite{Matt01}, on the basis of a model including and."
+. We therefore decided to adopt a more refined model. (?).. to fit the simultaneous pn/PDS spectrum.," We therefore decided to adopt a more refined model, \citep{nkk_xion}, to fit the simultaneous pn/PDS spectrum."
+ This model consists of reflection from a photoionized disc and the associated iron line. including relativistic effects.," This model consists of reflection from a photoionized disc and the associated iron line, including relativistic effects."
+" The adopted geometry ts that of a lamp-post. with free parameters of the model being the height of the X-ray source above the disk. hy. the accretion rate through the disc. m. the ratio between the luminosity of the source and that of the disc. Εις / Ly. the inner and outer disk radii (the first kept fixed to 3 R, in our fits) and the spectral index."," The adopted geometry is that of a lamp-post, with free parameters of the model being the height of the X-ray source above the disk, $_\mathrm{x}$, the accretion rate through the disc, $\mathrm{\dot{m}}$, the ratio between the luminosity of the X-ray source and that of the disc, $_\mathrm{x}$ / $_\mathrm{d}$, the inner and outer disk radii (the first kept fixed to 3 $_\mathrm{s}$ in our fits) and the spectral index."
+We further added a narrow line at 6.4 keV (which cannot be produced in an ionized disc) to this model.,We further added a narrow line at 6.4 keV (which cannot be produced in an ionized disc) to this model.
+ The fit is much worse than the one with reflection from neutral matter 31 with 2 less dof). mostly due to a difficulty in reproducing the neutral iron edge at 7.1 keV. Therefore. it seems likely that most (if not all) the observed Compton reflection should be produced by a Compton-thick. cold material. where the neutral iron line ts also likely produced.," The fit is much worse than the one with reflection from neutral matter $\Delta\chi^2=31$ with 2 less dof), mostly due to a difficulty in reproducing the neutral iron edge at 7.1 keV. Therefore, it seems likely that most (if not all) the observed Compton reflection should be produced by a Compton-thick, cold material, where the neutral iron line is also likely produced."
+ Thus. the following step was to add a further component. a Compton reflection from neutral matter modelled as with R fixed to a negative value (to include only the reflected spectrum) and the spectral index linked to that of.," Thus, the following step was to add a further component, a Compton reflection from neutral matter modelled as with R fixed to a negative value (to include only the reflected spectrum) and the spectral index linked to that of."
+ The result is now statistically as good as the one obtained with only the Compton reflection and the three lines (see Table 3))., The result is now statistically as good as the one obtained with only the Compton reflection and the three lines (see Table \ref{xionpexrav}) ).
+ Fig., Fig.
+ 7 shows the components of this model: most of the reflection at high energy. the iron edge and the narrow line are produced by neutral matter. while the broad component of the line originates from the tonized disc. whose reflected continuum does not contribute too much to the high energy spectrum.," \ref{xionmodel} shows the components of this model: most of the reflection at high energy, the iron edge and the narrow line are produced by neutral matter, while the broad component of the line originates from the ionized disc, whose reflected continuum does not contribute too much to the high energy spectrum."
+ However. the unbounded values of the distance of the illuminating lamp-post from the dise and the outer radius of the dise itself in the fit (see Table 3)) are clearly indicating that there are not relativistic signatures on the spectrum.," However, the unbounded values of the distance of the illuminating lamp-post from the disc and the outer radius of the disc itself in the fit (see Table \ref{xionpexrav}) ) are clearly indicating that there are not relativistic signatures on the spectrum."
+ Moreover. the value of Ly / Ly is very high. requiring an extreme fine-tuning for a magnetic flares model.," Moreover, the value of $_\mathrm{x}$ / $_\mathrm{d}$ is very high, requiring an extreme fine-tuning for a magnetic flares model."
+ These results are very similar to those obtained by ?— with a model includingPEXRAV.. andDISKLINE.. which is just a less sophisticated physical picture of the one we adopted: the line is not very broad and requires a large outer radius for the dise (lack of relativistic signatures) and Compton reflection from the ionized dise. is negligible if compared to that from neutral matter.," These results are very similar to those obtained by \citet{Matt01} with a model including, and, which is just a less sophisticated physical picture of the one we adopted: the line is not very broad and requires a large outer radius for the disc (lack of relativistic signatures) and Compton reflection from the ionized disc is negligible if compared to that from neutral matter."
+ These considerations make the ionized disc origin of the broad component less tenable than the identification as a blend of ionized lines. despite they are statistically equivalent.," These considerations make the ionized disc origin of the broad component less tenable than the identification as a blend of ionized lines, despite they are statistically equivalent."
+ The Chandra image of NGC 5506 clearly shows that the emission below | keV is extended on a radius of =350 pe (Fig. 8))., The $Chandra$ image of NGC 5506 clearly shows that the emission below 1 keV is extended on a radius of $\simeq350$ pc (Fig. \ref{soft}) ).
+ Furthermore. this region is clearly asymmetric with respect to the nuclear emission.," Furthermore, this region is clearly asymmetric with respect to the nuclear emission."
+ Interestingly the dimension and the asymmetry of the X-ray emission seems to agree fairly well with the radio maps presented by ?.., Interestingly the dimension and the asymmetry of the X-ray emission seems to agree fairly well with the radio maps presented by \citet{schm01}.
+ On the other hand. no clear evidence of a possible extension on larger scale. as claimed by ? on ROSAT HRI data. was found.," On the other hand, no clear evidence of a possible extension on larger scale, as claimed by \citet{colbert98} on $ROSAT$ HRI data, was found."
+ This extended gas ts naturally associated to the soft excess best observed in the two XMM spectra and the ASCA one. allowing us to exclude thermal bremsstrahlung or a blackbody as possible explanations. since it is hard to immagine such a spectral shape arising from a 350 pe scale gas.," This extended gas is naturally associated to the soft excess best observed in the two XMM spectra and the $ASCA$ one, allowing us to exclude thermal bremsstrahlung or a blackbody as possible explanations, since it is hard to imagine such a spectral shape arising from a 350 pc scale gas."
+ Therefore. in order to fit the residuals on the low energy part of the spectrum. we are left with two physically consistent models:," Therefore, in order to fit the residuals on the low energy part of the spectrum, we are left with two physically consistent models:"
+" We are interested in radiation-dominated epoch where µ distortions are generated, and we have H(z)=Hal!mel+z) and ne(z)=(tgo+2nHe0)(1ο)ποzy.addition we have kp=Αρ(1+z)°/? and d(1/kp?)/dz=-3Ap(1+ϱ)1. and the above equation simplifies to where The redshift and n; dependence obtained above matches that of ?.."," We are interested in radiation-dominated epoch where $\mu$ distortions are generated, and we have $H(z)=H_0\Omega_r^{1/2}(1+z)^2$ and $\Ne (z)=(n_{H0}+2n_{He0})(1+z)^3\equiv n_{e0}(1+z)^3$ we have $\kD =A_D^{-1/2}(1+z)^{3/2}$ and $\id(1/\kD ^2)/\id z=-3A_D(1+z)^{-4}$, and the above equation simplifies to where The redshift and $n_s$ dependence obtained above matches that of \citet{hss94}. ."
+ The diffusion scale at z=5x104 and kp=10*Mpc™! at z=2x10°., The diffusion scale $\kD =46\rm{Mpc}^{-1}$ at $z=5\times 10^4$ and $\kD =10^4\rm{Mpc}^{-1}$ at $z=2\times 10^6$.
+ Thus we are probing the primordial fluctuations on very small scales that are not accessible in any other way., Thus we are probing the primordial fluctuations on very small scales that are not accessible in any other way.
+" Thus for Harrison-Zeldovich spectrum with n,=1 We note the surprising fact that this redshift dependence is exactly the same as for the energy losses due to adiabatic cooling Eq.(I3),,?."," Thus for Harrison-Zeldovich spectrum with $\nS $ =1 We note the surprising fact that this redshift dependence is exactly the same as for the energy losses due to adiabatic cooling Eq.,."
+" For the currently value of n,=0.96 we have Figure[7] compares the rate of energy release multiplied by the blackbody visibility function, G(z)(1+z)dQ/dz/Ey, for different initial power spectra with the energy losses due to the adiabatic cooling of baryons, (1--z)$compion/Ey obtained by solving Eq."," For the currently value of $\nS =0.96$ we have Figure \ref{energyfig} compares the rate of energy release multiplied by the blackbody visibility function, $G(z)(1+z)\id Q/\id z/E_{\gamma}$ for different initial power spectra with the energy losses due to the adiabatic cooling of baryons, $(1+z){S_{\rm Compton}}/E_{\gamma}$ obtained by solving Eq."
+ for the standard recombination history (?) calculated using the effective multilevel approach and taking recent corrections into account (??)..," for the standard recombination history \citep{sss2000} calculated using the effective multilevel approach and taking recent corrections into account \citep{ah2011,ct2011}."
+ The spike due to increase in Silk damping during recombination is clearly abld/].The smaller spikes due toHell and Hel are also noticeable., The spike due to increase in Silk damping during recombination is clearly .The smaller spikes due toHeII and HeI are also noticeable.
+ Recombination of each species, Recombination of each species
+be smaller than this.,be smaller than this.
+ All observed Class I sources except Elias 29 are relatively low luminosity. 2.1 L.< Lig€13 L. (Wilking.Lada.&Young1989.hereafterWLY)..," All observed Class I sources except Elias 29 are relatively low luminosity, 2.4 $_{\odot} \leq$ $_{\rm bol} \leq 13$ $_{\odot}$ \citep[hereafter WLY]{WLY89}."
+" All observed flat-spect""Ld YSOs have 1 L.€ Ly«3 L. (WLY: Greeneetal. 1991)).", All observed flat-spectrum YSOs have 1 $_{\odot} \leq$ $_{\rm bol} < 3$ $_{\odot}$ (WLY; \citeauthor{GWAYL94} \citeyear{GWAYL94}) ).
+ These huniuositles are consistent with the sources beiug low-mass PMS YSOs (M « 1 M.) if the Class E objects are powere by acceretion. an assumption which is consistent with their high veiliugs (see also 81.2).," These luminosities are consistent with the sources being low-mass PMS YSOs (M $<$ 1 $_{\odot}$ ) if the Class I objects are powered by accretion, an assumption which is consistent with their high veilings (see also 4.2)."
+" In using VSSC lY as a spectral type template to measure veilines. we iive implicitly assumeL that all sour""es have spectral types near MOIV/V. a typical value for T Tauri stars."," In using VSSG 17 as a spectral type template to measure veilings, we have implicitly assumed that all sources have spectral types near M0IV/V, a typical value for T Tauri stars."
+ A young star of this spectral type has a mass of approximately. 0.1 M. if on the birhliueof the H-R Diagram 1988)., A young star of this spectral type has a mass of approximately 0.4 $_{\odot}$ if on the birthlineof the H–R Diagram \citep[see][]{S88}.
+. Loan observed YSO is really 1 M... then it wouk have a spectral type «X Ix3-HV/V (seeO88) and thus an intrinsic stellar CO eqivalent wicltrol ouly about 0.6 tinies tha ofan MOIV/V star (see Paper I.," If an observed YSO is really 1 $_{\odot}$, then it would have a spectral type of K3–4IV/V \citep[see][]{DM97, S88} and thus an intrinsic stellar CO equivalent width of only about 0.6 times that of an M0IV/V star (see Paper I)."
+" This would mean that 1s actual veiling is rn=6r,—0.1 (see 1.2 of Paper I) where rj is the veiling estimate iade assu1ing an MOIV/V spectral type."," This would mean that its actual veiling is $r'_{k} = 0.6 
+r_{k} - 0.4$ (see 4.2 of Paper I) where $r_{k}$ is the veiling estimate made assuming an M0IV/V spectral type."
+ Likewise. if an observed object were really a lower mass star near je brown cbwarf limit. then its true birthline spectral ty‘pe would ye near M6IV/V and its intrinsi CO absorption would be about ereater than that of the MOIV/V template.," Likewise, if an observed object were really a lower mass star near the brown dwarf limit, then its true birthline spectral type would be near M6IV/V and its intrinsic CO absorption would be about greater than that of the M0IV/V template."
+ Thus the true veiines of observed sources may cdifler from our derived oues by as 1wich as50%., Thus the true veilings of observed sources may differ from our derived ones by as much as.
+. However. it is urlikely that many oL the observed objects have such large cleviations (in either direction) because eir luminosities are consistent. with most having masses of approximately. 0.5 M... similar to that expected [οι our MOIV/V tempale.," However, it is unlikely that many of the observed objects have such large deviations (in either direction) because their luminosities are consistent with most having masses of approximately 0.5 $_{\odot}$, similar to that expected for our M0IV/V template."
+ We lave usec other published information on the observed objects when avallable to estimae their intrinsic syectval types before cle‘iving velliug estimates (see 1.1 and 1.2)., We have used other published information on the observed objects when available to estimate their intrinsic spectral types before deriving veiling estimates (see 4.1 and 4.2).
+ Finally. we stidied 1ow LOlatiol decreases the apparelt CO band. head. absorption depths.," Finally, we studied how rotation decreases the apparent CO band head absorption depths."
+" is elect can be separaed [rom veiliug by measuring the slope of the CO baud bead. but this is not possible for sources which clo 1ot show this absorptio1 featuο,"," This effect can be separated from veiling by measuring the slope of the CO band head, but this is not possible for sources which do not show this absorption feature."
+ The majority of sotrces in our sample do not show tus feature. so we do not know their rotation velocities.," The majority of sources in our sample do not show this feature, so we do not know their rotation velocities."
+ Our earlier study (Paper IL) showed that fIat-spectrum x)urces rotate significantly [aser thau Class Il YSOs. sO OUL rent sample of more highly. embedecd flat-spectrum aux Class I YSOs may 'Otate even Laster still (perhaps ve sin 7>50 km H.," Our earlier study (Paper II) showed that flat-spectrum sources rotate significantly faster than Class II YSOs, so our current sample of more highly embedded flat-spectrum and Class I YSOs may rotate even faster still (perhaps $v$ sin $i > 50$ km $^{-1}$ )."
+Eileed. receut N-ray observaious with tle ASCA satellite iudicate that one Class I source in our sample (IRS 13) is roating at least this rapicly 2000).," Indeed, recent X-ray observations with the ASCA satellite indicate that one Class I source in our sample (IRS 43) is rotating at least this rapidly \citep{MGTK00}."
+. We estimate that the eqtorial rotational breakup velocities of hese young stars (A~ολ. aud Rz 34.) are approNimately e2150 km 1. or e sin ioc150 km s! fora mean inclination 7=57deg.," We estimate that the equatorial rotational breakup velocities of these young stars $M 
+\simeq 0.5 M_{\sun}$ and $R \simeq 3 R_{\sun}$ ) are approximately $v 
+\simeq 180$ km $^{-1}$, or $v$ sin $i \simeq 150$ km $^{-1}$ for a mean inclination $i = 57\deg$."
+ Next. we sucdied how rotation decreases the maximiin CO band head absorption in YSOs by artificially rotating our WL 5 temdlate by convolving its observed spectrum with Iunb-darkened stellar broadeniug yroliles for 25 kia 1pο πα <175 kim I.," Next, we studied how rotation decreases the maximum CO band head absorption in YSOs by artificially rotating our WL 5 template by convolving its observed spectrum with limb-darkened stellar broadening profiles for 25 km $^{-1} \leq v$ sin $i \leq 175$ km $^{-1}$."
+ See Paper LI for more details :uxd examples of artificially rotated specTa., See Paper II for more details and examples of artificially rotated spectra.
+ These experinents showed that the maximum CO asorption depth of the VSSC 25 template with e sin =150 kin s Lds a factor of 1.37 weaker han the one with ο sin /=50 kms ?., These experiments showed that the maximum CO absorption depth of the VSSG 25 template with $v$ sin $i = 150$ km $^{-1}$ is a factor of 1.37 weaker than the one with $v$ sin $i = 50$ km $^{-1}$ .
+ Therefore if VSSC 17 were rotating at breakup. its maniuum CO absorption depth would be ouly as deep as low seen in its spectrum.," Therefore if VSSG 17 were rotating at breakup, its maximum CO absorption depth would be only as deep as now seen in its spectrum."
+ Thus if rotating near breakup. the featureless objects," Thus if rotating near breakup, the featureless objects"
+a region where the gas transitions from optically thin to optically thick where the temperatures differ.,a region where the gas transitions from optically thin to optically thick where the temperatures differ.
+ Figure 3. shows density (left) and gas temperature (right) plotted against radius as the collapse passes various values of the maximum density., Figure \ref{fig:BMmix} shows density (left) and gas temperature (right) plotted against radius as the collapse passes various values of the maximum density.
+ The first core ean be seen at a radius of ~It em where the density and temperature gradients change abruptly., The first core can be seen at a radius of $\sim 10^{14}$ cm where the density and temperature gradients change abruptly.
+" Similarly. the transition to the stellar core can be seen at a radius of ~10"" em."," Similarly, the transition to the stellar core can be seen at a radius of $\sim 10^{12}$ cm."
+ ?. give the density and temperature profiles when the central temperature reaches =1000 K. We reproduce their profiles as the thick solid black lines in Figure 3.., \citet{BM1992} give the density and temperature profiles when the central temperature reaches $\approx 1000$ K. We reproduce their profiles as the thick solid black lines in Figure \ref{fig:BMmix}.
+ Our results are in good agreement with theirs. the main ditference being that our results give slightly higher temperatures at radii ~10 em. presumably due to the fact that we use flux-limited ditfusion while ? use the Eddington approximation (the former retards radiation transport near the surface of the first core).," Our results are in good agreement with theirs, the main difference being that our results give slightly higher temperatures at radii $\sim 10^{15}$ cm, presumably due to the fact that we use flux-limited diffusion while \citet{BM1992} use the Eddington approximation (the former retards radiation transport near the surface of the first core)."
+ In Figure 4.. we plot temperature versus density at a series of snapshots during the highest resolution caleulation (since both temperature and density are monotonically decreasing functions of radius. we can plot these as lines for each snapshot).," In Figure \ref{fig:BMsnap}, we plot temperature versus density at a series of snapshots during the highest resolution calculation (since both temperature and density are monotonically decreasing functions of radius, we can plot these as lines for each snapshot)."
+ The figure clearly shows that even if it were possible to construct a barotropic equation of state to follow the evolution of maximum density and temperature. it would generally underestimate the temperature of the gas away from the maximum.," The figure clearly shows that even if it were possible to construct a barotropic equation of state to follow the evolution of maximum density and temperature, it would generally underestimate the temperature of the gas away from the maximum."
+ In the figure. the barotropic equation underestimates the temperature bymore than an order of magnitude at densities ~10 ecm ? at late times.," In the figure, the barotropic equation underestimates the temperature bymore than an order of magnitude at densities $\sim 10^{-13}$ g $^{-3}$ at late times."
+ Figure 5 shows the evolution of the maximum value of temperature and density for the three types of initial conditions. each performed with 50.000 particles.," Figure \ref{fig:50k} shows the evolution of the maximum value of temperature and density for the three types of initial conditions, each performed with 50,000 particles."
+ The spherically-symmetrie ? collapse (green line) begins to heat at a much lower density than the ?. collapse (blue)., The spherically-symmetric \citet{BB1979} collapse (green line) begins to heat at a much lower density than the \citet{BM1992} collapse (blue).
+ Essentially the only difference between these calculations is the initial density of the two clouds. the former is initially almost two orders of magnitude denser. and yet the temperature evolutions are quite different.," Essentially the only difference between these calculations is the initial density of the two clouds, the former is initially almost two orders of magnitude denser, and yet the temperature evolutions are quite different."
+ Figure 6 shows that. as with the ? collapse. the spatial and temporal evolution of the temperature is also complex.," Figure \ref{fig:BBsnap} shows that, as with the \citet{BM1992} collapse, the spatial and temporal evolution of the temperature is also complex."
+ Again. these results demonstrate that a barotropic equation of state cannot accurately describe the temperature distribution during such cloud collapses.," Again, these results demonstrate that a barotropic equation of state cannot accurately describe the temperature distribution during such cloud collapses."
+ The result of the ? initial conditions. including the #7=2 density perturbation and rotation. is to form a wide binary with a separation of =1000 AU (2)..," The result of the \citet{BB1979} initial conditions, including the $m=2$ density perturbation and rotation, is to form a wide binary with a separation of $\approx 1000$ AU \citep{BBP1995}."
+ Using these initial conditions. the evolution of the maximum temperature versus density is much cooler than that from its spherically-symmetrie counterpart (see Figure 5.. red line).," Using these initial conditions, the evolution of the maximum temperature versus density is much cooler than that from its spherically-symmetric counterpart (see Figure \ref{fig:50k}, red line)."
+ In facet. coincidentally. it is actually quite close to the evolution obtained with the ? initial conditions.," In fact, coincidentally, it is actually quite close to the evolution obtained with the \citet{BM1992} initial conditions."
+ There are likely to be two reasons for the cooler temperatures at the same maximum density in the rotating collapse., There are likely to be two reasons for the cooler temperatures at the same maximum density in the rotating collapse.
+ First. the collapse to form the stars is delayed by the rotational support (they form at ¢=1.30 initial cloud. free-fall times instead. of just over one free-fall time for the spherically symmetric ease).," First, the collapse to form the stars is delayed by the rotational support (they form at $t=1.30$ initial cloud free-fall times instead of just over one free-fall time for the spherically symmetric case)."
+ The slower collapse allows more time for energy to be radiated away from the object., The slower collapse allows more time for energy to be radiated away from the object.
+ Second. the presence of rotation means that the distribution of material around each of the collapsed objects is no longer spherically symmetric.," Second, the presence of rotation means that the distribution of material around each of the collapsed objects is no longer spherically symmetric."
+ In particular. there are low density cavities along the rotation axes and high densities in the dises.," In particular, there are low density cavities along the rotation axes and high densities in the discs."
+ Radiation can more easily escape through these cavities than in the spherically symmetric case., Radiation can more easily escape through these cavities than in the spherically symmetric case.
+ We also tested the etfect of resolution on the ? binary star initial conditions. using 50.000 and 150.000 particles (Figure 7.. red and green lines respectively).," We also tested the effect of resolution on the \citet{BB1979} binary star initial conditions, using 50,000 and 150,000 particles (Figure \ref{fig:BBbinarycomp}, , red and green lines respectively)."
+ As with the ?. initial conditions. increasing the resolution above 50.000 particles has little ettect," As with the \citet{BM1992} initial conditions, increasing the resolution above 50,000 particles has little effect"
+PL velation shows a break around ten davs.,$PL$ relation shows a break around ten days.
+ These results are controversial and are still being debated in the literature (see for instance Benedict οἱ al., These results are controversial and are still being debated in the literature (see for instance Benedict et al.
+ 2008. and Pietrzvuski et al.," 2008, and Pietrzynski et al."
+ 2007)., 2007).
+ Given the impact (hat Cepheids have (traditionally had on the determination of distances {ο ealaxies il 1s important (o confirm the effect. quantilv it. 22d (hen i£ it proves to be real. fully assess (he impact of However. it is not the purpose of this short contribution to attempt (hat assessment. rather we wish to show explicitly (hat the reddening-Iree form of (he relation. the so-called Wesenheit hinction (V). can be demonstrated to be largely immune to changes in the shape of the instability strip Chat may then be reflected in changes in the slope of the PL relation.," Given the impact that Cepheids have traditionally had on the determination of distances to galaxies it is important to confirm the effect, quantify it, and then if it proves to be real, fully assess the impact of However, it is not the purpose of this short contribution to attempt that assessment, rather we wish to show explicitly that the reddening-free form of the Period-Luminosity relation, the so-called Wesenheit function $W$ ), can be demonstrated to be largely immune to changes in the shape of the instability strip that may then be reflected in changes in the slope of the $PL$ relation."
+ Neeow Ixaubur (2005) have already demonstrated in one instance (the Large Magellanic Cloud Cepheid sample) the robust nature of W in (liis reg., Ngeow Kanbur (2005) have already demonstrated in one instance (the Large Magellanic Cloud Cepheid sample) the robust nature of W in this regard.
+ In the followinge we show why this is more egenerally true., In the following we show why this is more generally true.
+ We beein by noting that. for all stars. Stephans Law in combination with simple spherical geometry gives rise {ο a (wo-parameter mapping of radius ancl effective temperature into total luminosity.," We begin by noting that, for all stars, Stephan's Law in combination with simple spherical geometry gives rise to a two-parameter mapping of radius and effective temperature into total luminosity."
+" By the definition of effective temperature (Tip) we have L—AsloT),(1) Allowing a photometric color ((V-I);. sav) to stand in for temperature. and a monochromatic magnitude (M. sav) to substitute for bolometric Iuminositv. only one more observable is needed to complete the (transfer from theory to observation."," By the definition of effective temperature $T_{eff}$ ) we have $ L = 4\pi R^2 \sigma T^4_{eff} ~~~~~~~~~(1)$ Allowing a photometric color $_o$, say) to stand in for temperature, and a monochromatic magnitude $_V$, say) to substitute for bolometric luminosity, only one more observable is needed to complete the transfer from theory to observation."
+ For Cepheids. their (easily observed) periods can stand in lor their (had (to observe) radii: where lines of constant radius in the color-magnitude plane fncüonally substitute for lines of constant period.," For Cepheids, their (easily observed) periods can stand in for their (hard to observe) radii; where lines of constant radius in the color-magnitude plane functionally substitute for lines of constant period."
+" The physically motivated equation above then becomes the following equation based exclusively on observables. while necessarily retaining the form of an equation with two (and only (wo) independent variables: in (his case log(P) and (V.—/),."," The physically motivated equation above then becomes the following equation based exclusively on observables, while necessarily retaining the form of an equation with two (and only two) independent variables: in this case log(P) and $(V-I)_o$."
+ Thus My= olog(P-(V=D)5(2) ]t is important to note that both of these above equations describe and cover the infinite plane of the color-magnitude diagram., Thus $ M_V = \alpha $ $ + \beta (V-I)_o + \gamma ~~~~~~~~~~~(2)$ It is important to note that both of these above equations describe and cover the infinite plane of the color-magnitude diagram.
+ There is no period-Iuminositw relation in either of these equations., There is no period-luminosity relation in either of these equations.
+ More forcefully stated. Equation (2) is the Period-Lunminosity relation for Cepheids.," More forcefully stated, Equation (2) is the Period-Luminosity relation for Cepheids."
+ Any general correlation of period and Inminosity will only comeconstraint of these equations., Any general correlation of period and luminosity will only come of these equations.
+ That constraint. in the context of Cepheids. is imposed by the," That constraint, in the context of Cepheids, is imposed by the"
+observed. wavelength of 0.55.,observed wavelength of $0.5\mu {\rm m}$.
+ Note. of course. that these points are not simultaneous with the NII points.," Note, of course, that these points are not simultaneous with the NIR points."
+ The fitted svnchrotron component generally replicates well the wavelength dependence of the polarisation. although for some sources. such as 333. the »oints are equally well given by a constant polarisation component.," The fitted synchrotron component generally replicates well the wavelength dependence of the polarisation, although for some sources, such as $-$ 333, the points are equally well given by a constant polarisation component."
+ This is what vou would expect [rom a pure synchrotron component. and these sources are typically BL Lac objects. from which vou would expect to see a svnchrotron-domunated SED.," This is what you would expect from a pure synchrotron component, and these sources are typically BL Lac objects, from which you would expect to see a synchrotron-dominated SED."
+ A notable exception to his is 441. whose polarisation is not [it well by a constant svnchrotron component. but is fit better by a combination of a synchrotron anc a significant power aw component.," A notable exception to this is $-$ 441, whose polarisation is not fit well by a constant synchrotron component, but is fit better by a combination of a synchrotron and a significant power law component."
+ However. it is apparent from these. plots hat having simultaneous optical polarisation measurements would better help discriminate between the combined model and a constant. polarisation moclel.," However, it is apparent from these plots that having simultaneous optical polarisation measurements would better help discriminate between the combined model and a constant polarisation model."
+ Synchrotron Εαν emitted from the PLIES quasars is most likely to come from a thin. relativistic jet. and the radiation will not be isotropic.," Synchrotron flux emitted from the PHFS quasars is most likely to come from a thin, relativistic jet, and the radiation will not be isotropic."
+ This means that there will be very little svnchrotron radiation directed towards the emission line clouds. which have a much larger covering angle.," This means that there will be very little synchrotron radiation directed towards the emission line clouds, which have a much larger covering angle."
+ Thus. the ionisation of these clouds will be due to the continuum emission from the central accretion disk region — in other words. the Bie Blue Bump.," Thus, the ionisation of these clouds will be due to the continuum emission from the central accretion disk region – in other words, the Big Blue Bump."
+ One might expect that adding a non-ionising svnchrotron component to the continuum will have the effect of reducing the equivalent. width of the emission lines from. the BLR. due to the fact that the Hux in the emission lines is not changed but the continuum Lux is increased.," One might expect that adding a non-ionising synchrotron component to the continuum will have the effect of reducing the equivalent width of the emission lines from the BLR, due to the fact that the flux in the emission lines is not changed but the continuum flux is increased."
+ To test this prediction. we compared. the equivalent widths of five emission lines 11549. 11909. 22198.LL3.. and the doublet 44959.5007) with the ratio of svnchrotron to continuum tux at the line wavelength. to see if some form of an anti-correlation is present.," To test this prediction, we compared the equivalent widths of five emission lines 1549, 1909, 2798, and the doublet 4959,5007) with the ratio of synchrotron to continuum flux at the line wavelength, to see if some form of an anti-correlation is present."
+ The details of the observations will »esented elsewhere (Francisetal.2001)., The details of the observations will presented elsewhere \cite{francis-spectra}.
+. Objects that had spectra taken were essentially a random sample of the PLI (subject to visibility during the observing run)., Objects that had spectra taken were essentially a random sample of the PHFS (subject to visibility during the observing run).
+ In Fig. 10..," In Fig. \ref{fig-ew},"
+ we show the results of this comparison for he aand lines. which are the two broad lines with the longest wavelength (and hence the two lines most likely to show a reduction in equivalent width).," we show the results of this comparison for the and lines, which are the two broad lines with the longest wavelength (and hence the two lines most likely to show a reduction in equivalent width)."
+ We have also plotted. those sources best fit with the power law model., We have also plotted those sources best fit with the power law model.
+ The value of the ratio used for these sources was taken from fitting the combined model. and so are upper limits to the ratio.," The value of the ratio used for these sources was taken from fitting the combined model, and so are upper limits to the ratio."
+ The Iline does not show much relationship to the svnchrotron ratio. while the line does show a reduction in equivalent width with increasing amount of svnchrotron.," The line does not show much relationship to the synchrotron ratio, while the line does show a reduction in equivalent width with increasing amount of synchrotron."
+ This lends some support o the hypothesis that excess svynchrotron light is present., This lends some support to the hypothesis that excess synchrotron light is present.
+ The dilference in the two plots is likely due to the presence of the turnover in the svnchrotron Hux. so that is has ess elfect. at the shorter wavelength. ofi.," The difference in the two plots is likely due to the presence of the turnover in the synchrotron flux, so that is has less effect at the shorter wavelength of."
+ There are. iowever. several sources that have high svnchrotron ratios ogether with high equivalent widths.," There are, however, several sources that have high synchrotron ratios together with high equivalent widths."
+ The implications of hese results are ciscussed in Section 10.., The implications of these results are discussed in Section \ref{sec-discuss}.
+ So fu. we have shown that. for a number of sources. the opticalNIU photometry is well Gt with a power law plus a curved. component. which we have assumed to he the turnover of a synchrotron component.," So far, we have shown that, for a number of sources, the optical–NIR photometry is well fit with a power law plus a curved component, which we have assumed to be the turnover of a synchrotron component."
+ However could another moclel be used instead of syachrotron?, However could another model be used instead of synchrotron?
+ One possible alternative is blackbody emission resulting from hot dust., One possible alternative is blackbody emission resulting from hot dust.
+ To test. this model we usec a blackhody emission spectrum due to dust ata temperature of 1750. Ix he sublimation temperature characteristic of dust. grains consisting of eraphite and silicates. Laor Draine (1993))) emitted in the quasars’ rest. frame.," To test this model, we used a blackbody emission spectrum due to dust at a temperature of 1750 K (the sublimation temperature characteristic of dust grains consisting of graphite and silicates, Laor Draine \shortcite{laor93}) ) emitted in the quasars' rest frame."
+ A blackbody curve at this temperature would have its peak. in the quasar rest-frame. at. 1.667.," A blackbody curve at this temperature would have its peak, in the quasar rest-frame, at $1.66 \mu {\rm m}$."
+ This blackbody spectrum was combined with the same Ato power Law used in the combined model to produce a model that was fit to the data., This blackbody spectrum was combined with the same $\lambda^{-1.7}$ power law used in the combined model to produce a model that was fit to the data.
+ The fits generated. by this model were almost always worse than those of the svnchrotron model., The fits generated by this model were almost always worse than those of the synchrotron model.
+ This was due to the peak of the blackbody occurring at much longer observe wavelengths than the A band., This was due to the peak of the blackbody occurring at much longer observed wavelengths than the $K$ band.
+ In only 12 cases was the X7/v value from the dust model less than that of the svnchrotron model. and in most cases this was because the cust mocde jac one more degree of freedom.," In only 12 cases was the $\chi^2/\nu$ value from the dust model less than that of the synchrotron model, and in most cases this was because the dust model had one more degree of freedom."
+ For cach of these objects. he SED took the form of a blue power law (one object was À14. and the rest were bluer than AE tt) that had a sligh: amount of reddening at the Jf and A bands. which was fit by he presence of the dust. blackbody curve.," For each of these objects, the SED took the form of a blue power law (one object was $\lambda^{-1.4}$, and the rest were bluer than $\lambda^{-1.54}$ ) that had a slight amount of reddening at the $H$ and $K$ bands, which was fit by the presence of the dust blackbody curve."
+ For the majority of the sources. however. the hot cust model was a Lot worse han the svnchrotron model. and so can not provide the peak in the optical/NIIt that is required to explain many of the observed SEDs.," For the majority of the sources, however, the hot dust model was a lot worse than the synchrotron model, and so can not provide the peak in the optical/NIR that is required to explain many of the observed SEDs."
+ A number of authors (Sandersetal.1989.forexample) have argued For the existence of a near-LR. bump. somewhere around 3//m. corresponding to blackbocdy emission from hot dust.," A number of authors \cite[for example]{sanders89} have argued for the existence of a near-IR bump, somewhere around $3\mu {\rm m}$, corresponding to blackbody emission from hot dust."
+ Sanders et shortceitesanderssO.. also mention the presence of a local münimum at lym. which they note for its “universality”.," Sanders et \\shortcite{sanders89}, also mention the presence of a local minimum at $1\mu {\rm m}$, which they note for its “universality”."
+ They interpret this to be due to the finite. sublimation temperature of dust. causing a drop in the blackbocdy emission. combined with the rise in Dux of thermal emission from hot (Z10. 000K) gas.," They interpret this to be due to the finite sublimation temperature of dust, causing a drop in the blackbody emission, combined with the rise in flux of thermal emission from hot $T\gtrsim10,000$ K) gas."
+ Any dip observed in our data occurs at. wavelengths which are too short to be attributed to hot dust., Any dip observed in our data occurs at wavelengths which are too short to be attributed to hot dust.
+ To observe a lim dip. we would need photometry at longer wavelengths.," To observe a $1\mu {\rm m}$ dip, we would need photometry at longer wavelengths."
+The .X rav emission from clusters of galaxies ollers a unique means for studying the physies of cosmic barvons and their connection with the processes of galaxy formation and evolution.,The $X$ –ray emission from clusters of galaxies offers a unique means for studying the physics of cosmic baryons and their connection with the processes of galaxy formation and evolution.
+ The first attempt to. mocel the LOCAL in the framework of the hierarchical. clustering scenario assumed. its thermocynamical properties to. be entirely determined. by. gravitational processes. such as adiabatie compression. during collapse and. shock heating »v supersonic gas accretion (Ixaiser. 1986).," The first attempt to model the ICM in the framework of the hierarchical clustering scenario assumed its thermodynamical properties to be entirely determined by gravitational processes, such as adiabatic compression during collapse and shock heating by supersonic gas accretion (Kaiser 1986)."
+ Since gravity does not have characteristic scales. for an Einsteinde-Sitter cosmology this model. predicts rich massive clusters and oor groups to appear as scaled. versions of cach other.," Since gravity does not have characteristic scales, for an Einstein–de-Sitter cosmology this model predicts rich massive clusters and poor groups to appear as scaled versions of each other."
+ Under the assumptions of emissivity dominated by free.free »emsstrahlung and of hydrostatic equilibrium of the gas. his model predicts Lyκπα|z)*7? for the shape aud evolution of the relation between Yo rav luminosity and gas emperature.," Under the assumptions of emissivity dominated by free–free bremsstrahlung and of hydrostatic equilibrium of the gas, this model predicts $L_X\propto T^2(1+z)^{3/2}$ for the shape and evolution of the relation between $X$ –ray luminosity and gas temperature."
+ Furthermore. if we define the gas entropy as ο.”n?3 (n.: electron. number density: e.g. Eke ot al.," Furthermore, if we define the gas entropy as $S=T/n_e^{2/3}$ $n_e$: electron number density; e.g. Eke et al."
+ 998). then the selfsimilar ICM has SxZ(112)E7.," 1998), then the self–similar ICM has $S\propto T(1+z)^{-2}$."
+ This simple model was quickly recognized. to fail at accounting for several observational facts: (a) the Ly T relation for nearby clusters is steeper than predicted. with Lxox17 lor TZ2 keV clusters (e... David et al.," This simple model was quickly recognized to fail at accounting for several observational facts: (a) the $L_X$ $T$ relation for nearby clusters is steeper than predicted, with $L_X\propto T^{\sim 3}$ for $T\magcir 2$ keV clusters (e.g., David et al."
+ 1993. White. Jones Forman 1997. Allen Fabian 1998. Markeviteh. 1998. Arnaud Evrarcl 1999). with a possible further steepening at the group scale. Z2;1 keV. (Ponman et al.," 1993, White, Jones Forman 1997, Allen Fabian 1998, Markevitch 1998, Arnaud Evrard 1999), with a possible further steepening at the group scale, $T\mincir 1$ keV (Ponman et al."
+ 1996. Llelsclon Ponman 2000): (b) no evidence of evolution [for its amplitude has been detected out to zz1 (e.g. Mushotzky Scharf 1997. Donahue et al.," 1996, Helsdon Ponman 2000); (b) no evidence of evolution for its amplitude has been detected out to $z\magcir 1$ (e.g., Mushotzky Scharf 1997, Donahue et al."
+ 1999. Fairley et al.," 1999, Fairley et al."
+ 2000. Della Coca et al.," 2000, Della Ceca et al."
+ 2000. Borgani et al.," 2000, Borgani et al."
+ 2001b. Holden et al.," 2001b, Holden et al."
+ 2002): (c) the gas density profile in central regions of cooler groups is relatively softer than in clusters and. corresponcinely. he entropy is higher than predicted by selfsimilar scaling (ος... Ponman. Cannon Navarro 1999. LlovdDavis et al.," 2002); (c) the gas density profile in central regions of cooler groups is relatively softer than in clusters and, correspondingly, the entropy is higher than predicted by self–similar scaling (e.g., Ponman, Cannon Navarro 1999, Lloyd–Davis et al."
+ 2000. Finoguenov et al.," 2000, Finoguenov et al."
+ 2002): (d) colder clusters contain a relatively smaller amount of gas than very hot ones (e.g. eumann Arnaud 2001)., 2002); (d) colder clusters contain a relatively smaller amount of gas than very hot ones (e.g. Neumann Arnaud 2001).
+ A common interpretation for such observational [acts requires nongravitational energy. input. which should have aken place during the past history of the ICM.," A common interpretation for such observational facts requires non–gravitational energy input, which should have taken place during the past history of the ICM."
+ This extra jeating would place the gas on a higher acliabat and sustain he gas during the gravitational collapse of the cluster dark matter (DAL) halo.preventing it from reaching high density in central regions and suppressing the X. ray emission.," This extra heating would place the gas on a higher adiabat and sustain the gas during the gravitational collapse of the cluster dark matter (DM) halo,preventing it from reaching high density in central regions and suppressing the $X$ –ray emission."
+ For à ixed specific amount of heating energy per gas particle. the elect is larger for the smaller svstems. while being negligible or more massive. hotter systems.," For a fixed specific amount of heating energy per gas particle, the effect is larger for the smaller systems, while being negligible for more massive, hotter systems."
+ Ixaiser (1991) suggested or the first time that gas preheating ancl subsequent aciabatie collapse turns into a cdilferential steepening of the Lx T relation at the cluster scales., Kaiser (1991) suggested for the first time that gas pre–heating and subsequent adiabatic collapse turns into a differential steepening of the $L_X$ $T$ relation at the cluster scales.
+ Evrard Henry (1991). avarro. Frenk White (1995) and Bower (1997) assumed he gas in the cluster core to have a minimum entropy level. which was established at some preheating epoch.," Evrard Henry (1991), Navarro, Frenk White (1995) and Bower (1997) assumed the gas in the cluster core to have a minimum entropy level, which was established at some pre–heating epoch."
+ Cavaliere. Menci Pozzi (1998. 1999) were able to predict the correct Ly E slope. both at. the cluster and at the group scale. by assuming a preheated gas at a temperature Z770.5 keV. which is shocked with different strengths. when falling into groups ancl clusters.," Cavaliere, Menci Tozzi (1998, 1999) were able to predict the correct $L_X$ $T$ slope, both at the cluster and at the group scale, by assuming a pre–heated gas at a temperature $T\sim 0.5$ keV, which is shocked with different strengths when falling into groups and clusters."
+ X. preheating producing an isentropic gas distribution has been assumed by Balogh. Babul Patton (1999) and Tozzi Norman (2001. TNOL hereafter).," A pre--heating producing an isentropic gas distribution has been assumed by Balogh, Babul Patton (1999) and Tozzi Norman (2001, TN01 hereafter)."
+ In. particular. “PNOL worked oul a series of predictions of observables propertics of the ICM with isentropic preheating. including the amount of energy feedback required to produce the correct relation and entropy threshold at the group scale.," In particular, TN01 worked out a series of predictions of observables properties of the ICM with isentropic pre–heating, including the amount of energy feedback required to produce the correct relation and entropy threshold at the group scale."
+ Mighenti Mathews (2001) cliscussed the cllect of heating the eas either when it has still to collapse within clusters and groups (external heating) or when it is already accreted (internal heating)., Brighenti Mathews (2001) discussed the effect of heating the gas either when it has still to collapse within clusters and groups (external heating) or when it is already accreted (internal heating).
+ Phey concluded that internal heating is more efficient al accounting at the same time for both the slope of the relation and the excess entropy in poor clusters., They concluded that internal heating is more efficient at accounting at the same time for both the slope of the relation and the excess entropy in poor clusters.
+ While most of such analysis agree that an extra heating energy of 0.5 1 keV per gas particle is required. vet no general consensus has been reached about its astrophysical origin.," While most of such analysis agree that an extra heating energy of $\sim 0.5$ –1 keV per gas particle is required, yet no general consensus has been reached about its astrophysical origin."
+ Supernovae (SNe) have been advocated by several authors as à possible source for preheating (oe. Wu. Fabian Nulsen 1998) and metal enriching (e.g. Loewenstein Alushotzky 1996. Renzini 1997. F'inoguenov Ponman 1999. Pipino et al.," Supernovae (SNe) have been advocated by several authors as a possible source for pre–heating (e.g., Wu, Fabian Nulsen 1998) and metal enriching (e.g. Loewenstein Mushotzky 1996, Renzini 1997, Finoguenov Ponman 1999, Pipino et al."
+ 2002) the ICM., 2002) the ICM.
+ Although some analvses showed that SNe can actually provide an adequate energy budget (Alenci Cavaliere 2000. LlovelDavis ct al.," Although some analyses showed that SNe can actually provide an adequate energy budget (Menci Cavaliere 2000, Lloyd–Davis et al."
+ 2000). other authors claimed that this requires a very high. possibly unrealistic. efficiency. for the thermalization of the released energy (Wu. Fabian Nulsen 2000. Ixravtsov Yepes 2000. Jower et al.," 2000), other authors claimed that this requires a very high, possibly unrealistic, efficiency for the thermalization of the released energy (Wu, Fabian Nulsen 2000, Kravtsov Yepes 2000, Bower et al."
+ 2001)., 2001).
+ In this case. alternative sources for LOCAL heating are required. Like SN from a primordial star population. the socalled. Pop LL stars (ee. Loewenstein 2001). and nuclear galactic activity associated with QSOs (e.g... Valageas Silk 1999. Me Namara ct al.," In this case, alternative sources for ICM heating are required, like SN from a primordial star population, the so--called Pop III stars (e.g., Loewenstein 2001), and nuclear galactic activity associated with QSOs (e.g., Valageas Silk 1999, Mc Namara et al."
+ 2000. Yamada Fujita 2001.Nath Rovehowelhury 2002).," 2000, Yamada Fujita 2001,Nath Roychowdhury 2002)."
+ As for Pop LL stars. one expects them not to significantly. heat the cilfuse medium. in order not to destroy a absorbers at lower redshift. 22;3. and not to overpollute the IGM. with metals.," As for Pop III stars, one expects them not to significantly heat the diffuse medium, in order not to destroy $\alpha$ absorbers at lower redshift, $z\mincir 3$, and not to overpollute the IGM with metals."
+ To date. only few attempts have. been pursued to aclelress in detail the cllects of non gravitational heating on X ray observable quantities with numerical hyerodyvnamical simulations (Navarro. Frenk White 1995. Lewis et al.," To date, only few attempts have been pursued to address in detail the effects of non gravitational heating on $X$ –ray observable quantities with numerical hydrodynamical simulations (Navarro, Frenk White 1995, Lewis et al."
+ 2000. Boreani ct al.," 2000, Borgani et al."
+ 2001a. Valdarnini 2002).," 2001a, Valdarnini 2002)."
+ Aalek. Evrard Mohr (2001) have realized a large set of moderate resolution simulations of clusters and groups. setting different entropy loors at very high redshift.," Bialek, Evrard Mohr (2001) have realized a large set of moderate resolution simulations of clusters and groups, setting different entropy floors at very high redshift."
+ Phev found that agreement with observational constraints on X. ray cluster scaling properties requires an entropy. Hloor of about LOO keV. em?. which. at heir heating redshift ο2 21. correspond to an extra energy of about 0.2 keV/part.," They found that agreement with observational constraints on $X$ –ray cluster scaling properties requires an entropy floor of about 100 keV $^2$ which, at their heating redshift $z_h\simeq 21$ correspond to an extra energy of about 0.2 keV/part."
+al c4250 in C-stars.,at $\sim 4250$ in C-stars.
+ The spectral types shown in Table 1 are based on the C-svstem classification by Ixeenan revised by Yamashitaetal.(1977)., The spectral types shown in Table 1 are based on the C-system classification by \citet{kee41} revised by \citet{yam77}.
+. Our stars fullfil the criteria for the SC abundance by Keenan.(1993). to be classified as J-stars: ratio of MCRC A4T44 to Me! \4737 in blue: ratios of PCPC AGLOL to BCPC AGI68A. BCPC AGIO2 to PHC! 6122 and Ον \G260 to CUN. A6206A.," Our stars fullfil the criteria for the $^{13}$ C abundance by \citet{kee93} to be classified as J-stars: ratio of $^{13}$ $^{12}$ C $\lambda 4744$ to $^{12}$ $^{12}$ C $\lambda 4737$ in blue: ratios of $^{12}$ $^{12}$ C $\lambda 6191$ to $^{13}$ $^{12}$ C $\lambda
+6168$, $^{13}$ $^{12}$ C $\lambda 6102$ to $^{12}$ $^{12}$ C $\lambda 6122$ and $^{13}$ $^{14}$ N $\lambda 6260$ to $^{12}$ $^{14}$ N $\lambda 6206$."
+ However. for some of them the spectrum is misleading. which might produce to a problematic spectral classification.," However, for some of them the spectrum is misleading, which might produce to a problematic spectral classification."
+" For instance, WZ Cas shows weaker CN and Co band absorptions than the other J-stars in the sample."," For instance, WZ Cas shows weaker CN and $_2$ band absorptions than the other J-stars in the sample."
+ It also shows very strong Na D lines and its spectrum does not look as crowded as the rest of the J-stars., It also shows very strong Na D lines and its spectrum does not look as crowded as the rest of the J-stars.
+ In fact. atomic lines in WZ Cas are more easily identified.," In fact, atomic lines in WZ Cas are more easily identified."
+ Indeed. this happens when (he C/O ratio in the atmosphere is very close to unity. a characteristic which delines a SC-(vpe carbon star.," Indeed, this happens when the C/O ratio in the atmosphere is very close to unity, a characteristic which defines a SC-type carbon star."
+ Since our C/O estimatein (his star is ~1.01. we believe that WZ Cas has to be considered a SCstar rather than a (vpical J-star.," Since our C/O estimatein this star is $\sim 1.01$, we believe that WZ Cas has to be considered a SC-star rather than a typical J-star."
+ Furthermore. it is (he most luminous star in our sample (AL)~ —6.44). which could indicate that WZ Cas belongs to a different population (massive) of C-stars or that it is in a different evolutionary status (more evolved) (han the rest of the J-stars in our sample.," Furthermore, it is the most luminous star in our sample $_{\rm{bol}}\sim -6.44$ ), which could indicate that WZ Cas belongs to a different population (massive) of C-stars or that it is in a different evolutionary status (more evolved) than the rest of the J-stars in our sample."
+ However argues (hat it is difficult to separate J-tvpe stus and SC stars only by spectroscopy results., However \citet{lor96} argues that it is difficult to separate J-type stars and SC stars only by spectroscopy results.
+ Modeling the dusty envelope of this star he concludes that WZ Cas could be a J-tvpe carbon star., Modeling the dusty envelope of this star he concludes that WZ Cas could be a J-type carbon star.
+ Less obviously. the same peculiarities are observed in (he spectrum of WX Cve.," Less obviously, the same peculiarities are observed in the spectrum of WX Cyg."
+ In fact. Ohnaka&Tsuji(1996) considered this star to be SC tvpe although. as [ar as we know. no other work in the literature classifies (his star as SC-tvpe.," In fact, \citet{ohn96} considered this star to be SC type although, as far as we know, no other work in the literature classifies this star as SC-type."
+ For this star we do not have a clear opinion., For this star we do not have a clear opinion.
+ On the otherhand. FO Ser is also classified in the literature as alate RG type carbon star (Stephenson 1989)..," On the otherhand, FO Ser is also classified in the literature as a late R6 type carbon star \citep{ste89}. ."
+ We consider it a tvpical J-star but we, We consider it a typical J-star but we
+"Despite the comprehensive nature of their model. ? have vet to identify au appropriate mass loss nechanisin and in thei own words. ""the outer οσο of the disk. rq. is not well defined otler iu df be much further out than the centrifueal radius.","Despite the comprehensive nature of their model, \citet{ward10} have yet to identify an appropriate mass loss mechanism and in their own words, “the outer edge of the disk, $r_{\rm d}$, is not well defined other than it be much further out than the centrifugal radius”."
+ Caven a moderate external ΕΝ flux roni either the central star or nearby hieh-auass stars. photoevaporation could provide a natural uechanisn for both mass loss and truncation at 1e outer edge of eireumplanetary disks.," Given a moderate external FUV flux from either the central star or nearby high-mass stars, photoevaporation could provide a natural mechanism for both mass loss and truncation at the outer edge of circumplanetary disks."
+ ? asstune the same cireunplauetary disk evolution for both Jupiter aud Saturn aud that 1e dichotomy preseut in their satellite svstenis is explained by thestochastic timing of foriiation/iuieration and the depletion of the solar nebula (7).., \citet{canup02} assume the same circumplanetary disk evolution for both Jupiter and Saturn and that the dichotomy present in their satellite systems is explained by the stochastic timing of formation/migration and the depletion of the solar nebula \citep{canup06}.
+ In contrast. 2 surmise that the difference in the satellite systems of these two planets is the result of very different disk evolution scenarios.," In contrast, \citet{sasaki10} surmise that the difference in the satellite systems of these two planets is the result of very different disk evolution scenarios."
+ Auneneal siumlations show that Jupiter is large enough to open a complete gap in the solar nebula cutting off infall onto the Jovian subuchula on a 10?10! year timescale (2)., Numerical simulations show that Jupiter is large enough to open a complete gap in the solar nebula cutting off infall onto the Jovian subnebula on a $10^2-10^4$ year timescale \citep{sasaki10}.
+ In coutrast. Satin is too small to have opened a complete eap iu he solar nebula and iufall outo the Saturnia subuebula would have halted ou the ~109 vear iuescale for the dissipation of the solar nebula as a whole.," In contrast, Saturn is too small to have opened a complete gap in the solar nebula and infall onto the Saturnian subnebula would have halted on the $\sim 10^6$ year timescale for the dissipation of the solar nebula as a whole."
+" Although the assumption that Jupiter opened a gap and Saturn dil not depends on he asstuned: viscosity and scale height of the otoplanetary disk. the critical mass for opening a gap is larger in the onter regions of the disk,"," Although the assumption that Jupiter opened a gap and Saturn did not depends on the assumed viscosity and scale height of the protoplanetary disk, the critical mass for opening a gap is larger in the outer regions of the disk."
+" It is a reasonable asstuuption using standard disk om;uneters, aud further supported by the curent nasses of Jupiter and Saturn."," It is a reasonable assumption using standard disk parameters, and further supported by the current masses of Jupiter and Saturn."
+ For these reasons. ∙?. assune in. their. models that Jupiter. opened a eap in the solar nebula aud Saturn did uot.," For these reasons, \citet{sasaki10} assume in their models that Jupiter opened a gap in the solar nebula and Saturn did not."
+ ? have recently published results from a suite of simmuations iu which the growth aud dynamical evolution of proto-satellite euibryos was modeled., \citet{sasaki10} have recently published results from a suite of simulations in which the growth and dynamical evolution of proto-satellite embryos was modeled.
+" They seek to explain why the Jovian regular satellite system consists of four nearly equal mass satellites whereas the Saturuian system, contains oue large satellite.", They seek to explain why the Jovian regular satellite system consists of four nearly equal mass satellites whereas the Saturnian system contains one large satellite.
+" They propose that the Jovian satellite systein may have been ""frozen in place when Jupiter grow suthcicuthy large that a gap was opened in the solar uebula.", They propose that the Jovian satellite system may have been “frozen” in place when Jupiter grew sufficiently large that a gap was opened in the solar nebula.
+ Tu the Saturn system. where oulv au incomplete gap may have formed. the slower shutoff timescale for material infalliug from the solar nebula would have allowed its satellites to coutinue to dynamically evolve.," In the Saturnian system, where only an incomplete gap may have formed, the slower shutoff timescale for material infalling from the solar nebula would have allowed its satellites to continue to dynamically evolve."
+ The typical cud result is significant depletion of solids in the iuner disk aud the retention of a single laree satellite iu the outer disk that is similar to Titan., The typical end result is significant depletion of solids in the inner disk and the retention of a single large satellite in the outer disk that is similar to Titan.
+ This outcome may also be a natural result of the inside-out clearing of the solar nebula due to photoevaporation by the Sun (see discussion iu Section | below)., This outcome may also be a natural result of the inside-out clearing of the solar nebula due to photoevaporation by the Sun (see discussion in Section \ref{sec:discussion} below).
+ 7. ve able to produce four or five similarly sized satellites in the Jovian svstenài in SO of their ruus., \citet{sasaki10} are able to produce four or five similarly sized satellites in the Jovian system in $80\%$ of their runs.
+ Whereas. in the Saturni system. only one large satellite remains in 7054 of their runs.," Whereas, in the Saturnian system only one large satellite remains in $70\%$ of their runs."
+ Their models however rely ou a rapid dispersal mechanisni for cireunplanetaryv eas once the sap iu the solar nebula has been opened., Their models however rely on a rapid dispersal mechanism for circumplanetary gas once the gap in the solar nebula has been opened.
+ Photoevaporation could provide just such aduechanisin as well as help to determine a natural outer disk boundary., Photoevaporation could provide just such a mechanism as well as help to determine a natural outer disk boundary.
+ As with ?.. the models of ? rely on an ad hoc outer disk bowuclary.," As with \citet{ward10}, the models of \citet{sasaki10} rely on an ad hoc outer disk boundary."
+ Receut simmlatious have investigated the tidal truucation of circuplauctary disks (?).., Recent simulations have investigated the tidal truncation of circumplanetary disks \citep{martin10}.
+ These tidal truncation simulations produce disks that are truncated at a radius that occurs at ~(0.1ry. where rg is the planet Will radius.," These tidal truncation simulations produce disks that are truncated at a radius that occurs at $\sim 0.4\ r_{\rm H}$, where $r_{\rm H}$ is the planet Hill radius."
+" where à and AL, are the planct’s semi-major axis and mass aud AL. is the mass of the Sun.", where $a$ and $M_{\rm p}$ are the planet's semi-major axis and mass and $M_{\odot}$ is the mass of the Sun.
+ This outer disk radius is. however. too large to explain the compact configuration of the regular satellite systems of Jupiter aud Saturn which extend to less than 0.06ry.," This outer disk radius is, however, too large to explain the compact configuration of the regular satellite systems of Jupiter and Saturn which extend to less than $0.06\ r_{\rm H}$."
+ The primary objective ofthe preseut paper is to showdose how. a nominal: FUV ---flux can photoevaporatei. outer portions of cieunmplanetary iedisks., The primary objective of the present paper is to show how a nominal FUV flux can photoevaporate the outer portions of circumplanetary disks.
+ ?7? the demonstrate how photoevaporation creates a subsouicsubs outfow of gieas in the disk: that is: well inside the gravitational radius. re. where the retinal velocityof the hot disk atinosphere equals ie planet's escape velocity.," \citet{adams04} demonstrate how photoevaporation creates a subsonic outflow of gas in the disk that is well inside the gravitational radius, $r_{\rm g}$, where the thermal velocity of the hot disk atmosphere equals the planet's escape velocity."
+ We apply the ? oXhotoevaporatiou model to ciremuplanctary disks and find that the disks ave truncated well inside 16 gravitational radius. ry.," We apply the \citet{adams04} photoevaporation model to circumplanetary disks and find that the disks are truncated well inside the gravitational radius, $r_{\rm g}$."
+ It is important to note iat plotocvaporation ouly needs to remove gas Yon the plauct’s Till sphere in order to tuncate le cirennplunetary disk: it does not need to remove the eas from the solar svsteun., It is important to note that photoevaporation only needs to remove gas from the planet's Hill sphere in order to truncate the circumplanetary disk; it does not need to remove the gas from the solar system.
+ The ciffereuce in solar flux. duc to the difference I seni-najor axes. Ίαν further account for," The difference in solar flux, due to the difference in semi-major axes, may further account for"
+to compute 10 years of stellar time despite its lower numerical resolution.,to compute 10 years of stellar time despite its lower numerical resolution.
+ The CPU time needed for non-gray runs does scale almost linearly with the number of wavelength groups., The CPU time needed for non-gray runs does scale almost linearly with the number of wavelength groups.
+ 3D simulations start with an initial model that has an estimated radius. a certain envelope mass. a certain potential profile. and a prescribed luminosity as described above.," 3D simulations start with an initial model that has an estimated radius, a certain envelope mass, a certain potential profile, and a prescribed luminosity as described above."
+ During the run. the internal structure relaxes to something not to far away from the mitial guess.," During the run, the internal structure relaxes to something not to far away from the initial guess."
+ We needed several trials to get the right mitial model for the latest set., We needed several trials to get the right initial model for the latest set.
+ The average final stellar parameters are determined once the simulation has ended (Fig. 2))., The average final stellar parameters are determined once the simulation has ended (Fig. \ref{structure}) ).
+ For this purpose. we adopted the method reported in ? (hereafter Paper D).," For this purpose, we adopted the method reported in \cite{2009A&A...506.1351C} (hereafter Paper )."
+ The quantities are averages over spherical shells and over time. and the errors are one-sigma fluctuations with respect to the average over time.," The quantities are averages over spherical shells and over time, and the errors are one-sigma fluctuations with respect to the average over time."
+ We computed the average temperature and luminosity over spherical shells. 7(7). and Lir).," We computed the average temperature and luminosity over spherical shells, $T(r)$, and $L(r)$."
+ Then. we search for the radius R for which where c is the Stefan-Boltzmann constant.," Then, we search for the radius $R$ for which where $\sigma$ is the Stefan-Boltzmann constant."
+ Eventually. the effective temperature is Tay=ΤΙΝ).," Eventually, the effective temperature is $T_{\rm{eff}}=T(R)$."
+ As already pointed out in PaperI. the gray model st35em03n07 (top row of Fig. 2))," As already pointed out in Paper, the gray model st35gm03n07 (top row of Fig. \ref{structure}) )"
+ shows a drift in the first 2 years (-O0.5% per year) before stabilisation in the last 2.5 years., shows a drift in the first 2 years $-$ $\%$ per year) before stabilisation in the last 2.5 years.
+ This drift in also visible in the radius of the non-gray model st35em03n13 (second row) even though it is weaker (less than 0.3%)., This drift in also visible in the radius of the non-gray model st35gm03n13 (second row) even though it is weaker (less than $\%$ ).
+ The luminosity fluctuations are of the order of 1% and the temperature variations <0.3% for both simulations., The luminosity fluctuations are of the order of $\%$ and the temperature variations $<0.3\%$ for both simulations.
+ It must be noted that the last snapshot of the gray simulation has been used as the initial snapshot for the non-gray model., It must be noted that the last snapshot of the gray simulation has been used as the initial snapshot for the non-gray model.
+ The two hotter models. st36gm00n04 and st36gm00n05. are also somehow related because they both use a gray treatment of opacities and st36gm00n05 has been started using the last snapshot of st36gmOOnO4.," The two hotter models, st36gm00n04 and st36gm00n05, are also somehow related because they both use a gray treatment of opacities and st36gm00n05 has been started using the last snapshot of st36gm00n04."
+ The temperature (0.5%) and luminosity (2.2%) fluctuations of the higher resolution model are slightly larger than the temperature (0.4%) and lumimosity (1.59) fluctuations of st36em00n04., The temperature $\%$ ) and luminosity $\%$ ) fluctuations of the higher resolution model are slightly larger than the temperature $\%$ ) and luminosity $\%$ ) fluctuations of st36gm00n04.
+ None of the simulation shows a drift of the radius in the last years of evolution., None of the simulation shows a drift of the radius in the last years of evolution.
+ The bottom row in Fig., The bottom row in Fig.
+ 2. show the ratio between the turbulent pressure and the gas pressure (defined as in Sect. 22))., \ref{structure} show the ratio between the turbulent pressure and the gas pressure (defined as in Sect. \ref{temp_sect}) ).
+ This quantity shows that in the outer layers. just above the stellar radius. the turbulent pressure plays a significant role for the average pressure stratification and the radial velocities resulting from the vigorous convection are supersonic.," This quantity shows that in the outer layers, just above the stellar radius, the turbulent pressure plays a significant role for the average pressure stratification and the radial velocities resulting from the vigorous convection are supersonic."
+ Figure 3. displays some quantities spatially averaged over spherical shells for a snapshot of the gray simulation st35gmO3n07 in Table 1.., Figure \ref{1dquantities} displays some quantities spatially averaged over spherical shells for a snapshot of the gray simulation st35gm03n07 in Table \ref{simus}.
+ Figure 4. shows three-dimensional views of some quantities for the same snapshot and simulation., Figure \ref{cobold_quantities} shows three-dimensional views of some quantities for the same snapshot and simulation.
+ Radiation is of primary importance for many aspects of convection and the envelope structure in a RSG., Radiation is of primary importance for many aspects of convection and the envelope structure in a RSG.
+ It does not only cool the surface to provide à somewhat unsharp outer boundary for the convective heat transport., It does not only cool the surface to provide a somewhat unsharp outer boundary for the convective heat transport.
+ It also contributes significantly to the energy transport 1n the interior (top left panel of Fig. 3) ).," It also contributes significantly to the energy transport in the interior (top left panel of Fig. \ref{1dquantities}) ),"
+ Where convection never carries close to 1006€ of the luminosity., where convection never carries close to $100\%$ of the luminosity.
+ In the optically thick part the stratification is slightly far from radiative equilibrium and the entropy jump from the photosphere to the layers below ts fairly large (entropy s in Fig., In the optically thick part the stratification is slightly far from radiative equilibrium and the entropy jump from the photosphere to the layers below is fairly large (entropy $s$ in Fig.
+ 3 and Fig. 4))., \ref{1dquantities} and Fig. \ref{cobold_quantities}) ).
+ The He II/He tonization zone is visible in the entropy structure as a small minimum near the surface before the normal steep entropy decrease., The He /He ionization zone is visible in the entropy structure as a small minimum near the surface before the normal steep entropy decrease.
+ The opacity peak at around 0000K just below the photosphere causes a very steep temperature Jump (temperature 7 and opacity « in Fig. 4)), The opacity peak at around 000K just below the photosphere causes a very steep temperature jump (temperature $T$ and opacity $\kappa$ in Fig. \ref{cobold_quantities}) )
+ which is very prominent on top of upflow regions (opacity « in Fig. 3))., which is very prominent on top of upflow regions (opacity $\kappa$ in Fig. \ref{1dquantities}) ).
+ This causes a density inversion (densityp in Fig., This causes a density inversion (density $\rho$ in Fig.
+ 4+ and Fig.8)) which is a sufficient condition of convective instability resolved. the entropy drop occurs in a very thin layer. while the smearing due to the averaging procedure over nonstationary up- anddownflows leads to the large apparent extent of the mean superadiabatic layer.," \ref{cobold_quantities} and \ref{temperaturedensity_profile}) ) which is a sufficient condition of convective instability resolved, the entropy drop occurs in a very thin layer, while the smearing due to the averaging procedure over nonstationary up- and downflows leads to the large apparent extent of the mean superadiabatic layer."
+ The local radiative relaxation time scale in the photosphere is much shorter than a typical hydrodynamical time scale (time, The local radiative relaxation time scale in the photosphere is much shorter than a typical hydrodynamical time scale (time
+not included in (he present study so that the hydrodynamieal equilibrium of halo gas can be obtained [rom isolated cluster models.,not included in the present study so that the hydrodynamical equilibrium of halo gas can be obtained from isolated cluster models.
+ We can expect that the shocked gaseous regions in models with radiative cooling have significantlv higher gas densities ancl pressure. and therefore galaxies in the models can be more strongly influenced by cluster merging during (heir passage of the shocked regions.," We can expect that the shocked gaseous regions in models with radiative cooling have significantly higher gas densities and pressure, and therefore galaxies in the models can be more strongly influenced by cluster merging during their passage of the shocked regions."
+" We adopt a representative value of M,=0.136. (e.g... MeCarthy et al."," We adopt a representative value of $M_{\rm g}=0.136 M_{\rm cl}$ (e.g., McCarthy et al."
+ 2003)., 2008).
+ Galaxies in a cluster are represented by collisionless particles and their spatial distribution follows the NEW profile with e=3., Galaxies in a cluster are represented by collisionless particles and their spatial distribution follows the NFW profile with $c=3$.
+ The canonical Schechter function is adopted [or eeneraling a galaxy luminosity/mass function for Iuminosities ranging from Q.0LL* to 2.5£* in a cluster., The canonical Schechter function is adopted for generating a galaxy luminosity/mass function for luminosities ranging from $0.01L^{\ast}$ to $2.5L^{\ast}$ in a cluster.
+ The total mass (thus munber) of galaxiesin a cluster with A4 is determined by 1e niass-to-light-ratio. that itselE is dependent on M4 (Alarinoni Hudson 2002).," The total mass (thus number) of galaxiesin a cluster with $M_{\rm cl}$ is determined by the mass-to-light-ratio, that itself is dependent on $M_{\rm cl}$ (Marinoni Hudson 2002)."
+ Therefore.as an example. a cluster with A4=10!M. has 14 galaxies.," Therefore,as an example, a cluster with $M_{\rm cl}=10^{14} {\rm M}_{\odot}$ has 114 galaxies."
+ Cluster member galaxies in a ‘luster have an isotropic velocity dispersion just as the dark matter of the cluster does., Cluster member galaxies in a cluster have an isotropic velocity dispersion just as the dark matter of the cluster does.
+ Galaxies in the present study have no halo gas that might well shield ISM of disk galaxies and (hus weaken (he possible physical effects of ICM (e.g.. triggering star formation) on the ISAL.," Galaxies in the present study have no halo gas that might well shield ISM of disk galaxies and thus weaken the possible physical effects of ICM (e.g., triggering star formation) on the ISM."
+ Previous numerical simulations. however. showed that halo gas around. disk. galaxies can be elliciently stripped bx ICM even in isolated clusters (e.g.. Dekki 2009): the halo gas would not signilicantlv. weaken the effects of ICM on galaxy. evolution in merging groups and clusters.," Previous numerical simulations, however, showed that halo gas around disk galaxies can be efficiently stripped by ICM even in isolated clusters (e.g., Bekki 2009): the halo gas would not significantly weaken the effects of ICM on galaxy evolution in merging groups and clusters."
+ The ISM of galaxies is not included in the present models so Chat the details of star formation caused by cluster merging can not be properly investigated., The ISM of galaxies is not included in the present models so that the details of star formation caused by cluster merging can not be properly investigated.
+ The present study thus assumes (hat if pressure of ICAL around galaxies can become high enough. then the star formation histories can be significantly changed irrespective of galaxy. properties.," The present study thus assumes that if pressure of ICM around galaxies can become high enough, then the star formation histories can be significantly changed irrespective of galaxy properties."
+ This assumption would be reasonable. if GMC properties do not depend strongly on galaxy ones.," This assumption would be reasonable, if GMC properties do not depend strongly on galaxy ones."
+" The relative position and velocity of CL2 with respect to CL1 are describedas (Αν. Y. Z.) and (U,. Ἐν. UV). respectively."," The relative position and velocity of CL2 with respect to CL1 are describedas $X_{\rm r}$, $Y_{\rm r}$, $Z_{\rm r}$ ) and $U_{\rm r}$, $V_{\rm r}$ , $W_{\rm r}$ ), respectively."
+" For all models described in (he present study. Owhere H4 is the size of CLI). Z=0. U,=0. and Wy=0."," For all models described in the present study, $X_{\rm r}=2R_{\rm cl}$ (where $R_{\rm cl}$ is the size of CL1), $Z_{\rm r}=0$, $U_{\rm r}=0$, and $W_{\rm r}=0$."
+" Therelore Y, is the inpact parameter of cluster merging and Ἐν is the initialrelative velocity of merging two clusters.", Therefore $Y_{\rm r}$ is the impact parameter of cluster merging and $V_{\rm r}$ is the initialrelative velocity of merging two clusters.
+ Cluster merging processes and subsequent ICM. evolution depend strongly on the [our parameters. Ma. ma. Y. and Vi.," Cluster merging processes and subsequent ICM evolution depend strongly on the four parameters, $M_{\rm cl}$ , $m_{2}$ , $Y_{\rm r}$ , and $V_{\rm r}$."
+ Although we have conducted a large parameter study. we only show the results of the models with M4=LOMAL.. 0.1xoma<LI. M=0.544. and V;= —Vy.," Although we have conducted a large parameter study, we only show the results of the models with $M_{\rm cl}=10^{14} {\rm M}_{\odot}$, $0.1\leq m_{2} \leq 1$, $Y_{\rm r}=0.5 R_{\rm cl}$, and $V_{\rm r} = -V_{\rm cl}$ ."
+ Parameter dependencies of (he results. will be described in detail in our Iorthcoming papers (Bekki οἱ al., Parameter dependencies of the results will be described in detail in our forthcoming papers (Bekki et al.
+ 2009)., 2009).
+" Here we focus in particular on the ""standard model with M4=10. M... ms= 0.25. Ry=2.0 Mpe. Vy=595 km |. dy=2.0Gvr. X;=L0 Mpc. and V;——595 km LH because itshows oneofthetypical behaviors of the time evolution of the external gas"," Here we focus in particular on the “standard model” with $M_{\rm cl}=10^{14} {\rm M}_{\odot}$ , $m_{2}=0.25$ $R_{\rm cl}=2.0$ Mpc, $V_{\rm cl}=595$ km $^{-1}$ , $t_{\rm cl}=2.0$Gyr, $Y_{\rm r}=1.0$ Mpc, and $V_{\rm r} = -595$ km $^{-1}$ , because itshows oneofthetypical behaviors of the time evolution of the external gas"
+One artifact of optimizing the mixture for Criterion 1 is (hat. on average. smaller models are selected. compared with the mixtures selected by DIC.,"One artifact of optimizing the mixture for Criterion 1 is that, on average, smaller models are selected, compared with the mixtures selected by BIC."
+ For ESOLV on average 34 components were selected by (he former approach. while on average 61 components were selected by the latter.," For ESOLV on average 34 components were selected by the former approach, while on average 61 components were selected by the latter."
+ Furthermore. an average of 4.1 nonpredefined components were used with Criterion 1 model selection while an average of 8.9 were used with BIC.," Furthermore, an average of 4.1 nonpredefined components were used with Criterion 1 model selection while an average of 8.9 were used with BIC."
+ For the SDSS dala an average of 61 components. including 4.8 nonpredefined components. were selected using Criterion 1.," For the SDSS data an average of 61 components, including 4.8 nonpredefined components, were selected using Criterion 1."
+ Finding the best SDSS model by BIC we needed on average 71 components. including 6.4 nonpredefined components.," Finding the best SDSS model by BIC we needed on average 71 components, including 6.4 nonpredefined components."
+ While we learned models with up to 80 components. in some cases for the SDSS data the best models were using close to 80 components.," While we learned models with up to 80 components, in some cases for the SDSS data the best models were using close to 80 components."
+ This suggests it may be reasonable to evaluate solutions with even more components., This suggests it may be reasonable to evaluate solutions with even more components.
+ While the mean number of components selected by BIC was greater than that selected. according to Criterion 1. (he variance in the number of selected components is much ereater for selection according to Criterion 1.," While the mean number of components selected by BIC was greater than that selected according to Criterion 1, the variance in the number of selected components is much greater for selection according to Criterion 1."
+ This is consistent with the results in Figure 2. which indicate that. for best Criterion 1 performance. the number of components is significantly correlated with (he mass of the unknown classes (which varies greatly since the classes ave far from equally likely).," This is consistent with the results in Figure 2, which indicate that, for best Criterion 1 performance, the number of components is significantly correlated with the mass of the unknown classes (which varies greatly since the classes are far from equally likely)."
+ This further means that. in some cases. when the mass of the unknown classes is large. Criterion 1 selects more components than BIC.," This further means that, in some cases, when the mass of the unknown classes is large, Criterion 1 selects more components than BIC."
+ For example. for the ESOLV cata. class 2 occurs 46% of the time.," For example, for the ESOLV data, class 2 occurs $46\%$ of the time."
+ When this class is taken as unknown. BIC selected 63 components. while Criterion 1 selected 70.," When this class is taken as unknown, BIC selected 63 components, while Criterion 1 selected 70."
+ Another factor that influences model accuracy is the fact that the learning objective [function £ is multimodal. with significant potential for finding suboptimal local maxima. rather than the global maxinnun.," Another factor that influences model accuracy is the fact that the learning objective function ${\cal L}$ is multimodal, with significant potential for finding suboptimal local maxima, rather than the global maximum."
+ At each moclel order. we generated several solutions based on different initializalions and picked the one with greatest log-likelihood.," At each model order, we generated several solutions based on different initializations and picked the one with greatest log-likelihood."
+ However. there is anecdotal evidence in our results that we may only be finding locally optimal solutions al each model order.," However, there is anecdotal evidence in our results that we may only be finding locally optimal solutions at each model order."
+ Referring back to Figure 2 we see that the best model for unknown Classes 0 2 contains only 13 mixture components., Referring back to Figure 2 we see that the best model for unknown classes 0 2 contains only 13 mixture components.
+ This clearly is not in keeping with the trend that more mixture components are needed to explain the data with larger mass fraction of unknown classes., This clearly is not in keeping with the trend that more mixture components are needed to explain the data with larger mass fraction of unknown classes.
+ It appears in (this case that a particularly good solution was found., It appears in this case that a particularly good solution was found.
+ This likewise suggests that. at other orders. suboptimal. local maximum solutions were found.," This likewise suggests that, at other orders, suboptimal, local maximum solutions were found."
+ The Criterion 2 error is à measure of how well the algorithm can classify objects within the newly found classes., The Criterion 2 error is a measure of how well the algorithm can classify objects within the newly found classes.
+ This is à very. hard problem because we are asking (he algorithm {ο do two things., This is a very hard problem because we are asking the algorithm to do two things.
+ First. determine if some mixture components (the nonpredelimed components) are needed to describe objects that do not fit into the existing known class structure.," First, determine if some mixture components (the nonpredefined components) are needed to describe objects that do not fit into the existing known class structure."
+ Second. partition these objects correctly between the unknown class components.," Second, partition these objects correctly between the unknown class components."
+ This second step is effectivelv looking for substructure in the newly discovered classes., This second step is effectively looking for substructure in the newly discovered classes.
+ For the ESOLY data we found on average about a τος Criterion 2 error compared with, For the ESOLV data we found on average about a $70\%$ Criterion 2 error compared with
+agreement with COMPTEL.,agreement with COMPTEL.
+" 'To test the dependence of the sub-MeV emission on the assumed ISRF model, we calculate the diffuse emission for the case of no metallicity gradient."," To test the dependence of the sub-MeV emission on the assumed ISRF model, we calculate the diffuse emission for the case of no metallicity gradient."
+ Figure 5 shows the spectrum of the diffuse emission for this case., Figure \ref{fig:diffuse_no_metallicity} shows the spectrum of the diffuse emission for this case.
+ The major change in the ISRF is a ~30% reduction in the infrared emission which is due to the smaller amount of, The major change in the ISRF is a $\sim$ reduction in the infrared emission which is due to the smaller amount of
+FLAG = 0 and restricted our analysis to the enerev range 0.512 keV. The average PN source count rate in this enerev range. excluding N-ray bursts; was 230 + for observation 1. [70s + for observation 2. aud 620s + for observation 3.,"FLAG = 0 and restricted our analysis to the energy range 0.5–12 keV. The average PN source count rate in this energy range, excluding X-ray bursts, was 230 $^{-1}$ for observation 1, 470 $^{-1}$ for observation 2, and 620 $^{-1}$ for observation 3."
+ These rates are below the SOQO + timing uode pile-up lait above which the spectral response would be deteriorated by photon pile-up Users Haudbook)., These rates are below the 800 $^{-1}$ timing mode pile-up limit above which the spectral response would be deteriorated by photon pile-up User's Handbook).
+ The background count rate was less han of the source count rate aud did not exhibit VAienificaut faring., The background count rate was less than of the source count rate and did not exhibit significant flaring.
+ We fud a strong lhne-like feature iu he background spectrum near 0.15 keV. We interpret lis feature as the electronic noise peak for double eveuts hat is also preseut for PN inaagiug modes but is shifted o higher energies iu timine mode., We find a strong line-like feature in the background spectrum near 0.45 keV. We interpret this feature as the electronic noise peak for double events that is also present for PN imaging modes but is shifted to higher energies in timing mode.
+ IU 1636-536 was observed withRXTE (7) siuultaueouslv with each observation.," 4U 1636-536 was observed with \citep{1993A&AS...97..355B}
+ simultaneously with each observation."
+ The TEobservations where somewhat longer than the observations and completely overlapped with the EPIC) PN exposures., The observations where somewhat longer than the observations and completely overlapped with the EPIC PN exposures.
+ No earth occultation occurred diving these observations aud the target was observed continuously withRATE., No earth occultation occurred during these observations and the target was observed continuously with.
+ Iu this paper we consider only the data obtained with the TWENTE detector C2).., In this paper we consider only the data obtained with the HEXTE detector \citep{1998ApJ...496..538R}.
+ We did not use the PCA data because of their lower euerev resolution compared to the EPIC PN data aud because the PCA spectrum in the 312 keV cherev range showed au up to lueher flux than the EPIC PN spectrum., We did not use the PCA data because of their lower energy resolution compared to the EPIC PN data and because the PCA spectrum in the 3–12 keV energy range showed an up to higher flux than the EPIC PN spectrum.
+ Similar excesses of the PCA flux compared to other instrmucuts have previously bec reported (e.g.77).," Similar excesses of the PCA flux compared to other instruments have previously been reported \citep[e.g.][]{1999ApJ...512..892T,2003A&A...411L.343C}."
+ It was not possible to correct for the higher PCÀ flus by inchiding a iultiplicative cross-calibration factor because the flux excess is strongly energv dependent (~30% at 3 keV and —10 at 10 keV)., It was not possible to correct for the higher PCA flux by including a multiplicative cross-calibration factor because the flux excess is strongly energy dependent $\sim$ at 3 keV and $\sim$ at 10 keV).
+ We included ii our analysis only those ΠΤΙ. data that were taken simultaneously to the EPIC PN data., We included in our analysis only those HEXTE data that were taken simultaneously to the EPIC PN data.
+ The TWENTE data cover the cutive duration of the EPIC PN exposures with the exception of a single «1 ks data gap in cach observation., The HEXTE data cover the entire duration of the EPIC PN exposures with the exception of a single $<$ 1 ks data gap in each observation.
+ The TWENTE background rate and the telescope pointing were stable duiug both observations., The HEXTE background rate and the telescope pointing were stable during both observations.
+ We extracted TENTE data for the energy range 20100 keV using the REX data analysis script., We extracted HEXTE data for the energy range 20–100 keV using the REX data analysis script.
+ For observation 1. we imcluded in our analysis data frou both VENTE clusters.," For observation 1, we included in our analysis data from both HEXTE clusters."
+ For observations 2 aud 3. we ouly used the data from cluster DB. because cluster A was uo longer able to obtain background measurements.," For observations 2 and 3, we only used the data from cluster B because cluster A was no longer able to obtain background measurements."
+ For the spectral analysis. we removed all X-ray bursts from the data (three for observation 1 aud oue cach for observations 2 and 3).," For the spectral analysis, we removed all X-ray bursts from the data (three for observation 1 and one each for observations 2 and 3)."
+ The EPIC PN spectra were binned at about 1/3 of the FWIIAL detector resolution aud were fitted simultancously with the TWENTE data., The EPIC PN spectra were binned at about 1/3 of the FWHM detector resolution and were fitted simultaneously with the HEXTE data.
+ We did not iuclude a multiplicative cross-calibration factor between the two instruments., We did not include a multiplicative cross-calibration factor between the two instruments.
+ The spectral fitting was performed using NSPEC version 12.1 (7).., The spectral fitting was performed using XSPEC version 12.4 \citep{1996ASPC..101...17A}.
+ For the EPIC PN spectrmm from observation 3. we fud siguificaut residuals between data aud model at ~2.3 keV near the iustriuental Au edge.," For the EPIC PN spectrum from observation 3, we find significant residuals between data and model at $\sim$ 2.3 keV near the instrumental Au edge."
+ We attribute these residuals to a sanall offset in the photon cucrey cleterumination caused be a slightly inaccurate charee trauster efficiency. (CTE) calibration iu tinue mode (seee.g.?).., We attribute these residuals to a small offset in the photon energy determination caused be a slightly inaccurate charge transfer efficiency (CTE) calibration in timing mode \citep[see e.g.][]{2007A&A...461.1049S}.
+ We corrected for this energv offset bv introducing a mnultiplicative cain correction factor of 1.0020 which we determined from a fit., We corrected for this energy offset by introducing a multiplicative gain correction factor of 1.0020 which we determined from a fit.
+ This ein factor. which was fixed in our further spectral analysis. completely removed the residuals near the iustrumental Au edge.," This gain factor, which was fixed in our further spectral analysis, completely removed the residuals near the instrumental Au edge."
+ No eain correction was necessary for the spectra from observations 1 aud 2., No gain correction was necessary for the spectra from observations 1 and 2.
+ Initial fitting of the spectra from all three observations showed significant residuals between data and model below 2 keV near the K-shell absorption edges of ο. Ne. Me. and $i," Initial fitting of the spectra from all three observations showed significant residuals between data and model below 2 keV near the K-shell absorption edges of O, Ne, Mg, and Si."
+ We attribute these residuals to an overly simplified description of the interstellar absorption edees by current NSPEC models., We attribute these residuals to an overly simplified description of the interstellar absorption edges by current XSPEC models.
+ While these absorption models are generally adequate for EPIC PN spectra. the data preseuted here has an exceptionally high sieual-to-nolse ratio so that very small discrepancies between data aud model become apparent.," While these absorption models are generally adequate for EPIC PN spectra, the data presented here has an exceptionally high signal-to-noise ratio so that very small discrepancies between data and model become apparent."
+ We coufirmed thiρα interpretation with the RGS spectra which clearly show a conmlex structure with several narrow absorption features near the IK-shell absorption edges of oxvecn aud neon., We confirmed this interpretation with the RGS spectra which clearly show a complex structure with several narrow absorption features near the K-shell absorption edges of oxygen and neon.
+ This edge structure is not adequately described by any of the absorption models in NSPEC which model the edees as an exponential with a low-energv cut-off., This edge structure is not adequately described by any of the absorption models in XSPEC which model the edges as an exponential with a low-energy cut-off.
+ The structure of the oxvecn aud ucon ER-edees iu the RCS spectra is simular to the structure found by ?? for several A-rav binaries Guchiding IU 1636-536) and interpreted as absorption bv neutral aud ionized oxveen and neon oeo1 the interstellar uiedimm.," The structure of the oxygen and neon K-edges in the RGS spectra is similar to the structure found by \citet{2004ApJ...612..308J,2006ApJ...648.1066J} for several X-ray binaries (including 4U 1636-536) and interpreted as absorption by neutral and ionized oxygen and neon in the interstellar medium."
+ We find no emissiou lines iu the RCS spectra that could explain the residuals iu the EPIC) PN spectra., We find no emission lines in the RGS spectra that could explain the residuals in the EPIC PN spectra.
+ Because we are mainly interested iu the spectrum at higher energies near the Fe Ίνα linc. we simply excluded the data near the absorption edees with the largest residuals.," Because we are mainly interested in the spectrum at higher energies near the Fe $\alpha$ line, we simply excluded the data near the absorption edges with the largest residuals."
+ The following energy rauges were excluded from the fits: <0.5 keV (O K-edee at 0.5[ keV). 0.78.0.99 keV. (Ne K-edge at 0:57 keV). 1.131.15 keV (Mg K-edge at 1.3 keV). aud 1.751.57 keV (Si Ik-edge at 1.5 keV).," The following energy ranges were excluded from the fits: $<0.5$ keV (O K-edge at 0.54 keV), 0.78–0.99 keV (Ne K-edge at 0.87 keV), 1.13--1.45 keV (Mg K-edge at 1.3 keV), and 1.75–1.87 keV (Si K-edge at 1.8 keV)."
+ Tn order to determüne au appropriate model for the continui enmüssion. we mitiallv fitted the spectra while excluding the 2310 keV energv rauge.," In order to determine an appropriate model for the continuum emission, we initially fitted the spectra while excluding the 3–10 keV energy range."
+ This excludes any coutribution from a relativistically broadened. fluorescent Fe I&-o linewhich. according to the model in NSPEC (?).. is limited to the 3.10 keV energv range for accretion disks around neutron stars.," This excludes any contribution from a relativistically broadened, fluorescent Fe $\alpha$ linewhich, according to the model in XSPEC \citep{1989MNRAS.238..729F}, is limited to the 3–10 keV energy range for accretion disks around neutron stars."
+ The X-rav spectra of LAINBs are typically characterized by a hard component interpreted as Comptouization of soft plotous by a hot corona. a soft coumpoucut thought to be blackbody raciatiou from the accretion disk. aud a broad Fe Ίνα line at 6.L keV (e.g.?2)..," The X-ray spectra of LMXBs are typically characterized by a hard component interpreted as Comptonization of soft photons by a hot corona, a soft component thought to be blackbody radiation from the accretion disk, and a broad Fe $\alpha$ line at 6.4 keV \citep[e.g.][]{2000ApJ...533..329B}."
+ To fit the N-xav προσ of the contiuuun clission we therefore combined a compoucut (2) for a Coptonizing corona or boundary laver. a disk blackbody coniponeut (c.c.2). aud a compoucut for photoclectric absorption with variable abundances (nodel 1).," To fit the X-ray spectrum of the continuum emission we therefore combined a component \citep{1994ApJ...434..570T} for a Comptonizing corona or boundary layer, a disk blackbody component \citep[e.g.][]{1986ApJ...308..635M}, and a component for photoelectric absorption with variable abundances (model 1)."
+ In the model we used the approximation for a disk ecolctry aud fixedthe redshift at 0., In the model we used the approximation for a disk geometry and fixed the redshift at 0.
+ Iu the model we ouly varied the elemental abundances for O. Ne. Si. and Fe.," In the model we only varied the elemental abundances for O, Ne, Si, and Fe."
+ Abunudauces for the other elements were not well constrained by the fit aud were fixed at their solar values according to 7.., Abundances for the other elements were not well constrained by the fit and were fixed at their solar values according to \citet{2000ApJ...542..914W}.
+ The spectra from the three observations are reasonably well fitted by model 1 with &Z values of 1.11. 1.50. and 1.28 (Table 1)).," The spectra from the three observations are reasonably well fitted by model 1 with $\chi_\nu^2$ values of 1.44, 1.80, and 1.28 (Table \ref{spectralmodels}) )."
+ We find that the ft cau be siguificautlv inproved by adding a second blackbody coiiponieut (bbodyrad)) with a temperature of ~200 eV (anodel 2)., We find that the fit can be significantly improved by adding a second blackbody component ) with a temperature of $\sim$ 200 eV (model 2).
+ The additional colmponent improves the 42 for the three observations to 1.36. 1.13. aud 0.97. respectively.," The additional component improves the $\chi_\nu^2$ for the three observations to 1.36, 1.43, and 0.97, respectively."
+ A second blackbody conrponeut with aio similar temperature was also introduced by ? in their fitting of BeppoSAN spectra, A second blackbody component with a similar temperature was also introduced by \citet{2006ApJ...651..416F} in their fitting of spectra
+"between Sirius A and B for optical and longer wavelengths, the most successful observations of Sirius B have been space-based.","between Sirius A and B for optical and longer wavelengths, the most successful observations of Sirius B have been space-based."
+" spectra of Sirius B have led to the accurate determination of its effective temperature (T.g—25,193 K) and mass (0.978Mo;?), parallax measurements of Sirius A place the system 2.64 pc from the Sun (?),, and astrometric monitoring has determined Sirius B's orbital period (50.090 yrs), semi-major axis (7.500"" or 19.8 AU), and eccentricity (0.5923;?)/ad.usno.navy.mil/wds/orb6."," spectra of Sirius B have led to the accurate determination of its effective temperature $T_{\rm eff}=$ 25,193 K) and mass \citep[0.978 $M_{\sun}$, parallax measurements of Sirius A place the system 2.64 pc from the Sun \citep{1997A&A...323L..49P}, and astrometric monitoring has determined Sirius B's orbital period (50.090 yrs), semi-major axis (7.500"" or 19.8 AU), and eccentricity \citep[0.5923;][]{2001AJ....122.3472H}."
+"html.. Increasingly, ground-based adaptive optics (AO) systems are becoming adept at contrast imaging, usually with the goal of discovering faint planets/brown dwarfs near their bright host stars."," Increasingly, ground-based adaptive optics (AO) systems are becoming adept at high-contrast imaging, usually with the goal of discovering faint planets/brown dwarfs near their bright host stars."
+" Imaging a faint white dwarf like Sirius B around a bright, main-sequence A star like Sirius A is a similar problem."," Imaging a faint white dwarf like Sirius B around a bright, main-sequence A star like Sirius A is a similar problem."
+" Recently, ? used the ESO 3.6 meter telescope, along with the ADONIS AO system, to image Sirius B in the JHKs filters, which, before this work, were the longest wavelength photometric measurements of Sirius B. The ? measurements showed a small (1.70) Ks-band excess when compared to the models of ?.."," Recently, \citet{2008A&A...489..651B} used the ESO 3.6 meter telescope, along with the ADONIS AO system, to image Sirius B in the JHKs filters, which, before this work, were the longest wavelength photometric measurements of Sirius B. The \citet{2008A&A...489..651B} measurements showed a small $\sigma$ ) Ks-band excess when compared to the models of \citet{2006AJ....132.1221H}."
+" Similar near-IR excesses are present in the spectra of dusty whitedwarfs, which are characterized by large excesses in the mid-infrared (?).."," Similar near-IR excesses are present in the spectra of dusty whitedwarfs, which are characterized by large excesses in the mid-infrared \citep{2007ApJ...662..544V}."
+" The first infrared excess around a white dwarf (G29-38) was found by ?,, and was initially thought to be a brown dwarf."," The first infrared excess around a white dwarf (G29-38) was found by \citet{1987Natur.330..138Z}, and was initially thought to be a brown dwarf."
+ This hypothesis was ruled out by optical/near-IR pulse monitoring that suggested a disk-geometry source of the excess (??)..," This hypothesis was ruled out by optical/near-IR pulse monitoring that suggested a disk-geometry source of the excess \citep{1990ApJ...357..216G,1991ApJ...374..330P}."
+ Subsequent spectroscopy using //IRS revealed the presence of small dust grains (?).., Subsequent spectroscopy using /IRS revealed the presence of small dust grains \citep{2005ApJ...635L.161R}.
+" A /TRS survey found 4 dusty white dwarfs out of their sample of 124 (??),, and that each of the dusty white dwarfs were of type DAZ (metal-polluted, hydrogen atmosphere), leading to coin DAZ stars with mid-infrared excesses, DAZd."," A /IRS survey found 4 dusty white dwarfs out of their sample of 124 \citep{2007ApJS..171..206M,2007ApJ...662..544V}, and that each of the dusty white dwarfs were of type DAZ (metal-polluted, hydrogen atmosphere), leading \citet{2007ApJ...662..544V} to coin DAZ stars with mid-infrared excesses, DAZd."
+" The metals in DAZ white dwarfs are expected to settle below the white dwarf photosphere much faster than evolutionary timescales (?),, thus the presence of metals implies a recent accretion event."," The metals in DAZ white dwarfs are expected to settle below the white dwarf photosphere much faster than evolutionary timescales \citep{2006A&A...453.1051K}, thus the presence of metals implies a recent accretion event."
+" ~1/4 of DA (hydrogen atmosphere) white dwarfs are DAZ (?),, and ~1/3 of DB (helium atmosphere) white dwarfs are DBZ (?),, demonstrating that accretion must be a common phenomenon for white dwarfs."," $\sim$ 1/4 of DA (hydrogen atmosphere) white dwarfs are DAZ \citep{1998ApJ...505L.143Z}, and $\sim$ 1/3 of DB (helium atmosphere) white dwarfs are DBZ \citep{2010ApJ...722..725Z}, demonstrating that accretion must be a common phenomenon for white dwarfs."
+" The cause of this accretion, in many if not all cases, is thought to be tidally disrupted asteroid-sized objects from a remanent debris disk/planetary system (??).."," The cause of this accretion, in many if not all cases, is thought to be tidally disrupted asteroid-sized objects from a remanent debris disk/planetary system \citep{2008AJ....135.1785J,2010ApJ...722..725Z}."
+" We note that Sirius B is a DA white dwarf, but determining if it is metal polluted (DAZ) is complicated by the difficulty of taking high-resolution spectra of Sirius B from the ground."," We note that Sirius B is a DA white dwarf, but determining if it is metal polluted (DAZ) is complicated by the difficulty of taking high-resolution spectra of Sirius B from the ground."
+" This problem is exacerbated by the fact that Sirius B’s high temperature, impedes the detection and interpretation of photospheric metals (??).."," This problem is exacerbated by the fact that Sirius B's high temperature, \citep[25,193 K][]{2005MNRAS.362.1134B} impedes the detection and interpretation of photospheric metals \citep{2006A&A...453.1051K,1995ApJ...454..429C}."
+ Sirius B's high temperature also has implications for its potential to host a debris disk., Sirius B's high temperature also has implications for its potential to host a debris disk.
+" While typical DAZd white dwarfs have temperatures ranging from T227,000-15,000 K (?),, Sirius B's high temperature would sublimate dust out to its tidal truncation radius (?),, precluding tidal"," While typical DAZd white dwarfs have temperatures ranging from $\approx$ 7,000-15,000 K \citep{2009ApJ...694..805F}, , Sirius B's high temperature would sublimate dust out to its tidal truncation radius \citep{2003ApJ...584L..91J}, , precluding tidal"
+Massive star colliding wind systems continue to be an area of intense study.,Massive star colliding wind systems continue to be an area of intense study.
+ The colliding wind interaction region (CWIR) is intrinsically interesting as a problem of hydrodynamic flow (Girard Willson 1987: Shore Brown 1988: Stevens. Blondin. Pollock 1992: Usov 1992: Canto. Raga. Wilkin 1996: Gayley. Owocki. Cranmer 1997).," The colliding wind interaction region (CWIR) is intrinsically interesting as a problem of hydrodynamic flow (Girard Willson 1987; Shore Brown 1988; Stevens, Blondin, Pollock 1992; Usov 1992; Canto, Raga, Wilkin 1996; Gayley, Owocki, Cranmer 1997)."
+ The CWIRs lead to a number of interesting observable effects as well. such as the X-ray emissions (e.g.. Stevens 11996: Maeda 11999: Zhekov Skinner 2000: Pittard 22002: Henley. Stevens. Pittard 2003). ultraviolet (UV). optical. and infrared (CIR) emission line profile effects (e.g. Stevens 1993: Lühhrs 1997: Stevens Howarth 1999: Hill 22000: Hill. Moffat. St-Louis 2002). dust emission (e.g. Tuthill. Monnier. Danchi 1999: Monnier. Tuthill. Danchi 1999: Tuthill 22006: Tuthill 22008). and polarization (e.g.. Brown. McLean. Emslie 1978: Drissen 11986: Villar-Sbaffi 22003. 2005. 2006)," The CWIRs lead to a number of interesting observable effects as well, such as the X-ray emissions (e.g., Stevens 1996; Maeda 1999; Zhekov Skinner 2000; Pittard 2002; Henley, Stevens, Pittard 2003), ultraviolet (UV), optical, and infrared (IR) emission line profile effects (e.g., Stevens 1993; Lühhrs 1997; Stevens Howarth 1999; Hill 2000; Hill, Moffat, St-Louis 2002), dust emission (e.g., Tuthill, Monnier, Danchi 1999; Monnier, Tuthill, Danchi 1999; Tuthill 2006; Tuthill 2008), and polarization (e.g., Brown, McLean, Emslie 1978; Drissen 1986; Villar-Sbaffi 2003, 2005, 2006)."
+ Of chief interest between relating the observables to the hydrodynamical predictions are. the determinations of the mass-loss rates for the two stars. their orbital properties. and the opening angle of the colliding wind bow shock.," Of chief interest between relating the observables to the hydrodynamical predictions are the determinations of the mass-loss rates for the two stars, their orbital properties, and the opening angle of the colliding wind bow shock."
+ Obtaining the orbital solution implies deriving the masses of the two stars. which is important for understanding massive star evolution.," Obtaining the orbital solution implies deriving the masses of the two stars, which is important for understanding massive star evolution."
+ As is typical with unresolved binary systems. the viewing inclination ¢ creates a challenge to obtaining the full orbital solution and stellar masses.," As is typical with unresolved binary systems, the viewing inclination $i$ creates a challenge to obtaining the full orbital solution and stellar masses."
+ In the study of CWIRs of massive binaries. the interpretation of optically thin lines remains prominent. such as X-ray lines that encode information about opening angle and viewing inclination.," In the study of CWIRs of massive binaries, the interpretation of optically thin lines remains prominent, such as X-ray lines that encode information about opening angle and viewing inclination."
+ One area that has been largely ignored is the influence of the CWIR on the of forbidden emission line profiles. especially those in the IR.," One area that has been largely ignored is the influence of the CWIR on the of forbidden emission line profiles, especially those in the IR."
+ TheObservatory CISO) and the Spitzer have provided a large quantity of high quality IR spectra for massive) stars., The ') and the ') have provided a large quantity of high quality IR spectra for massive stars.
+ Especially in the case of the evolved Wolf-Rayet (WR) stars. the IR spectra are rich in emission," Especially in the case of the evolved Wolf-Rayet (WR) stars, the IR spectra are rich in emission"
+passing from the South-West of he SAGS field through to the Nort]East.,passing from the South-West of the SA68 field through to the North-East.
+ Along this direction. a little «wer half a degree from the center of SAGS. lic the N-ray clusters CLOOITG6|16 aid citchug95 with redshift +=0.5.," Along this direction, a little over half a degree from the center of SA68, lie the X-ray clusters CL0016+16 and \\cite{hug95}
+ with redshift $z=0.54$."
+ Further. as part of the SITARC survey. (see Romer et in these proceedres) we have ideutified au additional N-raw cluster in the vicinity of citecon96 (RN JO0018.8|1602).," Further, as part of the SHARC survey (see Romer et in these proceedings) we have identified an additional X-ray cluster in the vicinity of \\cite{con96} (RX J0018.8+1602)."
+ A recent observation with the ARC 3.51 lescope las confirmed the redsift of this cluster as +=0.511 (Figure 2)., A recent observation with the ARC 3.5m telescope has confirmed the redshift of this cluster as $z=0.541$ (Figure 2).
+The radiation. characterised. by the temperature. of. the super-heated [aver Zi is produced only in the super-heated Iaver.,The radiation characterised by the temperature of the super-heated layer $T_{\rm s}$ is produced only in the super-heated layer.
+ In other lavers. this radiation is not. produced. but. is absorbed or scattered.," In other layers, this radiation is not produced but is absorbed or scattered."
+ The super-heated. laver is vertically. optically thin for its own radiation as zi—ας=(WeONLPea«n1 because the grazing. angle à is+ small ancl we have (877ONL£8)exlx1.," The super-heated layer is vertically optically thin for its own radiation as $\tau_{\rm s,s}^{\rm ext}\equiv\kappa_{\rm s}^{\rm ext}\Sigma_{\rm s}
+=(\kappa_{\rm s}^{\rm ext}/\kappa_*^{\rm ext})\alpha\ll1$ because the grazing angle $\alpha$ is small and we have $(\kappa_{\rm s}^{\rm ext}/\kappa_*^{\rm ext})\le1$."
+ On the other hand. the total optical. depth of other lavers is very large.," On the other hand, the total optical depth of other layers is very large."
+ Thus. we consider a ecometry that a thin isothermal laver lies on a semi-infinite absorption and scattering slab.," Thus, we consider a geometry that a thin isothermal layer lies on a semi-infinite absorption and scattering slab."
+ The above of the super-heated laver is assumed to be vacuum: there is no downwards input radiation with the temperature Z2 at the top of the TIaver., The above of the super-heated layer is assumed to be vacuum; there is no downwards input radiation with the temperature $T_{\rm s}$ at the top of the layer.
+ However. there is upwards input radiation at the bottom of the laver because of the reflection by the semi-infinite interior below the laver.," However, there is upwards input radiation at the bottom of the layer because of the reflection by the semi-infinite interior below the layer."
+" The upwards and downwarels radiation intensities from the super-heated laver (2 and phe* respectively) in the two-stream Ecldington approximation (the cosine of the angle between the stream lines and the normal of the laver is set to be 41/43) become from equations (B12) and (B16) and where £L"" js the upwards input intensity from the interior. and e. ος. and ὃς are the thermal. rellection. and transmission coefficients in the super-heated [aver [or the radiation. characterised. by the temperature. Z5. (see eqs. "," The upwards and downwards radiation intensities from the super-heated layer $I^{\rm up}_{\rm s,s}$ and $I^{\rm down}_{\rm s,s}$, respectively) in the two-stream Eddington approximation (the cosine of the angle between the stream lines and the normal of the layer is set to be $\pm1/\sqrt{3}$ ) become from equations (B12) and (B16) and where $I^{\rm up}_{\rm i,s}$ is the upwards input intensity from the interior, and $a_{\rm s,s}$, $b_{\rm s,s}$, and $c_{\rm s,s}$ are the thermal, reflection, and transmission coefficients in the super-heated layer for the radiation characterised by the temperature $T_{\rm s}$ (see eqs. ["
+B13BI5]).,B13–B15]).
+" The rellected. intensity [rom the interior becomes where 5i, is the reflection coellicient in the interior for the radiation with Zi."," The reflected intensity from the interior becomes where $b_{\rm i,s}$ is the reflection coefficient in the interior for the radiation with $T_{\rm s}$."
+ Note that the interior itself does not emit radiation with 7i., Note that the interior itself does not emit radiation with $T_{\rm s}$.
+ We can always obtain the unique exact solution of £7[m eet and IU [rom equations (S10). whereas we derive an approximate solution of them in this paper.," We can always obtain the unique exact solution of $I^{\rm up}_{\rm s,s}$, $I^{\rm down}_{\rm s,s}$, and $I^{\rm up}_{\rm i,s}$ from equations (8–10), whereas we derive an approximate solution of them in this paper."
+ As found equation (DIS). for a semi-infinte. slab. we have bj.m(νοας| NL) where 4;=Yla. with a. being the scattering albedo for the radiation with 72.," As found equation (B18), for a semi-infinte slab, we have $b_{\rm i,s}\approx(1-\chi_{\rm s})/(1+\chi_{\rm s})$ , where $\chi_{\rm s}=\sqrt{1-\omega_{\rm s}}$ with $\omega_{\rm s}$ being the scattering albedo for the radiation with $T_{\rm s}$."
+ Since the super- laver is optically thin for its own radiation. we can approximate b...] and c.z1.," Since the super-heated layer is optically thin for its own radiation, we can approximate $b_{\rm s,s}\ll1$ and $c_{\rm s,s}\approx1$."
+ We also find a...zο...=πο...Ke from equation (1919) for a small optical depth.," We also find $a_{\rm s,s}\approx\sqrt{3}\chi_{\rm s}^2\tau_{\rm s,s}^{\rm ext}
+=\sqrt{3}\chi_{\rm s}^2(\kappa_{\rm s}^{\rm ext}/\kappa_*^{\rm ext})\alpha$ from equation (B13) for a small optical depth."
+" Then. we obtain £77=(11ban... LO""doscmalgdA. and £0zmbyaBL. op"," Then, we obtain $I^{\rm up}_{\rm s,s}\approx (1+b_{\rm i,s})a_{\rm s,s}B_{\rm s}$, $I^{\rm down}_{\rm s,s}\approx a_{\rm s,s}B_{\rm s}$, and $I^{\rm up}_{\rm i,s}\approx b_{\rm i,s}a_{\rm s,s}B_{\rm s}$."
+sThe upwards and downwards Iluxes from the super-heatec laver are and 4p99—(qvoaunPyia273).," The upwards and downwards fluxes from the super-heated layer are and $H^{\rm down}_{\rm s}=(I^{\rm down}_{\rm s,s}-I^{\rm up}_{\rm i,s})/(2\sqrt{3})$."
+" ""Pherefore. we obtain and As discussed in the section 3.1. we consider the micelle laver only when the laver is optically thin for its own radiation."," Therefore, we obtain and As discussed in the section 3.1, we consider the middle layer only when the layer is optically thin for its own radiation."
+ Although the middle laver is sandwiched between the super-heated laver and the interior. we neglect the ellect of the super-heated laver because the optical depth of the laver is very small (see above).," Although the middle layer is sandwiched between the super-heated layer and the interior, we neglect the effect of the super-heated layer because the optical depth of the layer is very small (see above)."
+ The optical depth of the interior is so large that we can regard it as a semi-infinite medium., The optical depth of the interior is so large that we can regard it as a semi-infinite medium.
+ Thus. we have the same setting as the super-heated laver. other than the optical thickness of the laver. zu=Gunp).," Thus, we have the same setting as the super-heated layer, other than the optical thickness of the layer, $\tau_{\rm m,m}^{\rm ext}
+=(\kappa_{\rm m}^{\rm ext}/\kappa_{\rm s}^{\rm ext})$."
+" Following the section 3.1.2. we have and where D is the integrated. Planck function with the temperature of the middle laver Zi. au, is the scattering albedo for the radiation of the middle laver. and xiV/l Vau "," Following the section 3.1.2, we have and where $B_{\rm m}$ is the integrated Planck function with the temperature of the middle layer $T_{\rm m}$, $\omega_{\rm m}$ is the scattering albedo for the radiation of the middle layer, and $\chi_{\rm m}=\sqrt{1-\omega_{\rm m}}$ ."
+We always consider the interior to be optically thick. for its own radiation., We always consider the interior to be optically thick for its own radiation.
+ On the other hand. other lavers above the interior are. considered. to be always. optically thin.," On the other hand, other layers above the interior are considered to be always optically thin."
+ Η we neglect. the clfect of the upper layers. the interior is regarded: as a semi-infinite isothermal medium without incident Hux of the radiation with the temperature Zi.," If we neglect the effect of the upper layers, the interior is regarded as a semi-infinite isothermal medium without incident flux of the radiation with the temperature $T_{\rm i}$."
+" In this case. based on equation (B12). the upwards outbound flux becomes LE""?=aiiBi(24/3) with ei; being the thermal coellicient in the interior for the radiation with Zi ancl D being the integrated. Planck function with 7;."," In this case, based on equation (B12), the upwards outbound flux becomes $H_{\rm i}^{\rm up} = a_{\rm i,i} B_{\rm i} / (2\sqrt{3})$ with $a_{\rm i,i}$ being the thermal coefficient in the interior for the radiation with $T_{\rm i}$ and $B_{\rm i}$ being the integrated Planck function with $T_{\rm i}$."
+" For à semi-infinite medium. ej;72xyi1/(E|xi). where yi=Vlai with a, being the single scattering albedo for the interior radiation (eq. "," For a semi-infinite medium, $a_{\rm i,i}\approx2\chi_{\rm i}/(1+\chi_{\rm i})$, where $\chi_{\rm i}=\sqrt{1-\omega_{\rm i}}$ with $\omega_{\rm i}$ being the single scattering albedo for the interior radiation (eq. ["
+D17]).,B17]).
+ Therefore. we have Note that the factor in 20] ds equal or tess than 1/2.," Therefore, we have Note that the factor in $[~~~]$ is equal or less than 1/2."
+ Vhe equal1 is true without scattering., The equal is true without scattering.
+8 Pherefore. in eeneral. scattering reduces the radiation energy loss from the interior. i.e. the ereen-house effect.," Therefore, in general, scattering reduces the radiation energy loss from the interior, i.e. the green-house effect."
+ The radiation energy conservation in the super-heated Laver is, The radiation energy conservation in the super-heated layer is
+[function for region candidates with a power-law index at 6 cm of z—1.2. but exclude sources with spectral indices of a>0.2.,"function for region candidates with a power-law index at 6 cm of $\approx -1.2$, but exclude sources with spectral indices of $\alpha > 0.2$."
+ As result. the most optically thick sources were excluded from (heir sample.," As result, the most optically thick sources were excluded from their sample."
+ Never the less. (he seven sources presented in with a>0 are also included in (his paper (our sources 1. 2. 4. 6. 7. 8. and 10 in Tables 3. and 6)).," Never the less, the seven sources presented in \citet{hyman00} with $\alpha > 0$ are also included in this paper (our sources 1, 2, 4, 6, 7, 8, and 10 in Tables \ref{tbl-3}
+ and \ref{N6946.tab}) )."
+" Ou primary (ool for distinguishing UUs from other (vpes of radio sources. such as supernovae remnants (SNR). is the radio spectral index a (where 5,x»)."," Our primary tool for distinguishing s from other types of radio sources, such as supernovae remnants (SNR), is the radio spectral index $\alpha$ (where $S_{\nu} \propto
+\nu^{\alpha}$ )."
+ While optically (hin sources have a=—0.1. optically thick sources have a2—0.1 (6 approaches 2 in the black body limit). and tvpical non-thermal objects have a<—0.4.," While optically thin sources have $\alpha = -0.1$, optically thick sources have $\alpha > -0.1$ $\alpha$ approaches 2 in the black body limit), and typical non-thermal objects have $\alpha <
+-0.4$."
+" Unfortunately. not all SNRs obev (his convention in particular. there is a class of “composite” or ""plerionic"" SNRs (ascoinedbyWeiler&Panagia1978). that are known to have atvpically flat spectral indices (70.3<a< 0.0)."," Unfortunately, not all SNRs obey this convention — in particular, there is a class of “composite” or “plerionic” SNRs \citep[as coined by][]{weiler78} that are known to have atypically flat spectral indices $-0.3 < \alpha < 0.0$ )."
+ In fact. in the first few hundred days after a supernova explosion. a can actually be positive (Weileretal.1986).," In fact, in the first few hundred days after a supernova explosion, $\alpha$ can actually be positive \citep{weiler86}."
+. ILowever. given this extremely short timescale. we would have to be extremely fortunate (or unfortunate) to have one of these objects contained in our sample.," However, given this extremely short timescale, we would have to be extremely fortunate (or unfortunate) to have one of these objects contained in our sample."
+ A second set of observations taken with adecuate time separation. such as (those emploved by Nobulnicky&Johnson(1999) would resolve (his issue with confidence.," A second set of observations taken with adequate time separation, such as those employed by \citet{kj99} would resolve this issue with confidence."
+ Consequently. in order to select candidates from (he radio data. we have applied the condition (hat these sources must have radio spectral indices a70 (higher [luxes ab shorter wavelengths).," Consequently, in order to select candidates from the radio data, we have applied the condition that these sources must have radio spectral indices $\alpha > 0$ (higher fluxes at shorter wavelengths)."
+ This eriterion should exclude most other classes of raclio-emil(ing objects. but it max also exclude genuine regions depending upon the exact instrumental angular resolution and [frequencies observed.," This criterion should exclude most other classes of radio-emitting objects, but it may also exclude genuine regions depending upon the exact instrumental angular resolution and frequencies observed."
+ For example. low-lrequency observations of oplically thick thermal sources could be contaminated by steep-spectrum non-thermal sources. such as SNRs. and (hus fail the a>0 test. especially at longer wavelengths.," For example, low-frequency observations of optically thick thermal sources could be contaminated by steep-spectrum non-thermal sources, such as SNRs, and thus fail the $\alpha > 0$ test, especially at longer wavelengths."
+" Furthermore. more clistant sources. or sources observed using beam sizes larger (han the (thermal emitting region, mav be contaminated bv steep-specirum. non-thermal background. emission (hat dominates the disks ancl halos of most galaxies."," Furthermore, more distant sources, or sources observed using beam sizes larger than the thermal emitting region, may be contaminated by steep-spectrum, non-thermal background emission that dominates the disks and halos of most galaxies."
+ Both of these effects will act to disguise an oplically thick thermal bremsstralilung signature., Both of these effects will act to disguise an optically thick thermal bremsstrahlung signature.
+ Thermal sources with the lowest emission measure. and thus the lowest Iree-DIree optical depth will be affected most strongly.," Thermal sources with the lowest emission measure, and thus the lowest free-free optical depth will be affected most strongly."
+ Sources with the highest emission measures are more likely to maintain a positive spectral signature and be correctly. identified., Sources with the highest emission measures are more likely to maintain a positive spectral signature and be correctly identified.
+ Fortunately. it is precisely these sources with high emission measures that are most likely to host ultra-voung stellar clusters.," Fortunately, it is precisely these sources with high emission measures that are most likely to host ultra-young stellar clusters."
+projection of the matrix of observations onto the chosen basis of eigenvectors: The selection of the number of eigenvectors that are dominated by polarimetric signal is the fundamental free parameter of the denoising procedure.,projection of the matrix of observations onto the chosen basis of eigenvectors: The selection of the number of eigenvectors that are dominated by polarimetric signal is the fundamental free parameter of the denoising procedure.
+ It is not an easy task to efficiently select it and sometimes requires subjective criteria., It is not an easy task to efficiently select it and sometimes requires subjective criteria.
+ We have verified that the following criterium works quite well for many of the tested cases and ts also based on the properties of the PCA decomposition., We have verified that the following criterium works quite well for many of the tested cases and is also based on the properties of the PCA decomposition.
+ For each value of the variance of the noise added to the simulated profiles. we compute an observation matrix equal to the observed one but made only of uncorrelated noise.," For each value of the variance of the noise added to the simulated profiles, we compute an observation matrix equal to the observed one but made only of uncorrelated noise."
+ In real situations. this matrix has to be built based on the estimation of the noise present in the observations.," In real situations, this matrix has to be built based on the estimation of the noise present in the observations."
+ The eigenvalues of the pure noise matrix are calculated and compared with those of the real observations., The eigenvalues of the pure noise matrix are calculated and compared with those of the real observations.
+ If the variance of the noise has been correctly estimated. the two eigenvalue distributions will overlap except for those eigenvalues associated with correlated signals.," If the variance of the noise has been correctly estimated, the two eigenvalue distributions will overlap except for those eigenvalues associated with correlated signals."
+ Consequently. we select those eigenvectors 5; with eigenvalues higher than a factor f of those corresponding to the pure noise case.," Consequently, we select those eigenvectors $b_i$ with eigenvalues higher than a factor $f$ of those corresponding to the pure noise case."
+ We have verified that £=1.2 gives good results., We have verified that $f \approx 1.2$ gives good results.
+ It is also instructive not to rely on automatic selection methods but verify the shape and weight of each eigenvector., It is also instructive not to rely on automatic selection methods but verify the shape and weight of each eigenvector.
+ The direct analysis of the eigenvectors can show many important details about the hidden signal and some tricks can be used to enhance the possibility of recovering it., The direct analysis of the eigenvectors can show many important details about the hidden signal and some tricks can be used to enhance the possibility of recovering it.
+ A detailed discussion of the properties of the eigenvalues and the choice of the cutoff for the SVD problem can be found in Christensen-Dalsgaardet=al.(1993) and in a series of papers by Hansen(1992). Hansenetal.(1992) and Hansenetal.(1993).," A detailed discussion of the properties of the eigenvalues and the choice of the cutoff for the SVD problem can be found in \cite{christensen93} and in a series of papers by \cite{hansen92}, \cite{hansen_sekii92} and \cite{hansen93}."
+. It is important to point out that the results we are presenting here are not the optimal case., It is important to point out that the results we are presenting here are not the optimal case.
+ For a field of 1000 G and since we do not include any additional line broadening mechanism. the majority of the lines are not in the weak field regime of the Zeeman effect.," For a field of 1000 G and since we do not include any additional line broadening mechanism, the majority of the lines are not in the weak field regime of the Zeeman effect."
+ Therefore. they present different. shapes (as a consequence of the Zeeman saturation) and part of the correlation is lost.," Therefore, they present different shapes (as a consequence of the Zeeman saturation) and part of the correlation is lost."
+ We show that the PCA denoising technique is very performant in this non-optimal case., We show that the PCA denoising technique is very performant in this non-optimal case.
+ In stars with broadened lines (due to whatever mechanism). the weak field regime can be expected for larger magnetic field strengths and our PCA denoising algorithm will work even better.," In stars with broadened lines (due to whatever mechanism), the weak field regime can be expected for larger magnetic field strengths and our PCA denoising algorithm will work even better."
+high as LO°GeV. a [actor 5-10 larger than one could find without including the magnetic backreaction.,"high as $10^6\rmn{GeV}$, a factor 5-10 larger than one could find without including the magnetic backreaction."
+ In. principle even larger maximum energies can be achieved by taking into account the non-resonant streaming instability in the early phases of the SNR evolution., In principle even larger maximum energies can be achieved by taking into account the non-resonant streaming instability in the early phases of the SNR evolution.
+ Our entire analysis was carried out [or the case of a shock propagating5 in the direction parallel to the ambient magnetic field., Our entire analysis was carried out for the case of a shock propagating in the direction parallel to the ambient magnetic field.
+ The amplification of the background field due to cosniüc rav streamüng soon changes (his situation. in that the magnetic field upstream becomes highlv turbulent.," The amplification of the background field due to cosmic ray streaming soon changes this situation, in that the magnetic field upstream becomes highly turbulent."
+ Changing the field obliquity should not change much the picture discussed in the present paper as long as the projection of the shock velocity along the fieldl lines remains supersonic., Changing the field obliquity should not change much the picture discussed in the present paper as long as the projection of the shock velocity along the field lines remains supersonic.
+ If the supernova shock propagates in a medium in which the background field is mainly oriented in a given direction. there will be locations ad which the shock is quasi-parallel ancl others al which it is euasi-perpendicular.," If the supernova shock propagates in a medium in which the background field is mainly oriented in a given direction, there will be locations at which the shock is quasi-parallel and others at which it is quasi-perpendicular."
+ The implications for (he morphology of the non-thermal emission require a dedicated investigation. in that the quasi-perpendicular regions could be expected to be bright because of the more efficient perpendicular configuration. but quasi-parallel regions could be expected to be bright because of more efficient magnetic field amplification.," The implications for the morphology of the non-thermal emission require a dedicated investigation, in that the quasi-perpendicular regions could be expected to be bright because of the more efficient perpendicular configuration, but quasi-parallel regions could be expected to be bright because of more efficient magnetic field amplification."
+ Clearly the degree to which these expectations mav be true depends on many details which at present are not under control. although observations of specilic astrophysical objects (e.g. SN 1006) will provide us with important inlormation on this issue.," Clearly the degree to which these expectations may be true depends on many details which at present are not under control, although observations of specific astrophysical objects (e.g. SN 1006) will provide us with important information on this issue."
+ Finally. we investigated the effect of turbulent. heating on the shock modification.," Finally, we investigated the effect of turbulent heating on the shock modification."
+ We generalized the formalism introduced by Volk&MelIxenzie(1981) and (1982) in order (ο account lor the possibility (hat a finite fraction of the wave energy is damped on the background plasma., We generalized the formalism introduced by \cite{v-mck81} and \cite{mck-v82} in order to account for the possibility that a finite fraction of the wave energy is damped on the background plasma.
+ We find that. although in general the heating of the upstream plasma results in a decrease of the compressibilitv. the effect is al most competitive. if not subdominant. compared with the dynamical feedback of the amplified magnetic field on the plasma.," We find that, although in general the heating of the upstream plasma results in a decrease of the compressibility, the effect is at most competitive, if not subdominant, compared with the dynamical feedback of the amplified magnetic field on the plasma."
+ When turbulent heating is indeed dominant. a fraction very close to unity of the wave energy is transformed in thermal energy. and (he main effect is (hat of suppressing the erowth of the magnetic field. so that it becomes challenging to explain X-ray observations in terms of severe svnchrotron losses.," When turbulent heating is indeed dominant, a fraction very close to unity of the wave energy is transformed in thermal energy, and the main effect is that of suppressing the growth of the magnetic field, so that it becomes challenging to explain X-ray observations in terms of severe synchrotron losses."
+ Moreover in this case the maximum energy of accelerated particles is drastically reduced aud it becomes much smaller (han the knee., Moreover in this case the maximum energy of accelerated particles is drastically reduced and it becomes much smaller than the knee.
+ ]t is worth stressing that while the magnetic feedback on the background plasma is based on well known physics. the turbulent heating is (treated al best in a very phenomenological wav. Wilh little attention to the specific physical processes that may be responsible of the non-acliabatic heating and on how modes with different wavelengths are damped.," It is worth stressing that while the magnetic feedback on the background plasma is based on well known physics, the turbulent heating is treated at best in a very phenomenological way, with little attention to the specific physical processes that may be responsible of the non-adiabatic heating and on how modes with different wavelengths are damped."
+ In these, In these
+as an edge was apparent in the data residuals at 007 keV. The derived line strength is rather high (EW ~900 eV): the line is also unusually broad (σ71.4 keV) and redshifted (I 5.6 keV).,as an edge was apparent in the data residuals at $>7$ keV. The derived line strength is rather high (EW $\sim900$ eV); the line is also unusually broad $\sigma\sim1.4$ keV) and redshifted (E $\sim5.6$ keV).
+ We next tried a lit to highly ionised iron (model 4). the (narrow) line energy was fixed at 6.7 keV. (corresponding to Le-like Fe) and the edge energy was ο as a free parameter.," We next tried a fit to highly ionised iron (model 4), the (narrow) line energy was fixed at 6.7 keV (corresponding to He-like Fe) and the edge energy was left as a free parameter."
+ In this case an acceptable fit was found to the data MLsx 1}: with a deep and hiehly ionisecl iron Woedge at 8.7 keV with τ=0.7640.12., In this case an acceptable fit was found to the data $\chi_{\nu}^{2}\approx1$ ); with a deep and highly ionised iron K edge at 8.7 keV with $\tau=0.76\pm0.12$.
+ A iron emission line was no longer required in the fit (EN«77 eV)., A iron emission line was no longer required in the fit $<77$ eV).
+ The energy of the edge is consistent with iron of ionisation from Fe Fexxv., The energy of the edge is consistent with iron of ionisation from Fe -- Fe.
+.. Figure 3 clearly illustrates the high ionisation state of the reprocessing material., Figure 3 clearly illustrates the high ionisation state of the reprocessing material.
+ Note that we also tried fitting a smeared edge model (width 1 keV). as might be expected from the inner disk. although the fit obtained was very similar to above.," Note that we also tried fitting a smeared edge model (width 1 keV), as might be expected from the inner disk, although the fit obtained was very similar to above."
+ Including abread iron line in the model (in addition to the edge. model 5) lead to a further improvement in the spectral fit (N47~ L1). at a significance evel of ~99% confidence (E-test. 2 additional parameters).," Including a iron line in the model (in addition to the edge, model 5) lead to a further improvement in the spectral fit $\Delta\chi^{2}\sim11$ ), at a significance level of $\sim99\%$ confidence (F-test, 2 additional parameters)."
+ The line equivalent width is ~350 eV: the [ine is broad (o1 keV) and possibly redshilted (I2 ~6 keV)., The line equivalent width is $\sim350$ eV; the line is broad $\sigma\sim1$ keV) and possibly redshifted (E $\sim6$ keV).
+ We note hat the detection of the ionised iron edge is robust. even after the inclusion of a broad line in the model fit.," We note that the detection of the ionised iron edge is robust, even after the inclusion of a broad line in the model fit."
+ Lavine established. the presence of reprocessing by vighly ionised. matter. we then attempted. to model these reprocessing features in terms of reflection of N-ravs olf the surface lavers of an ionised accretion disk.," Having established the presence of reprocessing by highly ionised matter, we then attempted to model these reprocessing features in terms of reflection of X-rays off the surface layers of an ionised accretion disk."
+ Phe mocel in NSPEC was used (Magdziarz Zedziarski 1995). assuming a disk surface temperature of 105 Ix (e.g. Zycki 1994: ltoss. Fabian Young 1999) and a disk inclination angle of30.," The model in XSPEC was used (Magdziarz Zdziarski 1995), assuming a disk surface temperature of $^{6}$ K (e.g. Zycki 1994; Ross, Fabian Young 1999) and a disk inclination angle of."
+. The disk ionisation parameter (£) and the strength of the reflection component It (—O/2z) were both left as free parameters in the fit. (, The disk ionisation parameter $\xi$ ) and the strength of the reflection component R $\Omega/2\pi$ ) were both left as free parameters in the fit. (
+"Note that £=L ni?, where n is the number density of the material at a clistance r from an ionising source of luminosity £L. where £ is defined fron 5 eV to 20 keV.) An ionised. rellector. provided a good Lit to the hard X-ray. spectrum and the iron Ix edge (mocdel 6. v»7 1). giving a value of It close to 1 and a high ionisation (£=6400 1).","Note that $\xi=L/nr^{2}$ , where $n$ is the number density of the material at a distance $r$ from an ionising source of luminosity $L$, where $L$ is defined from 5 eV to 20 keV.) An ionised reflector provided a good fit to the hard X-ray spectrum and the iron K edge (model 6, $\chi_{\nu}^{2}\approx1$ ), giving a value of R close to 1 and a high ionisation $\xi=6400$ $^{-1}$ )."
+ Although a narrow iron Ix line is not required in this fit. the presence of a broad line was not ruled out (EW —250 eV).," Although a narrow iron K line is not required in this fit, the presence of a broad line was not ruled out (EW $\la250$ eV)."
+ We also attempt to model the hard. X-ray spectrum of PDS 456 with a power-law mocified by a warm line-of-sight absorber (model 7). using the routine in NSPEC (note we assume a temperature of 1=3 Ix.," We also attempt to model the hard X-ray spectrum of PDS 456 with a power-law modified by a warm line-of-sight absorber (model 7), using the routine in XSPEC (note we assume a temperature of $T=3\times10^{5}$ K)."
+ This model provides a good alternative description of the X-ray data. modelling both the ionised edge ancl providing a good overall fit (422 1).," This model provides a good alternative description of the X-ray data, modelling both the ionised edge and providing a good overall fit $(\chi_{\nu}^{2}\approx1$ )."
+ However both a large column of material CNgz4.11075 em. 7) and a very high degree of ionisation (C=1.5lot 1) are required to reproduce the energy and depth of the edge feature., However both a large column of material $N_{H}\approx4.7\times10^{23}$ $^{-2}$ ) and a very high degree of ionisation $\xi=1.5\times10^{4}$ $^{-1}$ ) are required to reproduce the energy and depth of the edge feature.
+ Extrapolation of the preferred. hard X-ray spectral model to lower energies (0.7. keV) does not. provide an aclequate description of the broad-band data., Extrapolation of the preferred hard X-ray spectral model to lower energies (0.7 keV) does not provide an adequate description of the broad-band data.
+ Phere remains a significant deficit of counts present below ~2 keV. (see figure 4. panel 1). indicating that significant intrinsic absorption (in excess of Galactic) is present.," There remains a significant deficit of counts present below $\sim$ 2 keV (see figure 4, panel 1), indicating that significant intrinsic absorption (in excess of Galactic) is present."
+ A neutral absorber (Ny~l077em 7) gave a poor fit: the spectrum is poorly modelled and an excess of counts is present. below lkeV in the delata., A neutral absorber $N_{\rm H}\sim10^{22}\rm{cm}^{-2}$ ) gave a poor fit; the spectrum is poorly modelled and an excess of counts is present below 1 keV in the data.
+ However anzZonised absorber provides à goocl Lit (figure 4. panel 2). with Ny~5+1.5]1077em7 and €—350+100 (again assuming —3.107 IX).," However an absorber provides a good fit (figure 4, panel 2), with $N_{\rm H}\sim5\pm1.5\times10^{22}{\rm cm}^{-2}$ and $\xi\sim350\pm100$ (again assuming $=3\times10^{5}$ K)."
+ Note that this material is not sullicientIv ionised to reproduce the deep. ionised. iron Ix. edge seen in the hard. X-ray spectrum: the absorber effectively. contributes a negligible amount 0.1) to the iron Ix edge.," Note that this material is not sufficiently ionised to reproduce the deep, ionised, iron K edge seen in the hard X-ray spectrum; the absorber effectively contributes a negligible amount $\tau<0.1$ ) to the iron K edge."
+ Phe absence of anv appreciable optical absorption (aside from Galactic) may also. imply either a low intrinsic dust-to-gas ratio. or that most of the dust close to the central engine hassublimated.," The absence of any appreciable optical absorption (aside from Galactic) may also imply either a low intrinsic dust-to-gas ratio, or that most of the dust close to the central engine hassublimated."
+ Light curves were constructed by binning the data into 96 min orbital bins. for the backerounc-subtracted ((210 keV) and," Light curves were constructed by binning the data into 96 min orbital bins, for the background-subtracted (2–10 keV) and"
+"is got equal to fy...p; Where Tis the unnber of zoues that have crossed 7;,. aud the new fractional increment à ds calculated from Eq. |A1]|","is set equal to $\mu_{j_{min}+l}$, where $l$ is the number of zones that have crossed $r_{in}$, and the new fractional increment $\alpha$ is calculated from Eq. \ref{finite:mu}] ]"
+ with the new value of inner boundary., with the new value of inner boundary.
+ Afterwards. all the variables are interpolated on the new eric.," Afterwards, all the variables are interpolated on the new grid."
+ As for the spatial ceutering of the variables. B=og aud wyaxe evaluated at iid zone (ityye). while wy at zone boundary (yj).," As for the spatial centering of the variables, $B=a_R T^4$ and $w_0$are evaluated at mid zone $\mu_{j-1/2}$ ), while $w_1$ at zone boundary $\mu_j$ )."
+ Coucerumg the tine centering. all quantities are evaluated at tho full time level £.," Concerning the time centering, all quantities are evaluated at the full time level $t^n$."
+ With this centeriug and following a staudard Lagrangian approach for the discretization of the equations aud the derivatives. the finite difference form of Eqs. |1|]. [2]]," With this centering and following a standard Lagrangian approach for the discretization of the equations and the derivatives, the finite difference form of Eqs. \ref{eq:energy}] ], \ref{eq:w0}] ]"
+ aud [3]] cau be written as: nil As €. p. kg and kp are in general rather couples. non-linear functions of p aud B (see. c.g. Eqs. |," and \ref{eq:w1}] ] can be written as: where As $\epsilon$, $p$, $k_R$ and $k_P$ are in general rather complex, non-linear functions of $\rho$ and $B$ (see, e.g., Eqs. ["
+"L1] aud {15} in Zauipicrictal. 1998)). relations |À2]|-A form a highly non-linear svsteii of (3/,,4,-3) equations having 3,443) uukuowns. that we solve by mucans of the lJ]Newtou-Raphsou iterative method (sec.c.g..Pressetal.1996.for details). ","14] and [15] in \citealt{zampieri98}) ), relations \ref{finite:energy}] \ref{finite:w1}] ] form a highly non-linear system of $j_{max}$ -3) equations having $j_{max}$ -3) unknowns, that we solve by means of the Newton-Raphson iterative method \citep[see, e.g.,][for details]{NumRec}. ."
+Using this approach. the finite difference equations are linearized aud solved iteratively for the corrections at the new," Using this approach, the finite difference equations are linearized and solved iteratively for the corrections at the new"
+simple average at the levels of ( = 1. 2. 3. aud f.,"simple average at the levels of $ \ell $ = 1, 2, 3, and 4."
+" Π is factor 10 larger than the siniple ave""age al í5."," It is factor 10 larger than the simple average at $ \ell \, = \, 5 $."
+ These result coulirm that our scheme gives the secoud order accurate gravity iu tlle cels ueitlier adjacent to the grid level bouudary nor lying ou the sharp edge of the density distributio, These result confirm that our scheme gives the second order accurate gravity in the cells neither adjacent to the grid level boundary nor lying on the sharp edge of the density distribution.
+ Iu the cells adjacert to the grid level bouudary. the relative error decreases 1 proportion Nt except at í1.," In the cells adjacent to the grid level boundary, the relative error decreases in proportion to $ N ^{-1} $ except at $ \ell \, = \, 1 $."
+ This is because our scheme is ouly the first order accurate at the erid lev boundary., This is because our scheme is only the first order accurate at the grid level boundary.
+ This does no mean that the error is always large ouly at the grid level )ouucdary., This does not mean that the error is always large only at the grid level boundary.
+ No that the average relative error is iore or less the same between in the cells adja‘ent to the eric level boundary aud iutje rest cells when Ν΄lA32.," Note that the average relative error is more or less the same between in the cells adjacent to the grid level boundary and inthe rest cells when $ N \, \le \, 32 $."
+" Only when N is sulliciently la‘ee, the error is appreciably large at the eric level bouudary."," Only when $ N $ is sufficiently large, the error is appreciably large at the grid level boundary."
+ At {1. the relative error decreases rapidly with increase in WV.," At $ \ell \, = \, 1 $, the relative error decreases rapidly with increase in $ N $."
+ When ;N.>|L28. the binary is resolved eve Lin the grid of (=1.," When $ N \, \ge \, 128 $, the binary is resolved even in the grid of $ \ell \, = 1 $."
+ Then the dipole and quadruple moineus of the gravity are evaluated with a good accuracy aud the outer bouudary couclitiou is improves| greatly., Then the dipole and quadruple moments of the gravity are evaluated with a good accuracy and the outer boundary condition is improved greatly.
+ Accordingly he relative er‘or is sinall at {tL when N>12s.," Accordingly the relative error is small at $ \ell \, = \, 1 $ when $ N \, \ge \, 128 $."
+ Remember that the e‘ror due to the outer )ouncdary COMMition dominates at {1.," Remember that the error due to the outer boundary condition dominates at $ \ell \, = \, 1 $."
+ We mace another exaliple to confirm that our scheme can reproduce no only the 1/r gravity iut also the qtaclraple moment of the gravity accurately., We made another example to confirm that our scheme can reproduce not only the $ 1/r $ gravity but also the quadraple moment of the gravity accurately.
+ Figure 7 slows tie second. example in which 4 uuifor1u density spheres are jxaced. near the center of the grid., Figure \ref{Q-pole.eps} shows the second example in which 4 uniform density spheres are placed near the center of the grid.
+ Tje centers. radii. aud passes of the spheres are listed in Tale 2.," The centers, radii, and masses of the spheres are listed in Table 2."
+ Note that the total mass is zero., Note that the total mass is zero.
+ When measured in he coarse eric Lol {p2 he cell aveage density. vaulshes in all the cells since ;V61 in this example.," When measured in the coarse grid of $ \ell \, \le \, 2 $, the cell average density vanishes in all the cells since $ N \, = \, 64 $ in this example."
+ Iu οἱher words. ἰiere are no source terms iu the discrete Poisson equation at the level of (=1.," In other words, there are no source terms in the discrete Poisson equation at the level of $ \ell \, = \, 1 $."
+ Nevertljeless. the quadraple mouent of gravity is reproduced in Figure 7..," Nevertheless, the quadraple moment of gravity is reproduced in Figure \ref{Q-pole.eps}."
+ This proves tliat Equation (17)) counects tle gravitatioial potential between the coarse and fine grids., This proves that Equation \ref{gauss}) ) connects the gravitational potential between the coarse and fine grids.
+ We measwed reduction iu the residual. the «difference iu the rieh and left. Baud: sides of Equatio1 (13). to evaluate the efficieucy ofthe iteration.," We measured reduction in the residual, the difference in the right and left hand sides of Equation \ref{dpoisson1}) ), to evaluate the efficiency of the iteration."
+ Figu ‘eS shows he residual as a fuuctior of the number ol the FMC iteration cycle lor ie sample problem lor ju:uy shown iu Ll., Figure \ref{res-nmg.eps} shows the residual as a function of the number of the FMG iteration cycle for the sample problem for binary shown in \ref{accuracy}.
+ The a»olute value of the residual redices by a factor of 300 per oue cycle of un|ti-grk )teratior intle first four times of the EMG ite‘ation., The absolute value of the residual reduces by a factor of 300 per one cycle of multi-grid iteration in the first four times of the FMG iteration.
+ The residual reialiis CODslant after the five times F1C ieratioL, The residual remains constant after the five times FMG iteration.
+ The remaining residualal is due to the ‘ound o [Te‘ror in the gravitationa potential, The remaining residual is due to the round off error in the gravitational potential.
+ The resilal is proportional to the 1iverse square of the ce| size in Figure & siice the second «erlvaive ls slotted., The residual is proportional to the inverse square of the cell size in Figure \ref{res-nmg.eps} since the second derivative is plotted.
+ The rouud olf error is extremely small aic is iegligible compared with the cliscretiza lone‘ror discussed iu Ll. , The round off error is extremely small and is negligible compared with the discretization error discussed in \ref{accuracy}. .
+Trying ses'eral yocel density distributions. we have confirmed that the reduction depends little ou the spatial distribution of the residual.," Trying several model density distributions, we have confirmed that the reduction depends little on the spatial distribution of the residual."
+which case. upstream. «—1.,"which case, upstream, $u \rightarrow 1$."
+ For w=1. this reduces to Since 1+ji0 everywhere iu the integration interval except of course at ys=—1. aid since [>0 we see that. for vw=1 we must have f.xo(p+1) where d(r) is Dirac’s delta.," For $u=1$, this reduces to Since $1+\mu > 0$ everywhere in the integration interval except of course at $\mu = -1$, and since $f\geq 0$ we see that, for $u = 1$ we must have $f \propto \delta(\mu+1)$ where $\delta(x)$ is Dirac's delta."
+ Thus. in his limit. the augular dependence factors out.," Thus, in this limit, the angular dependence factors out."
+ For reasons to be explaiued in the next Subsection. we shall also need f lor 1—wu«1. but still 40.," For reasons to be explained in the next Subsection, we shall also need $f$ for $ 1-u \ll 1$, but still $\neq 0$."
+ To search [or such a solution. we let ourselves je guided by the solution at &=I: thus. we let the angular dependence factor out. aud use theAnsalz f=gleyu(Ge+Y)/h(u)).," To search for such a solution, we let ourselves be guided by the solution at $u =1$: thus, we let the angular dependence factor out, and use the $f = g(z) w((\mu+1)/h(u))$."
+ Here h(i) is an as yet undetermined function of the pre-shock Iuid speed such that Go)—0 as τε—1., Here $h(u)$ is an as yet undetermined function of the pre–shock fluid speed such that $h(u) \rightarrow 0$ as $u \rightarrow 1$.
+ In this way. as the speed grows. the augular dependeuce )ecomes more and more coucentrated toward jp——1. as required by the previously found solution Or uw—1.," In this way, as the speed grows, the angular dependence becomes more and more concentrated toward $\mu = -1$, as required by the previously found solution for $u = 1$."
+ Introduciug ourAnsatz into Eq., Introducing our into Eq.
+" 1 I fiud where I defined DD4,=DG—1.p). and I factored the eigenvalue for future convenience."," \ref{main} I find where I defined $D_{-1} \equiv D(\mu = -1, p)$, and I factored the eigenvalue for future convenience."
+ on1⋅ the angular part. and defining (jt⋅⋅ 4]1)/h(u)⋅Ξy.aud w= ευη./οὐ» z2dew/dy-./Concentrating I Gud We are Interested in a solution of the above equation ouly in the limit //(:)—0. the ouly one in which our Lactored is a good approximation to the true f.," Concentrating on the angular part, and defining $(\mu+1)/h(u) \equiv y$, and $\dot{w} \equiv dw/dy$, $\ddot{w} \equiv 
+d^2 w/d y^2$, I find We are interested in a solution of the above equation only in the limit $h(u)\rightarrow 0$, the only one in which our factored is a good approximation to the true $f$."
+ Iu this case. the teri xie is clearly subdomiuant. and can be neglected iu a first approximation (this techuique is called. domiuaut balance. Bender aud Orszag 1978).," In this case, the term $\propto \dot{w}$ is clearly subdominant, and can be neglected in a first approximation (this technique is called dominant balance, Bender and Orszag 1978)."
+ Furthermore. for (4)—0. we expect wy) to be more aud more concentrated around fo——1. so that in this range we can approximate the term D(gi.p)(1— pi). aud E obtain azΑπο with obvious solution wzwpexp(—A(gi+1)/h(u)).," Furthermore, for $h(u)\rightarrow 0$, we expect $w(\mu)$ to be more and more concentrated around $\mu = -1$, so that in this range we can approximate the term $D(\mu,p)(1-\mu^2) \approx 2 D_{-1} (1+\mu)$ , and I obtain $\ddot{w} \approx 
+\lambda^2 w$ with obvious solution $w \approx w_0 \exp{(-\lambda 
+(\mu+1)/h(u))}$."
+ The factor Affi() cau be determined by inserting this approximate expression for w into Eq., The factor $\lambda/h(u)$ can be determined by inserting this approximate expression for $w$ into Eq.
+ κ)+) 3aa. A trivial computation yields A/h(u)=1/(1— a)., \ref{conserved2}a a. A trivial computation yields $\lambda/h(u) = 1/(1-u)$ .
+ Now. going back to the equation for g. the spatial part of f. we fiud. also using the above. from which. in the eud. I find an approximate solution for the distribution function iuthe limit ul: It is thus seen that the detailed shape of the pitch angle scattering function Dr.p) is irrelevant. and that what is left of it (its value D.4 at pp= —1) only enters the spatial part of thedistribution[uuction f. not the augular one.," Now, going back to the equation for $g$, the spatial part of $f$ , we find, also using the above, from which, in the end, I find an approximate solution for the distribution function inthe limit $u\rightarrow 1$: It is thus seen that the detailed shape of the pitch angle scattering function $D(\mu, p)$ is irrelevant, and that what is left of it (its value $D_{-1}$ at $\mu = -1$ ) only enters the spatial part of thedistributionfunction $f$ , not the angular one."
+The first step in this analysis is to compute the diffusion map representation of each SN.,The first step in this analysis is to compute the diffusion map representation of each SN.
+" In Figure[}] we plot diffusion map coordinates, coloured by true class, for the entire SN set, D (top) and spectroscopic set, S (bottom)."," In Figure \ref{fig:impnoz} we plot diffusion map coordinates, coloured by true class, for the entire SN set, $\mathcal{D}$ (top) and spectroscopic set, $\mathcal{S}$ (bottom)."
+ These coordinates were computed without using host photometric redshifts., These coordinates were computed without using host photometric redshifts.
+ Large discrepancies between the diffusion map distributions of these two sets are obvious., Large discrepancies between the diffusion map distributions of these two sets are obvious.
+" The left panels of Figure show the first two diffusion coordinates, where clear Blseparation between the Type I and Type II SNe is obvious."," The left panels of Figure \ref{fig:impnoz} show the first two diffusion coordinates, where clear separation between the Type I and Type II SNe is obvious."
+" 'The right panels plot the two most important coordinates (3 7), as estimated by a random forest classifier trained only on the spectroscopic data."," The right panels plot the two most important coordinates (3 7), as estimated by a random forest classifier trained only on the spectroscopic data."
+" With respect to these two coordinates, there is separation between the Type Ia and Ibc supernovae in the spectroscopic set."," With respect to these two coordinates, there is separation between the Type Ia and Ibc supernovae in the spectroscopic set."
+" In Fig. [],"," In Fig. \ref{fig:lcdmap},"
+ we plot the average 4-band light curves of the SNe within each of 4 regions in diffusion space., we plot the average 4-band light curves of the SNe within each of 4 regions in diffusion space.
+" Stepping across this coordinate system (with respect to coordinates 3 and 7), we see that the relative strength of the g-band emission increases and the light curves become more plateaued."," Stepping across this coordinate system (with respect to coordinates 3 and 7), we see that the relative strength of the $g$ -band emission increases and the light curves become more plateaued."
+ This corresponds to a gradual transition from Type Ia to Type II SNe., This corresponds to a gradual transition from Type Ia to Type II SNe.
+" We also explore the behavior, in diffusion space, of SN redshift."," We also explore the behavior, in diffusion space, of SN redshift."
+ Fig., Fig.
+" [7] plots the first two diffusion coordinates of the 5088 Ia SNe, coloured by their true redshift."," \ref{fig:zdmap} plots the first two diffusion coordinates of the 5088 Ia SNe, coloured by their true redshift."
+" Even though we did not use any redshift information in the computation of these coordinates, we see a gradual trend in redshift with respect to this coordinate system, with an increase in z as one moves diagonally from bottom right to top left."," Even though we did not use any redshift information in the computation of these coordinates, we see a gradual trend in redshift with respect to this coordinate system, with an increase in $z$ as one moves diagonally from bottom right to top left."
+ This means that our construction of distance measure for the diffusion map captures the incremental changes with respect to redshift that occur in the light curves., This means that our construction of distance measure for the diffusion map captures the incremental changes with respect to redshift that occur in the light curves.
+ Note that using the entire set of data to estimate the diffusion map coordinates has allowed us a fine enough sampling with respect to redshift; this would generally not be the case if we were to employ only the spectroscopic SNe to build the, Note that using the entire set of data to estimate the diffusion map coordinates has allowed us a fine enough sampling with respect to redshift; this would generally not be the case if we were to employ only the spectroscopic SNe to build the
+hydrogen οσοπια deusities do allow a crude upper linüt. however.,"hydrogen column densities do allow a crude upper limit, however."
+ À ranec of values are derived (see references above). but all imply MyE7<1075 ec.7.," A range of values are derived (see references above), but all imply $N_{\rm H} \la
+7\times10^{21}$ $^{-2}$."
+ Similarly. a rouge of eas-to-dust scalines exist. but that of Savage.&Drake (1978).. with an assumed the variance between lines of sight.," Similarly, a range of gas-to-dust scalings exist, but that of \citet{Bohlin:1978a}, , with an assumed the variance between lines of sight."
+ Using this scaliug and the upper limit ou Ng nuples E(B—V)z1.6., Using this scaling and the upper limit on $N_{\rm H}$ implies $\ebv \la 1.6$.
+ The actual value could well be 1iuch less than this., The actual value could well be much less than this.
+ Aw independent check is provided by the Galactic dust maps of Schlegel.Fiukbeiuer.&Davis(1998) which imply LBVV)~0.9. although their use for Galactic latitudes loss than 5° is not recomiuenuded by the autlors.," An independent check is provided by the Galactic dust maps of \citet{Schlegel:1998a} which imply $E(B-V)\sim0.9$, although their use for Galactic latitudes less than $5^{\circ}$ is not recommended by the authors."
+ The Iuninositv of the transition between ligh/soft aud low/liud. states in black holes is of iuterest (Maccarone2003) and it may even be an approximate standard candle., The luminosity of the transition between high/soft and low/hard states in black holes is of interest \citep{Maccarone:2003a} and it may even be an approximate standard candle.
+" As already. noted. Maccarone(2003). estimated a ΙΙ. distance of kkpe assimnuiuse the miuimuiuu ack hole mass is elven by the mass function of 5.8+JS MM, (ντοςetal.2003).."," As already noted, \citet{Maccarone:2003a} estimated a minimum distance of kpc assuming the minimum black hole mass is given by the mass function of $5.8\pm0.5$ $_{\odot}$ \citep{Hynes:2003a}."
+ The tue mass is likely to 0 slenificantly larecr than this. however. ax the mass Wuction is only approached for high inclinations aud here is no evidence that thas a high inclination: indeed 339)it has been argued that he inclination is very low (Wietal.2001)..," The true mass is likely to be significantly larger than this, however, as the mass function is only approached for high inclinations and there is no evidence that has a high inclination; indeed it has been argued that the inclination is very low \citep{Wu:2001a}."
+ For actual ack holes masses of 10. 15. or MAL... distances of iyout 10. 12. or kkpe respectively would be expected sed on the method of Maccarone(2003)..," For actual black holes masses of 10, 15, or $_{\odot}$, distances of about 10, 12, or kpc respectively would be expected based on the method of \citet{Maccarone:2003a}."
+. These are very reasonable if lios somewhat bevoud the tangent point as we have argued., These are very reasonable if lies somewhat beyond the tangent point as we have argued.
+ If it is really bevoud 15kkpc. however. then either the black hole mass is relatively high. or the ransition huuinositv is rather Ligh (but not implausibly πο).," If it is really beyond kpc, however, then either the black hole mass is relatively high, or the transition luminosity is rather high (but not implausibly so)."
+ Zdziarskictal.(1998). have instead tried to use the naxiuuni observed IDmuumositv as a constraint on the distance., \citet{Zdziarski:1998a} have instead tried to use the maximum observed luminosity as a constraint on the distance.
+" We can reverse the arguineut. however aud exaiine the maxiniuun huninosity for different possible distances,"," We can reverse the argument, however and examine the maximum luminosity for different possible distances."
+ At 6kkpe. their quoted maxima observed muunosity corresponds to just 10M. black hole.," At kpc, their quoted maximum observed luminosity corresponds to just $_{\odot}$ black hole."
+ This secnm rather low. and a Heher black hole mass would further exacerbate this.," This seems rather low, and a higher black hole mass would further exacerbate this."
+ Iu contrast. for a distance of kkpe and black hole nasses of NINE. the Eddington ratio is 0.3.0.6. which seems much more in line with other Galactic black tole X-ray binarics.," In contrast, for a distance of kpc and black hole masses of $_{\odot}$, the Eddington ratio is 0.3–0.6, which seems much more in line with other Galactic black hole X-ray binaries."
+ More recent observations provide nore stringeut constraimts. viclding au kkeV Bux of 5.1«10Seereccau 28s 3 (extrapolated frou a I00XkeV fit: Woman et iin preparation).," More recent observations provide more stringent constraints, yielding an keV flux of $5.4\times10^{-8}$ $^{-2}$ $^{-1}$ (extrapolated from a keV fit; Homan et in preparation)."
+ For a MINI. black hole at distances of 6 aud 15kkpe respectively this corresponds to readily be reconciled with a black hole somewhat more massive than LOAMAL..., For a $_{\odot}$ black hole at distances of 6 and kpc respectively this corresponds to readily be reconciled with a black hole somewhat more massive than $_{\odot}$.
+ Thus the Ipuuinositv is consistent with a relatively huge distance given the implied high mass for the black hole., Thus the luminosity is consistent with a relatively large distance given the implied high mass for the black hole.
+" The X-aav bhuunmositv in the lowest observed state (Gallo.Feuder.&Corbel2003). correspouds to 2«&lü Poor Pos10° + for distances of 6 or 15kkpe respectively,"," The X-ray luminosity in the lowest observed state \citep{Gallo:2003a}
+ corresponds to $2\times10^{33}$ $^{-1}$ or $1\times10^{34}$ $^{-1}$ for distances of 6 or kpc respectively."
+ Most quiescent black holes lave hinuinositics of betwee12.107 and 6&10 sst (Careiaetal.2001: 81tariietal.2002:: Tlameuryetal.2003: Toussicketal.2003:: MceCliutocketal.2003)3., Most quiescent black holes have luminosities of between $2\times10^{30}$ and $6\times10^{31}$ $^{-1}$ \citealt{Garcia:2001a};; \citealt{Sutaria:2002a}; \citealt{Hameury:2003a}; \citealt{Tomsick:2003a}; \citealt{McClintock:2003a}) ).
+ The exception is VAOL Cre (ος109 eeygss1: Garciaetal.2001))., The exception is V404 Cyg $5\times10^{33}$ $^{-1}$; \citealt{Garcia:2001a}) ).
+" This is a long period object ddav) aud may be something of an anomaly amoung quiesceut black holes: indeed it exhibits dramatic N-vav variability (Wagnerotal.199I: Nongetal.x2002: IIvnes et iin preparation) which suggests that it nav not be im a fully quiescenut state,", This is a long period object day) and may be something of an anomaly among quiescent black holes; indeed it exhibits dramatic X-ray variability \citealt{Wagner:1994a}; \citealt{Kong:2002a}; Hynes et in preparation) which suggests that it may not be in a fully quiescent state.
+ aappears to339. have a comparable or larger hnuuiünositv. at muni light. so its lowest observed state may also not correspond to the typical quicscent states of trausicut black hole binarics.," appears to have a comparable or larger luminosity at minimum light, so its lowest observed state may also not correspond to the typical quiescent states of transient black hole binaries."
+ A relatively active system is appealing in the light of the non-detection of spectral features from the conrpaiion star (Shahbaz.Fender.&Charles2001).., A relatively active system is appealing in the light of the non-detection of spectral features from the companion star \citep{Shahbaz:2001a}.
+ Two explanations exist for this: either the accretion flow is so hunuinous that it masks the companion star. or the colupalion is an carly-type star with uo spectral features in the region studied by Shahbaz.Fender.&Charles(2001)..," Two explanations exist for this: either the accretion flow is so luminous that it masks the companion star, or the companion is an early-type star with no spectral features in the region studied by \citet{Shahbaz:2001a}."
+. As noted above. a relatively active accretion flow already favors a larger distance. and even kkpe eives a plausible luminosity.," As noted above, a relatively active accretion flow already favors a larger distance, and even kpc gives a plausible luminosity."
+ If we require an early-type companion then its luminosity will be larger than for the typical I& or AL companions of Galactic black holes. aud so a large distance might also required to avoid it being visible.," If we require an early-type companion then its luminosity will be larger than for the typical K or M companions of Galactic black holes, and so a large distance might also required to avoid it being visible."
+ With no detection of the photospheric Lines of the companion. its nature is poorly constrained.," With no detection of the photospheric lines of the companion, its nature is poorly constrained."
+ The colupanion star should be somewhat evolved. as a L.iddayv orbital period (σοςetal.2003) implies a iiemn density for the companion of only 0.06ccm7," The companion star should be somewhat evolved, as a day orbital period \citep{Hynes:2003a} implies a mean density for the companion of only $^{-3}$."
+ This is a factor of teu below that obtained by σσShalibaz.Feuder.&Charles(2001) who assumed a shorter orbital period., This is a factor of ten below that obtained by \citet{Shahbaz:2001a} who assumed a shorter orbital period.
+ Such a low density is reached on the ain sequence only for BO stars: it is uulikelv that the colupalion is such a star. however. as it would then be expected to dominate the total liebt at low Iuninosities. and also significantly affect the motion of the black hole.," Such a low density is reached on the main sequence only for B0 stars; it is unlikely that the companion is such a star, however, as it would then be expected to dominate the total light at low luminosities, and also significantly affect the motion of the black hole."
+ Hence the companion of hhas a lower density than XNa main-sequence star of the sale temperature., Hence the companion of has a lower density than a main-sequence star of the same temperature.
+ It is probably a sub-giaut. as late-type eiauts have even lower deusities than we derive forL.," It is probably a sub-giant, as late-type giants have even lower densities than we derive for."
+. For relatively carly types. the companion mass ds also significant and should result iu measurable orbital motion of the accretion disk.," For relatively early types, the companion mass is also significant and should result in measurable orbital motion of the accretion disk."
+ ντοςetal.(2003). found uo such evidence (only a weak modulation was present. with the wrong pliase-depeudeuce) aud suggested a mass ratio of 4z0.08.," \citet{Hynes:2003a} found no such evidence (only a weak modulation was present, with the wrong phase-dependence) and suggested a mass ratio of $q \la 0.08$."
+ This argues for a coupanuion mass of NIME. or less unless the black hole mass is significantly larger than iu other objects (2 20MM.) and hence a spectral type of Foor later unless the companion is hotterthan a iain sequence star of the same mass., This argues for a companion mass of $_{\odot}$ or less unless the black hole mass is significantly larger than in other objects $\ga 20$ $_{\odot}$ ) and hence a spectral type of F or later unless the companion is hotterthan a main sequence star of the same mass.
+ Iu support of these aremuenuts. JJ1655.10 has a comparable orbital period (2.6dd). aud an F companion (Orosz&Bailvn 1997: Shahbazctal. 19993).," In support of these arguments, J1655–40 has a comparable orbital period d), and an F companion \citealt{Orosz:1997a}; ; \citealt{Shahbaz:1999a}) )."
+ Its mass ratio is 0.12£0.03 (Shahbaz 2003)..however. the disk lues move with an amplitude of 76.2+7.5 | (Soriaetal. 1998).. and the companion totally dominates the optical light iu quiescence aud remains visible even," Its mass ratio is $0.42\pm0.03$ \citep{Shahbaz:2003a}, ,however, the disk lines move with an amplitude of $76.2\pm7.5$ $^{-1}$ \citep{Soria:1998a}, , and the companion totally dominates the optical light in quiescence and remains visible even"
+The prompt cussion of gamma-ray bursts is characterized by the diversity of the observed. temporal profiles.,The prompt emission of gamma-ray bursts is characterized by the diversity of the observed temporal profiles.
+ Some bursts show a simple shape with a fast rise followed by a slower decay while others have a complex structure with a succession of pulses which can be overlapping or separated by intervals with almost no emission., Some bursts show a simple shape with a fast rise followed by a slower decay while others have a complex structure with a succession of pulses which can be overlapping or separated by intervals with almost no emission.
+ Couverselv the spectra are niore uniformi. ecucrally well fitted by two snootlly connected power laws (the so-called Band spectrum: Daud et al..," Conversely the spectra are more uniform, generally well fitted by two smoothly connected power laws (the so-called Band spectrum; Band et al.,"
+ 1993)., 1993).
+ May studies have tried to ink the temporal aud spectral properties of bursts with ie objective to eain insight mto the physical processes eoverning the prompt emission., Many studies have tried to link the temporal and spectral properties of bursts with the objective to gain insight into the physical processes governing the prompt emission.
+ Already before and duriug ie BATSE cra several relations between hardness aud duration (Ixouveliotou et al. 1993). iteusitv (Colenetskil et ab.," Already before and during the BATSE era several relations between hardness and duration (Kouveliotou et al, 1993), intensity (Golenetskii et al.,"
+ 1983) aud fluence (Liang INargatis. 1996) were uad.," 1983) and fluence (Liang Kargatis, 1996) were found."
+ Followingo the discovery of the afterelows and re measure of the first redshifts. intrinsic quautities such as the LDunünositvi or the total radiated ΟΠΟΙΟΥ: vecamne accessible and new relations appeared: tle Amati (Amati ct al. 2002) aud Chirlauda (Cchirlanda et al.," Following the discovery of the afterglows and the measure of the first redshifts, intrinsic quantities such as the luminosity or the total radiated energy became accessible and new relations appeared: the Amati (Amati et al, 2002) and Ghirlanda (Ghirlanda et al.,"
+ 200 relations between the peak. enerev of the global spectrum and the euergy release in gamma-rays (assuming isotropic euuission iu the Amati relation auc corrected or beaming iu the Chirlanda relation). the bpuuinositv-variability relation (Reichart et al.," 2004) relations between the peak energy of the global spectrum and the energy release in gamma-rays (assuming isotropic emission in the Amati relation and corrected for beaming in the Ghirlanda relation), the luminosity-variability relation (Reichart et al.,"
+ 2001) illustrating the enudeuev of Iuminous bursts to be more variable and 1e lag-Iuniinositv relation (hereafter LLR) discovered by Norris et al. (, 2001) illustrating the tendency of luminous bursts to be more variable and the lag-luminosity relation (hereafter LLR) discovered by Norris et al. (
+2000).,2000).
+ Spectral lags are a way to quantifv the changes in the burst profiles observed in different energy xuids., Spectral lags are a way to quantify the changes in the burst profiles observed in different energy bands.
+ When viewed at high energy. pulses are narrower and peak earlier.," When viewed at high energy, pulses are narrower and peak earlier."
+ Norris ct al. (, Norris et al. (
+2000) cross-correlated ie profiles between BATSE bands b aud 3 aud fone iat the resulting lags were decreasing with imereasiug nust peak bhuuinositv.,2000) cross-correlated the profiles between BATSE bands 1 and 3 and found that the resulting lags were decreasing with increasing burst peak luminosity.
+ As for the other relations between iuninositv and spectral or temporal properties. the LER offers clues to the plivsies of the prompt enission but also oxovides a potential method to evaluate CRB distances roni observations at high energy oulv.," As for the other relations between luminosity and spectral or temporal properties, the LLR offers clues to the physics of the prompt emission but also provides a potential method to evaluate GRB distances from observations at high energy only."
+ Spectral lags are a direct consequence of the burs spectral evolution since a fixed. constant spectrum. would lead to proportional profiles in all energy binds.," Spectral lags are a direct consequence of the burst spectral evolution since a fixed, constant spectrum, would lead to proportional profiles in all energy bands."
+ Iu a firs paper (Wafizi Mochkovitch. 2007) we computed spectral lags of pulses. defined as the time interval between pulse aNd in two different bands.," In a first paper (Hafizi Mochkovitch, 2007) we computed spectral lags of pulses, defined as the time interval between pulse maximum in two different bands."
+" We obtained au explici expression for the lags and assmuiug the validity of an ~Amati-like” relation between Iumiuositv and the value of L, at pulse maximus we were able to connect lags au DIuninuositv.", We obtained an explicit expression for the lags and assuming the validity of an “Amati-like” relation between luminosity and the value of $E_{\rm p}$ at pulse maximum we were able to connect lags and luminosity.
+ Ilakkila et al. (, Hakkila et al. (
+2008) have receutlv recousidered the LLR auc obtained a new relation which applies to individual pulses while the original Norris et al. (,2008) have recently reconsidered the LLR and obtained a new relation which applies to individual pulses while the original Norris et al. (
+2000) LLR considered the burst as a whole.,2000) LLR considered the burst as a whole.
+ Moreover. they also found a correlation between pulse duration aud Iuninositv (hereafter DLR).," Moreover, they also found a correlation between pulse duration and luminosity (hereafter DLR)."
+ These results offer the vossibility. to directly test and extend our previous work (Παν Mochisovitch. 2007).," These results offer the possibility to directly test and extend our previous work (Hafizi Mochkovitch, 2007)."
+ We start iu Sect.2 by comparing to observations our theoretical results for the LLR., We start in Sect.2 by comparing to observations our theoretical results for the LLR.
+ Simce they rely ou the validity of the DLR we propose to test it iu Sect.3 by comparing a svuthetic population of GRB pulses to the observed peak flux duration diagram of a sample of pulses collected by Takkila Cruubee (2009)., Since they rely on the validity of the DLR we propose to test it in Sect.3 by comparing a synthetic population of GRB pulses to the observed peak flux – duration diagram of a sample of pulses collected by Hakkila Cumbee (2009).
+ We discuss our results in Sect. [aud Sect.5 is the conclusion., We discuss our results in Sect.4 and Sect.5 is the conclusion.
+ Iu paper I (Ilafizi Mochisovitch. 2007) we presented a simple analvtical model to caleulate spectral lags.," In paper I (Hafizi Mochkovitch, 2007) we presented a simple analytical model to calculate spectral lags."
+ We recognized that lags were better defined using individual pulses rather than the whole burst profile., We recognized that lags were better defined using individual pulses rather than the whole burst profile.
+ Pulses iu the sale burst cau have differeut lags and IHakkila ct al. (, Pulses in the same burst can have different lags and Hakkila et al. (
+2008) have shown that the elobal lag represeuts some average where the brightest pulse Gvhich gencrally has the shortest lag) makes the dominant contribution.,2008) have shown that the global lag represents some average where the brightest pulse (which generally has the shortest lag) makes the dominant contribution.
+ Looking for correlation between lag aud luninosity it is therefore preferable to consider cach pulse separately., Looking for correlation between lag and luminosity it is therefore preferable to consider each pulse separately.
+ Iakkila et al. (, Hakkila et al. (
+2008) obtained,2008) obtained
+reduction of η aud the dry. air case represents the worst case designing.,reduction of $n$ and the dry air case represents the worst case designing.
+ To obtain the light path through the atmosphere. we need to know the vertical profiles of p(/) aud T(/). where / is altitude.," To obtain the light path through the atmosphere, we need to know the vertical profiles of $p(h)$ and $T(h)$, where $h$ is altitude."
+ We adopt au atmosphere inodel of the Earth used in the field of aeronautics. in which the pressure and temperature are given as analytic functions of the altitude. A (http://www.gre.tasa.gov/W," We adopt an atmosphere model of the Earth used in the field of aeronautics, in which the pressure and temperature are given as analytic functions of the altitude, $h$ (http://www.grc.nasa.gov/WWW/K-12/airplane/atmosmet.html)."
+ This model covers from the troposphere (1>0 kia) to the stratosphere (h<50 km). while turbulent lavers of sienificance[we for optical wave propagatione is located. at altitucdes below the lower stratosphere (7;<25 km).," This model covers from the troposphere $h > 0$ km) to the stratosphere $h < 50$ km), while turbulent layers of significance for optical wave propagation is located at altitudes below the lower stratosphere $h < 25$ km)."
+ Then the refractive iudex nm becomes a fuuctiou of À. aud A.," Then the refractive index $n$ becomes a function of $\lambda$, and $h$."
+ If tlie zenith distauce of a star incident ou the upper atmosphere froin space Is zo. the appareut zenitli distance. z(A.Zi). is related to zy by At an arbitrary altitude in the atmosphere. /. Suell’s law can be expressed in a differential form as. Or Iu principle. we must integrate this differential relation (1)) aloug the light path.," If the zenith distance of a star incident on the upper atmosphere from space is $z_0$, the apparent zenith distance, $z(\lambda,h)$, is related to $z_0$ by At an arbitrary altitude in the atmosphere, $h$, Snell's law can be expressed in a differential form as, or In principle, we must integrate this differential relation \ref{snell-2}) ) along the light path."
+ However. for the practical purpose. a more useful approximate expression for the bending angle. 0A./1). not iu a differential form. is obtained. since the index of refraction of the air is very close to that of the vacuum and the total beudiug augle is also small.," However, for the practical purpose, a more useful approximate expression for the bending angle, $\Delta \theta(\lambda,h)$, not in a differential form, is obtained, since the index of refraction of the air is very close to that of the vacuum and the total bending angle is also small."
+ By takiug the first order terms of eq(2)). To evaluate CS. we need to know the separation of the two beans at Ας aud Agey. perpeucdicular to the line of sight iu leugth. r. at a given altitude. // aud zenith distauce. zo.," By taking the first order terms of \ref{snell3}) ), To evaluate CS, we need to know the separation of the two beams at $\lambda_{WFS}$ and $\lambda_{SCI}$, perpendicular to the line of sight in length, $r$, at a given altitude, $h$ and zenith distance, $z_0$."
+ Iu actual observing with an AO system. an atinosphlieric dispersion corrector will be used aud the two beams are aligned iuskle the optical system.," In actual observing with an AO system, an atmospheric dispersion corrector will be used and the two beams are aligned inside the optical system."
+ The two beams gradually separate towards the upper atinosphliere aid eventually become parallel above tlie stratosphere., The two beams gradually separate towards the upper atmosphere and eventually become parallel above the stratosphere.
+ Since the angle between the two beams. a is given by," Since the angle between the two beams, $\alpha$ is given by"
+Section 4..,Section \ref{sum}.
+ Our model is shown as a schematic picture in Figure 1.., Our model is shown as a schematic picture in Figure \ref{modpic}.
+ The emission process we consider is shown in the left panel in Figure 1. and when comparing with the observational data. we determine the direction of the observer as shown in (he right panel in Figure l..," The emission process we consider is shown in the left panel in Figure \ref{modpic}
+ and when comparing with the observational data, we determine the direction of the observer as shown in the right panel in Figure \ref{modpic}."
+ First of all. electrons and/or positrons are injected isotropically al the location of the compact object.," First of all, electrons and/or positrons are injected isotropically at the location of the compact object."
+ They cool only by IC process as stated Section 1.. and radiate hish energy photons in various directions in 3-dimensional space. when they interact with thermal radiation from the companion star.," They cool only by IC process as stated Section \ref{intro}, and radiate high energy photons in various directions in 3-dimensional space, when they interact with thermal radiation from the companion star."
+ The high enerev photons also interact wilh photons from the companion and create ἐπ pairs., The high energy photons also interact with photons from the companion and create $e^{\pm}$ pairs.
+ We ienore the propagation of these electrons and positrons. so that they radiate IC photons at the created (and injected) places.," We ignore the propagation of these electrons and positrons, so that they radiate IC photons at the created (and injected) places."
+ Thus. these electrons and positrons initiate (he cascade processes.," Thus, these electrons and positrons initiate the cascade processes."
+ We count photons which undergo such a cascade process aud escape from the svstem in anv direction., We count photons which undergo such a cascade process and escape from the system in any direction.
+ By rotating the direction of the observer. we can obtain spectra bv photons escaping in the direction of a certain range of azimutlial angle (the right panel of Figure 1)).," By rotating the direction of the observer, we can obtain spectra by photons escaping in the direction of a certain range of azimuthal angle (the right panel of Figure \ref{modpic}) )."
+" These are equivalent to spectra by escaped photons during a relevant range of the orbital phase (hereafter we call the spectra ""phase-divided spectra”)."," These are equivalent to spectra by escaped photons during a relevant range of the orbital phase (hereafter we call the spectra ""phase-divided spectra"")."
+ At the same time. we can obtain the light curves in GeV-TeV energy range in any inclination angle. when we assume that the emission is stationary during each divided phase.," At the same time, we can obtain the light curves in GeV-TeV energy range in any inclination angle, when we assume that the emission is stationary during each divided phase."
+ We also calculate spectra of svuchrotron radiation bv electrons aud positrons in the svstem. assuming uniform magnetic field and using the electron distribution in steady state derived by cascade process.," We also calculate spectra of synchrotron radiation by electrons and positrons in the system, assuming uniform magnetic field and using the electron distribution in steady state derived by cascade process."
+ Below. we reler (o each process (Section 2.1.. 2.2)). validity. (Section 2.3)) aud implication (Section 2.4)) of the assumptions. aud caleulational method (Section 2.5)).," Below, we refer to each process (Section \ref{inj}, , \ref{radipro}) ), validity (Section \ref{validity}) ) and implication (Section \ref{maglim}) ) of the assumptions, and calculational method (Section \ref{method}))."
+observation.,observation.
+" However, as can be seen in Figure 1 the model is a good first order approximation of the data at a few percentage level."," However, as can be seen in Figure \ref{spec} the model is a good first order approximation of the data at a few percentage level."
+" T'he essence of this spectral analysis is that with a high energy spectral index of ~1.7, moderate intrinsic absorption and soft excess, the spectrum of Mrk 1040 is typical of AGNS of its class."," The essence of this spectral analysis is that with a high energy spectral index of $\sim 1.7$, moderate intrinsic absorption and soft excess, the spectrum of Mrk 1040 is typical of AGNs of its class."
+" The top panel of Figure 2 shows the background subtracted X-ray 640 sec binned lightcurves of Mrk 1040 in soft, 0.2-2.0 keV (circles) and hard, 2.0-10.0 keV (triangles) energy bands."," The top panel of Figure \ref{lcurve} shows the background subtracted X-ray $640$ sec binned lightcurves of Mrk 1040 in soft, $0.2$ $2.0$ keV (circles) and hard, $2.0$ $10.0$ keV (triangles) energy bands."
+ There is clear strong correlated variability in the two bands on ~10 sec timescales., There is clear strong correlated variability in the two bands on $\sim 10^4$ sec timescales.
+" The bottom panel shows that the ratio of the hard to soft count rates (i.e. the hardness ratio) significantly decreases from ~0.6 to ~0.45 during the course of the observation, indicating spectral evolution of the source on long (~10° sec) timescales."," The bottom panel shows that the ratio of the hard to soft count rates (i.e. the hardness ratio) significantly decreases from $\sim 0.6$ to $\sim 0.45$ during the course of the observation, indicating spectral evolution of the source on long $\sim 10^5$ sec) timescales."
+ It is interesting to note that the variation in the hardness ratio is clearly uncorrelated with the variability of the lightcurves., It is interesting to note that the variation in the hardness ratio is clearly uncorrelated with the variability of the lightcurves.
+ This means that the primary variability of the source in ~104 sec timescales is not associated with significant spectral variation and hence its origin cannot be due to changes in the absorption medium which would effect primarily the soft band., This means that the primary variability of the source in $\sim 10^4$ sec timescales is not associated with significant spectral variation and hence its origin cannot be due to changes in the absorption medium which would effect primarily the soft band.
+" Lightcurves binned at 64 secs were generated for seven energy bands corresponding to 0.3-0.5, 0.5- 0.8, 0.8-1.0, 1.0-2.0, 2.0-3.0, 3.0-5.0 and 5.0-10.0 keV. Cross-correlation analysis was undertaken between the reference band 1.0-2.0 keV and others."," Lightcurves binned at $64$ secs were generated for seven energy bands corresponding to $0.3$ $0.5$, $0.5$ $0.8$, $0.8$ $1.0$, $1.0$ $2.0$, $2.0$ $3.0$, $3.0$ $5.0$ and $5.0$ $10.0$ keV. Cross-correlation analysis was undertaken between the reference band $1.0$ $2.0$ keV and others."
+ The XRONOS task “crosscor” was used for the analysis., The XRONOS task “crosscor” was used for the analysis.
+" The functions were normalized by dividing by the square root of the product of the excess variances of the two series (i.e. with “crosscor” flag “norm = 2"").", The functions were normalized by dividing by the square root of the product of the excess variances of the two series (i.e. with “crosscor” flag “norm = 2”).
+" The entire lightcurve was taken as a single segment and the analysis was undertaken in the ""slow"" mode (i.e. not in fast Fourier transform mode).", The entire lightcurve was taken as a single segment and the analysis was undertaken in the “slow” mode (i.e. not in fast Fourier transform mode).
+" In such a case, the task computes errors on the cross-correlation by propagating the measurement errors of each time bin."," In such a case, the task computes errors on the cross-correlation by propagating the measurement errors of each time bin."
+" Dividing the light-curve into five or more segments, gave negative variance for some segments, rendering the analysis unfeasible."," Dividing the light-curve into five or more segments, gave negative variance for some segments, rendering the analysis unfeasible."
+ T'he cross-correlations in the delay time range of —2000 to 2000 secs are plotted in Figure 3.., The cross-correlations in the delay time range of $-2000$ to $2000$ secs are plotted in Figure \ref{cross}.
+ The lightcurves are well correlated with the cross-correlation function peaking at values >0.8., The lightcurves are well correlated with the cross-correlation function peaking at values $> 0.8$.
+" In this delay range, we fit the functions with a Gaussian and obtain the best fit centroid."," In this delay range, we fit the functions with a Gaussian and obtain the best fit centroid."
+ The best fit Gaussians are plotted in the figure lines) and their centroid values are marked with a (solidhorizontal dashed line which is taken as the time-delay between the, The best fit Gaussians are plotted in the figure (solid lines) and their centroid values are marked with a horizontal dashed line which is taken as the time-delay between the
+results of Smallονοἱal.(1996).,results of \citet{sswa96}.
+. Our abundance analysis confirms that 2204188 is a mid metallic line star., Our abundance analysis confirms that 204188 is a mild metallic line star.
+ The age obained for this stzw is determined. with not very high. accusον., The age obtained for this star is determined with not very high accuracy.
+ We agree with Ixünzli&North(1998) who noticed the dilliculty. of determing the stellar. age properly when the star is near ZAMS.," We agree with \citet{kn98}
+ who noticed the difficulty of determing the stellar age properly when the star is near ZAMS."
+ Following the svstematic search for abundance anomalies of Am stars driven by tidal ellects we continued with studying the depencences between the chemical abundances: and the orbital elements of the binary systems. the projected rotational velocity and the physical parameters of Ami stars like mass. age. temperature.," Following the systematic search for abundance anomalies of Am stars driven by tidal effects we continued with studying the dependences between the chemical abundances and the orbital elements of the binary systems, the projected rotational velocity and the physical parameters of Am stars like mass, age, temperature."
+ The main advantage. of our analysis is the homogeneity of the observational material we used - all data were obtained at one telescope with the same spectrograph and. detector as well as processed. with the same data reduction package and analysed with the same code Synspoec., The main advantage of our analysis is the homogeneity of the observational material we used - all data were obtained at one telescope with the same spectrograph and detector as well as processed with the same data reduction package and analysed with the same code Synspec.
+ As d is well known Am peculiarities are. mainly manifested ov Ca deficit anc Fe overabundances anc they could. be represented in a reliable way through the values of Ca/Fe] where Ca/EFo]2Ca/H]-Fe/H]., As it is well known Am peculiarities are mainly manifested by Ca deficit and Fe overabundances and they could be represented in a reliable way through the values of [Ca/Fe] where [Ca/Fe]=[Ca/H]-[Fe/H].
+ This ratio would multiply the ellect of Am peculiarities because these two elements have shown an opposite behaviour., This ratio would multiply the effect of Am peculiarities because these two elements have shown an opposite behaviour.
+ Besides the lines of both elements. ect weaker with the increasing temperature. and consequently Ca/Fe] ratio is not that sensitive to the uncertainties of the elective temperature., Besides the lines of both elements get weaker with the increasing temperature and consequently [Ca/Fe] ratio is not that sensitive to the uncertainties of the effective temperature.
+ Up to now fifteen stars from our sample were Lully processed., Up to now fifteen stars from our sample were fully processed.
+ The results are taken from this work ane also from Paper Land Paper LH., The results are taken from this work and also from Paper I and Paper II.
+ Two of the stars in this sample. 1178449 and LISTTS were found not to be Am stars.," Two of the stars in this sample, 178449 and 18778 were found not to be Am stars."
+ Thev are put on the figures but not included in the analysis., They are put on the figures but not included in the analysis.
+ First we investigate the dependences of. Ca/Ec] from the eccentricity (see Figure 5))., First we investigate the dependences of [Ca/Fe] from the eccentricity (see Figure \ref{feccentr}) ).
+ Despite of some scatter in the data there seems to be a trend like that mentioned. earlier bv Budaj(1997). and. Hievetal.(1998)., Despite of some scatter in the data there seems to be a trend like that mentioned earlier by \citet{Budaj97} and \citet{ibzbz98}.
+.. Ehe correlation coellicient is -0.552:0.05., The correlation coefficient is $\pm$ 0.05.
+ One star. 1198391. denoted by an asterisk in the figures. is distinguished from the common trend.," One star, 198391 denoted by an asterisk in the figures, is distinguished from the common trend."
+ This star is the hottest star amongst the sample and it is not a tvpical Am star (Paper D)., This star is the hottest star amongst the sample and it is not a typical Am star (Paper I).
+ The star could. be a transitional object between Am anc the hotter HgMn stars because of its position in ΗΛΙΟ close to the region occupied by Lle\In stars and also of its abundances (Lor more details see Paper D)., The star could be a transitional object between Am and the hotter HgMn stars because of its position in HRD close to the region occupied by HgMn stars and also of its abundances (for more details see Paper I).
+ Namely. that Am peculiarities increase (Ca/Fe] decreases) with increasing cecentricity of the binary system.," Namely, that Am peculiarities increase ([Ca/Fe] decreases) with increasing eccentricity of the binary system."
+ The dependence of Ca/Ec] from the orbital period shown in Figure 6 is not so clear but still there is a tendency he metalicitv to increase towards the longer periods., The dependence of [Ca/Fe] from the orbital period shown in Figure \ref{fperiod} is not so clear but still there is a tendency the metalicity to increase towards the longer periods.
+ Of course. these two parameters characterizing the binary system. eecentricity ancl period. are not fully independent recause Of the synchronization and circularization. of the orbits.," Of course, these two parameters characterizing the binary system, eccentricity and period, are not fully independent because of the synchronization and circularization of the orbits."
+ This was the main reason to analyse only stars with »iods between LO and 200 cdavs., This was the main reason to analyse only stars with periods between 10 and 200 days.
+ The comparison of the hree parameters simultaneously. in 3D-graph. (sce Fieure 7T))- metalicity. eccentricity and period showed us that there was not anv correlation between period and eccentricity in he whole range of periods which extends even bevond 150 clas.," The comparison of the three parameters simultaneously in 3D-graph (see Figure \ref{feperiod}) ) - metalicity, eccentricity and period showed us that there was not any correlation between period and eccentricity in the whole range of periods which extends even beyond 150 days."
+ We study. also the dependence of the Ani peculiarities on the projected rotational velocity of the Am stars (see Figure 9))., We study also the dependence of the Am peculiarities on the projected rotational velocity of the Am stars (see Figure \ref{fvsini}) ).
+ In general. there is a clear trend. of increasing re peculiarity (decreasing of Ca/Fe]) towards small values of esini - the correlation coellicient is. 0:85zE0.04..," In general, there is a clear trend of increasing the peculiarity (decreasing of [Ca/Fe]) towards small values of $\vsini$ - the correlation coefficient is $\pm0.04$."
+ Besides the hottest star already mentioned. 1198391. two other stars. 1138213 and 2204188. both marginal Am stars. do not follow the common trend.," Besides the hottest star already mentioned, 198391, two other stars, 138213 and 204188, both marginal Am stars, do not follow the common trend."
+ The fact that these two stars depart [rom the clear smooth correlation of metallicity nd esin/ indicates that the rotation is not the only agent responsible for this peculiarity., The fact that these two stars depart from the clear smooth correlation of metallicity and $\vsini$ indicates that the rotation is not the only agent responsible for this peculiarity.
+ Lt confirms our claims from Mievetal.(1998) that the low Am peculiarity in these two systems is due to their small eccentricity.," It confirms our claims from \citet{ibzbz98}
+ that the low Am peculiarity in these two systems is due to their small eccentricity."
+ Again in order to check the possible correlation between the parameters, Again in order to check the possible correlation between the parameters
+existence of an instability strip within which the natural period of oscillation is excited. ancl manifest as a periodic change of luminosity.,"existence of an instability strip within which the natural period of oscillation is excited, and manifest as a periodic change of luminosity."
+ The instability strip. expressed as an additional and external mathematical constraint on the above equations. produces the observed luminosity relations by narrowing down (he range of luminosities seen al anv given. period.," The instability strip, expressed as an additional and external mathematical constraint on the above equations, produces the observed period-luminosity relations by narrowing down the range of luminosities seen at any given period."
+ The slope of the red aud blue boundaries to (he instability strip as they cut lines of constant period determines the slope of the PL relation., The slope of the red and blue boundaries to the instability strip as they cut lines of constant period determines the slope of the PL relation.
+ This was first illustrated ancl emphasized in Aladore Freedman (1991). their Figure 3.," This was first illustrated and emphasized in Madore Freedman (1991), their Figure 3."
+ Belore advancing further. we introduce the Wesenheit function Wo I).," Before advancing further, we introduce the Wesenheit function $W = V
+- R_{VI} \times (V-I)$ ."
+ This specific combination of magnitude and color is explicitly constructed so as {ο inalhematically guarantee that their sum will be numerically independent of the total amount ol interstellar extinction., This specific combination of magnitude and color is explicitly constructed so as to mathematically guarantee that their sum will be numerically independent of the total amount of interstellar extinction.
+" With foreknowledge that (bv definition) ly=Ry;xBW—I). and again having the added definitions Chat V=V+Ay and (V—7)2(V—I)4-E(V—1). with simple algebra Π-V-hx(V-l) ihen —Vede—-d;x((Q-—l)s-dxEV-I) which upon regrouping terms becomes οxBW=2) By construction and definition (he last (wo terms sum to zero. leaving Γτννιx(he-1),=1, In other words. WW. as constructed. from observed (redclenecl) colors ancl (extinetedd) magnitudes. is numerically equivalent to VW. as it would be calculated [rom intrinsic colors and intrinsic magnitudes."," With foreknowledge that (by definition) $A_V = R_{VI} \times E(V-I)$, and again having the added definitions that $V = V_o + A_V$ and $(V-I) =
+(V-I)_o + E(V-I)$, with simple algebra $W = V - R_{VI} \times (V-I) $ then $= V_o + A_V - R_{VI} \times (V-I)_o - R_{VI} \times E(V-I)$ which upon regrouping terms becomes $ W = [V_o - R_{VI} \times (V-I)_o] + [A_V - R_{VI} \times E(V-I)]$ By construction and definition the last two terms sum to zero, leaving $ W = V_o - R_{VI} \times (V-I)_o = W_o$ In other words, $W$, as constructed from observed (reddened) colors and (extincted) magnitudes, is numerically equivalent to $W_o$, as it would be calculated from intrinsic colors and intrinsic magnitudes."
+ This occurs. it must be noted. without any explicit knowledge of (he reddenings and/or extinctions involved.," This occurs, it must be noted, without any explicit knowledge of the reddenings and/or extinctions involved."
+ Simply put. VW. is reddening-independent.," Simply put, $W$ is reddening-independent."
+ The geometrical interpretation of VV is illuminating., The geometrical interpretation of $W$ is illuminating.
+" As Figure | illustrates. lines of constant VW fall. by construction. identically along reddening trajectories having a slope ol Ry,=AS/E(V—D)."," As Figure 1 illustrates, lines of constant $W$ fall, by construction, identically along reddening trajectories having a slope of $R_{VI} =
+A_V/E(V-I)$."
+ Whether a star dis at position D because it was reddened from its inirinsic position at A. or whether it is there because that is its intrinsic color and magnitude. is irrelevant: HW collapses the coloranagnitude diagram in such a wav that the distinction usually made by reddening is moot.," Whether a star is at position B because it was reddened from its intrinsic position at A, or whether it is there because that is its intrinsic color and magnitude, is irrelevant: $W$ collapses the color-magnitude diagram in such a way that the distinction usually made by reddening is moot."
+ All stars along a given reddening trajectory will have, All stars along a given reddening trajectory will have
+Llore Mou is the outflow mass rate ancl Maa the Eddington aceretion rate.,Here $\dot M_{\rm out}$ is the outflow mass rate and $\dot M_{\rm Edd}$ the Eddington accretion rate.
+ Ro=2CAL/e> is the Schwarzschilel raclius for mass AZ and yp0.1 is the assumed accretion efficiency., $R_{\rm s} = 2GM/c^2$ is the Schwarzschild radius for mass $M$ and $\eta \sim 0.1$ is the assumed accretion efficiency.
+ Since the outflow is optically thick. most of the photons have scattered ancl given up at least their original momentum.," Since the outflow is optically thick, most of the photons have scattered and given up at least their original momentum."
+ This will be assured if photon leakage from the sides of the outllow cone is small. ie the collimation of the outllow is not loo extreme.," This will be assured if photon leakage from the sides of the outflow cone is small, ie the collimation of the outflow is not too extreme."
+ As in ΙΟΣ we assume that the outllow is launched [rom the photosphere at the local escape velocity. ic With Agi—0.7M. ‘appropriate to a black hole mass of AM~3.10°M. we then find Adv AL. 1," As in KP03 we assume that the outflow is launched from the photosphere at the local escape velocity, ie With $\dot M_{\rm Edd}$ $\Msun$ $^{-1}$ , appropriate to a black hole mass of $M \sim 3 \times 10^{7}\Msun$, we then find $\dot M_{\rm out}$ $\Msun$ $^{-1}$."
+ Assuming the measured outllow velocity is the same as the launch velocity. the associated kinetie energy is 1013 erg 1.," Assuming the measured outflow velocity is the same as the launch velocity, the associated kinetic energy is $\times$$10^{44}$ erg $^{-1}$."
+ Since the outllow is optically thick for Aaa~Mg. much of the accretion. luminosity generated. deep in the potential well near A. must emerge às. blackbodw.Like emission from it.," Since the outflow is optically thick for $\mo \sim \me$, much of the accretion luminosity generated deep in the potential well near $\rs$ must emerge as blackbody–like emission from it."
+" The quasispherical racliating area is with an effective blackbody temperature. (IKDOS. equation 15) of where Al,=Mu(1M.ve+).AM,M/IOM..."," The quasi–spherical radiating area is with an effective blackbody temperature (KP03, equation 18) of where $\dot M_1 =
+\mo/(1\msun~{\rm yr}^{-1}), M_8 = M/10^8\msun$."
+ Assuming b1 in this gives Zur=6 10!S. consistent with the optically thick How being a major component of the BBB (and bolometric luminosity) of349.," Assuming $\sim$ 1 in this gives $T_{\rm eff} = 6\times 10^4$ K, consistent with the optically thick flow being a major component of the BBB (and bolometric luminosity) of."
+. Confirmation of the physical consistency of the above analysis can be mace by noting that. for a Compton thick wind driven by radiation pressure. the momentum in the outflow must be of the same order as that in the Edclineton- radiation field (IXDO03. equation 12). as is indeed the case for349.," Confirmation of the physical consistency of the above analysis can be made by noting that, for a Compton thick wind driven by radiation pressure, the momentum in the outflow must be of the same order as that in the Eddington-limited radiation field (KP03, equation 12), as is indeed the case for."
+. Furthermore. the kinetic energy in the outflow. should be lower than that ofthe radiation field by a factor of order efe (equation 13 in IXDO3). a requirement. which again is met by the outflow parameters of349.," Furthermore, the kinetic energy in the outflow should be lower than that of the radiation field by a factor of order $v/c$ (equation 13 in KP03), a requirement which again is met by the outflow parameters of."
+. In summary. we conclude that the radiation. driven. optically thick wind described. in. IXDO3. represents a physically consistent model by which to assess the high velocity. highly ionised outllows detected in aand several other AGN.," In summary, we conclude that the radiation driven, optically thick wind described in KP03 represents a physically consistent model by which to assess the high velocity, highly ionised outflows detected in and several other AGN."
+ A question remains as to how common are these phenomena. and - in turn - how common is supercritical accretion?," A question remains as to how common are these phenomena, and - in turn - how common is supercritical accretion?"
+ As noted earlier the bolometric luminosity of iis ↔⋅of order 3.⋅107⊥↽∣ erg s.1. which ⋠⋅for a reverberation. mass of M3.10°M. (Ixaspi 22000) implies accretion at close to the Eddington rate.," As noted earlier the bolometric luminosity of is of order $3\times$$ 10^{45}$ erg $^{-1}$, which for a reverberation mass of $M \sim 3 \times 10^{7}\Msun$ (Kaspi 2000) implies accretion at close to the Eddington rate."
+ D€12111143 has a very similar luminosity and. black hole mass (Ixaspi 22000) and - taken together - the oobservations suggest that à high. velocity. highly. ionised outllow may be a new signature of a high (super-Ixington?) aceretion rate in AGN.," PG1211+143 has a very similar luminosity and black hole mass (Kaspi 2000) and - taken together - the observations suggest that a high velocity, highly ionised outflow may be a new signature of a high (super-Eddington?) accretion rate in AGN."
+ A strong soft. X-ray excess (seen in earlier data as a steep soft. X-ray spectrum) then follows naturally from the physical model. sketched in Section 7.1., A strong soft X-ray excess (seen in earlier data as a steep soft X-ray spectrum) then follows naturally from the physical model sketched in Section 7.1.
+ It is interesting to note that the detailed differences in the observed. properties of aand POI211]148 are also. broadly consistent with the ohvsical model described. in. ΙΟ., It is interesting to note that the detailed differences in the observed properties of and PG1211+143 are also broadly consistent with the physical model described in KP03.
+ In. particular. we associate the more highly ionised/higher velocity outllow for wwith a smaller launch radius.," In particular, we associate the more highly ionised/higher velocity outflow for with a smaller launch radius."
+ The smaller and hotter Whotosphere matches in a general way the lower UV. flux Πο compared m13.3 from the simultaneous OM data). out “hotter soft. X-ray excess in. ccompared. with POGI211]143 (compare figure 2 in this (uper with figure 3 in Pounds 22003).," The smaller and hotter photosphere matches in a general way the lower UV flux $\sim$ 14 compared $\sim$ 13.3 from the simultaneous OM data), but `hotter' soft X-ray excess in compared with PG1211+143 (compare figure 2 in this paper with figure 3 in Pounds 2003)."
+ Another interesting dillerence evident in the simple »ower law fits to the EPIC data of aancd PG 1211|143 is the absence of an extreme broad Ee Ix emission line in the former case., Another interesting difference evident in the simple power law fits to the EPIC data of and PG 1211+143 is the absence of an extreme broad Fe K emission line in the former case.
+" While. if truly a relativistic ke Ix line. this could. relate to a cilfercnee in. reflection from the innermost accretion disc. à simpler alternative is that the partial covering by a column of moderately ionisecl gas (proposed to explain the extreme Fe Ix ""line in DOGI1211|143) is simply absent [rom the more highly ionised flow in349."," While, if truly a relativistic Fe K line, this could relate to a difference in reflection from the innermost accretion disc, a simpler alternative is that the partial covering by a column of moderately ionised gas (proposed to explain the extreme Fe K `line' in PG1211+143) is simply absent from the more highly ionised flow in."
+. In summary. the oobservations of aand DCGI211]143. provide strong evidence of à massive outllow. of ionised gas. which may. well be a characteristic of AGN accreting at or above the Iddington limit.," In summary, the observations of and PG1211+143 provide strong evidence of a massive outflow of ionised gas, which may well be a characteristic of AGN accreting at or above the Eddington limit."
+ The question of how common are these outllows might. then be transposed to the more fundamental question of how common is Eddington or super-IExldington aceretion in AGN?, The question of how common are these outflows might then be transposed to the more fundamental question of how common is Eddington or super-Eddington accretion in AGN?
+ Until recently the evidence for high velocity outllows in AGN has come from UV studies of BAL quasars. where a variety of. mainly high-ionisation. UV resonance transitions exhibit velocity widths up to 30000 kms | (e.g. Wevmann et al.," Until recently the evidence for high velocity outflows in AGN has come from UV studies of BAL quasars, where a variety of, mainly high-ionisation, UV resonance transitions exhibit velocity widths up to $\sim 30000$ km $^{-1}$ (e.g. Weymann et al."
+ 1991)., 1991).
+ The. BAL property. involving some of optically selected: quasars. is usually interpreted. as viewing the central continuum source along a particular line-of-sight tangential to the accretion disc.," The BAL property, involving some of optically selected quasars, is usually interpreted as viewing the central continuum source along a particular line-of-sight tangential to the accretion disc."
+ At X-ray energies. most BAL quasars are “x-ray weak. a property explained by πουoFsight absorption in moderately tonised gas with column densities up to Ng=lo“2em (e.g. Hamann 1998: Sabra Hamann 2001: Gallagher 22002).," At X-ray energies most BAL quasars are `x-ray weak', a property explained by line-of-sight absorption in moderately ionised gas with column densities up to $N_{H}=10^{23}\rm{cm}^{-2}$ (e.g. Hamann 1998; Sabra Hamann 2001; Gallagher 2002)."
+ Crucially. these observations havenot. revealed whether the N-ray absorbing column is moving or stationary.," Crucially, these observations havenot revealed whether the X-ray absorbing column is moving or stationary."
+ The extended: oobservation of the gravitationally lensed high redshift BAL quasar APAL 08279|5255 (Chartas, The extended observation of the gravitationally lensed high redshift BAL quasar APM 08279+5255 (Chartas
+(2009).,.
+. Prior to the caleulatious done iu Natarajan&Sclavarz(2008):BelikovIIooper(20094) all the authors have only included the homogeneous DN conrponeut. ie. have neglected. the clustering effect.," Prior to the calculations done in \citet{2008PhRvD..78j3524N,2009arXiv0904.1210B} all the authors have only included the homogeneous DM component, i.e. have neglected the clustering effect."
+ Despite the inclusion of the clustering there are some problems with the formalism of Natarajan&Sclovarz(2008):Belikov&Tooper (2009a):: the treatineut for the evolution of the matter temperature is simplistic.," Despite the inclusion of the clustering there are some problems with the formalism of \citet{2008PhRvD..78j3524N,2009arXiv0904.1210B}: the treatment for the evolution of the matter temperature is simplistic."
+ Also. in Natarajan&Sclavarz(2008).. the optical depth for he 5-ravs is calculated iucorrectlv.," Also, in \citet{2008PhRvD..78j3524N}, the optical depth for the $\gamma$ -rays is calculated incorrectly."
+ The origina work by Padinanabhan&Finkbemer(2005). aud in particular. he exteuded newer version by Slatveretal.(2009) Mises nore accurate formalis. uufortunatelv the authors have rot included he effect of DM clustering.," The original work by \citet{2005PhRvD..72b3508P}, and in particular, the extended newer version by \citet{2009arXiv0906.1197S} uses more accurate formalism, unfortunately the authors have not included the effect of DM clustering."
+ During completion of our study couple of papers appeared calculating the diffuse y-ray backeround due o the auuihilatiug DAL e.g. Kawasakictal.(2009):Profunuo&Jeltema(2009):DelikovIIooper (2," During completion of our study couple of papers appeared calculating the diffuse $\gamma$ -ray background due to the annihilating DM, e.g. \citet{2009arXiv0904.3626K,2009arXiv0906.0001P,2009arXiv0906.2251B}."
+009). INawasalkictal.2009) have neglected the 5-ravs from he inverse Compton scattering of the final sta un€ which gives even stronger bound than the primary 5-ravs.," \citet{2009arXiv0904.3626K} have neglected the $\gamma$ -rays from the inverse Compton scattering of the final state $e^-/e^+$, which gives even stronger bound than the primary $\gamma$ -rays."
+ Profumo&Jeltema(2009) have used cüffereut fonualisa or the absorptiou of 5-ravs considering the ποασ] approximation from Chen&Ixanuionkowski(2001)., \citet{2009arXiv0906.0001P} have used different formalism for the absorption of $\gamma$ -rays considering the numerical approximation from \citet{2004PhRvD..70d3502C}.
+. Tu lis paper we iuclude he relevant absorption processes directly (see the discussions below)., In this paper we include the relevant absorption processes directly (see the discussions below).
+ Our paper is organized as follows: In Section 2 we describe the source term. ic. the energy iuput from the DAL annihilation.," Our paper is organized as follows: In Section 2 we describe the source term, i.e. the energy input from the DM annihilation."
+ In Section 3 we solve the radiative transfer equation aud prescut a few example 5-ray spectra along with some results for the evolution of the ionizatiou fraction and matter temperature., In Section 3 we solve the radiative transfer equation and present a few example $\gamma$ -ray spectra along with some results for the evolution of the ionization fraction and matter temperature.
+ Section L preseuts our main results and our sunuuary is elven in Section 5., Section 4 presents our main results and our summary is given in Section 5.
+ lu our study we cousider the DM annihilation to Standard Model charged leptous following the leptonic annihilation scenario presented in Cirellictal.(2009).., In our study we consider the DM annihilation to Standard Model charged leptons following the leptonic annihilation scenario presented in \citet{2009NuPhB.813....1C}.
+ The final state distributions of clectrous/positrous. photous. aud nueutrinos/antineutrimos from the primary two-body states produced in DAL aunihilations are computed using MonteCarlo code PYTIILÀ (SjóstraudetaL.2008)..," The final state distributions of electrons/positrons, photons, and neutrinos/antineutrinos from the primary two-body states produced in DM annihilations are computed using MonteCarlo code PYTHIA \citep{2008CoPhC.178..852S}."
+ The example final particle distributious for all three anuihilation channels assumiue the input DAL particle with amass mparg= TeV are shown in Fie. l.., The example final particle distributions for all three annihilation channels assuming the input DM particle with mass $m_{DM}=1$ TeV are shown in Fig. \ref{fig1}.
+ The partition of energv between final particles is shown iu Table 1.., The partition of energy between final particles is shown in Table \ref{tab1}.
+ The released clectrous/positrons imuuediately interact with he CMB photous aud upscatter hose to hieh energies via the inverse Compton mechanisia., The released electrons/positrons immediately interact with the CMB photons and upscatter those to high energies via the inverse Compton mechanism.
+ Thus. in addition to the prompt photons released directly in the anuihilation process ourphotou spectrum contains the other part: the upscattered inverse Couptou photous.," Thus, in addition to the prompt photons released directly in the annihilation process, ourphoton spectrum contains the other part: the upscattered inverse Compton photons."
+ Having au aunihilatiug DAL particle with mass ipay and with a thermally averaged cross-section (6405 the enütted power per volume and per steradian. ro. the chussion coefficient can be given as It is the cmission coefficient for the homogeneous distribution of DM.," Having an annihilating DM particle with mass $m_{DM}$ and with a thermally averaged cross-section $\cs$ the emitted power per volume and per steradian, i.e. the emission coefficient can be given as It is the emission coefficient for the homogeneous distribution of DM."
+ To include the effect of DAL cblusteriug Ppa!2) should be replaced. bx pay!ysPalOGIàti)s= Bp) , To include the effect of DM clustering $\bar{\rho}_{DM}^2(z)$ should be replaced by $\langle\rho_{DM}^2(z)\rangle\equiv \bar{\rho}_{DM}^2(z)\langle(1+\delta(z))^2\rangle\equiv B(z)\bar{\rho}_{DM}^2(z)$ .
+Tere 6=pfp— lids the density contrast and we have also defined a halo boost factor D(:)=(1|à:1|(a?(i)., Here $\delta\equiv\rho/\bar{\rho}-1$ is the density contrast and we have also defined a halo boost factor $B(z)\equiv\langle(1+\delta(z))^2\rangle=1+\langle\delta^2(z)\rangle$.
+ To calculate. the boost factor we use the halo model which approximates the matter distribution iu the Universe as a superposition of DM halos (for a comprehensive review ou halo model see Coorav&Sheth (2002)))., To calculate the boost factor we use the halo model which approximates the matter distribution in the Universe as a superposition of DM halos (for a comprehensive review on halo model see \citet{2002PhR...372....1C}) ).
+" Within this model B(:) can be given as where ους the matter density at += 0. A,=200 is the overdensity at which the halos are defined. Mj, 1s the niunuuni halo miss. το»(AL.2) as the halo mass function. e(AL.2) veprescuts the halo concentration parameter and"," Within this model $B(z)$ can be given as where $\bar{\rho}_{m,0}$ is the matter density at $z=0$ , $\Delta_c=200$ is the overdensity at which the halos are defined, $M_{\rm min}$ is the minimum halo mass, $\frac{{\rm d}n}{{\rm d}M}(M,z)$ is the halo mass function, $c(M,z)$ represents the halo concentration parameter and"
+In 1989. Hasinger and van der Klis (1989)) divided the bright low mass X-ray binaries into two groups. which they called Z sources and atoll sources after the patterns these sources trace out in the X-ray colour-colour diagram (CD).,"In 1989, Hasinger and van der Klis \cite{hk89}) ) divided the bright low mass X-ray binaries into two groups, which they called Z sources and atoll sources after the patterns these sources trace out in the X-ray colour-colour diagram (CD)."
+" The CDs of Z sources show three branches arranged in a roughly ""Z-shaped pattern. whereas the CDs of atoll sources show a so-called island and banana. usually arranged together in one broad. curved branch."," The CDs of Z sources show three branches arranged in a roughly `Z'-shaped pattern, whereas the CDs of atoll sources show a so-called island and banana, usually arranged together in one broad, curved branch."
+ The island state is characterized by lower count rates. much less motion in the CD and much stronger band limited noise than the banana state.," The island state is characterized by lower count rates, much less motion in the CD and much stronger band limited noise than the banana state."
+ iis à burst source that was classified by Hasinger and van der Klis as an atoll source., is a burst source that was classified by Hasinger and van der Klis as an atoll source.
+ From their paper. aappears to be a special case among the atoll sources m the sense that it has a “hard island” state (1.e.. in its island state it has a harder 1—4.5 keV spectrum than in its banana state) whereas the other atoll sources which show both a banana and an island state 11735—44. 11636-54 and 11820-30) have an island state in which the spectrum in the 1—41.5 keV band is softer than in the banana.," From their paper, appears to be a special case among the atoll sources in the sense that it has a `hard island' state (i.e., in its island state it has a harder 1–4.5 keV spectrum than in its banana state) whereas the other atoll sources which show both a banana and an island state 1735–44, 1636–54 and 1820–30) have an island state in which the spectrum in the 1–4.5 keV band is softer than in the banana."
+ The island state of ccan be unambiguously identified as such from its rapid X-ray variability: it shows strong (~20 %)) band limited noise with a cutoff frequency of about 0.3 Hz in this state (Hasinger and van der Klis. 1989.. Langmeter et al.. 1989)).," The island state of can be unambiguously identified as such from its rapid X-ray variability: it shows strong $\sim 20$ ) band limited noise with a cutoff frequency of about 0.3 Hz in this state (Hasinger and van der Klis, \cite{hk89}, Langmeier et al., \cite{paper2}) )."
+ iis one of the two atoll sources known to exhibit noise with such strength., is one of the two atoll sources known to exhibit noise with such strength.
+ The other 1s (CYoshida et al. 1993))., The other is (Yoshida et al. \cite{yoshida}) ).
+ The motivation for our present analysis is to quantitatively investigate the suspected similarity (van der Klis. 1994b)) both in shape and strength of the noise component found in these two atoll sources and that in black hole candidates in the low state. which was up to now based on mostly qualitative considerations.," The motivation for our present analysis is to quantitatively investigate the suspected similarity (van der Klis, \cite{vdkapjsupp}) ) both in shape and strength of the noise component found in these two atoll sources and that in black hole candidates in the low state, which was up to now based on mostly qualitative considerations."
+ To this aim. we re-examine the EXOSAT observations of1705—44.. and perform. for the first time. a systematic. quantitative and direct comparison with aas well as with the black hole candidate (νο ΧΙ.," To this aim, we re-examine the EXOSAT observations of, and perform, for the first time, a systematic, quantitative and direct comparison with as well as with the black hole candidate Cyg X-1."
+ An additional reason to carefully look at these data again came from our earlier work on, An additional reason to carefully look at these data again came from our earlier work on
+ 18.5 8.9 122 22 1100 I 5 (1.7 6 , $18.5$ $8.9$ $122$ $22$ $\sim 1100$ $1$ $8$ $1.7$ $6$ 
+have shown that standard non-rotating single star moclels do not agree with these observed rates and that mass loss must be enhanced. in their case by rotation. to find an agreement between predicted SN ratios and those inferred [rom observations.,"have shown that standard non-rotating single star models do not agree with these observed rates and that mass loss must be enhanced, in their case by rotation, to find an agreement between predicted SN ratios and those inferred from observations."
+ The studies mentioned so far only compare single star models to the observations., The studies mentioned so far only compare single star models to the observations.
+ This ignores the observed. [act that many stars are in binaries., This ignores the observed fact that many stars are in binaries.
+ Binary evolution has been studied widely. (c.g. Pocsiacdlowski.Joss&Lsu1992:Well-&BelkusRamirez-Ruiz2007:VAjzquezctal.DrayPout 2007)).," Binary evolution has been studied widely, (e.g. \citealt{podsibin1,langerbin,belcz,izzysne,vanbevwr,vaz07,DT07}) )."
+ The methods employed to estimate the cllect of binarics is based either on rapid population synthesis or a small number of detailed models., The methods employed to estimate the effect of binaries is based either on rapid population synthesis or a small number of detailed models.
+ Population synthesis employs formulae or tables based on detailed models to follow the evolution. (e.g. Liurley.Pols&Tout2000. and the evolution of a star can be caleulated in a fraction of a second rather than many minutes taken by a detailed code.," Population synthesis employs formulae or tables based on detailed models to follow the evolution, (e.g. \citealt{HPT00} and the evolution of a star can be calculated in a fraction of a second rather than many minutes taken by a detailed code."
+ The speed of caleulations means that the full range of parameters that govern the outcome of binary evolution can be studied and their relative ellects evaluated., The speed of calculations means that the full range of parameters that govern the outcome of binary evolution can be studied and their relative effects evaluated.
+ The disadvantage is that in the complex phases of evolution. such as when the envelope is close to being stripped. the results of population svnthesis can be quite spurious.," The disadvantage is that in the complex phases of evolution, such as when the envelope is close to being stripped, the results of population synthesis can be quite spurious."
+ We have constructed a set. of binary stellar. mocels covering a wide range of binary parameter space using a detailed: stellar evolution code (IEldridge&Tout.2004a)., We have constructed a set of binary stellar models covering a wide range of binary parameter space using a detailed stellar evolution code \citep{E03}.
+. With these models we have investigated. the effect. binary evolution has on the lifetimes of the various phases of stellar evolution and therefore the cllect it has on the relative populations of massive stars and the relative numbers of different types of SNe., With these models we have investigated the effect binary evolution has on the lifetimes of the various phases of stellar evolution and therefore the effect it has on the relative populations of massive stars and the relative numbers of different types of SNe.
+ While the binary parametor-space resolution is low compared to rapid. population synthesis studies it is still one of the highest. resolution studies unelertaken with a detailed evolution codo., While the binary parameter-space resolution is low compared to rapid population synthesis studies it is still one of the highest resolution studies undertaken with a detailed evolution code.
+ We note that we co not consider. here the elleet of rotation or magnetic fields., We note that we do not consider here the effect of rotation or magnetic fields.
+" ""Phe elfect of rotation on the structure and evolution of stars is a complicated: process and has been studied by Heger.Langer&Woosley(2000) and Leger&Langer(2000) and the extensive study beginning in Mlevnet&Maecder(1997). ancl continuing up to most recently in Alevnet&Maecer(2007).", The effect of rotation on the structure and evolution of stars is a complicated process and has been studied by \citet{hegrot1} and \citet{hegrot2} and the extensive study beginning in \citet{mm1997} and continuing up to most recently in \citet{mmgrb}.
+.. Phe elect o£ magnetic fields on stellar structure ancl evolution has also been investigated by Alevnet&Maeder(2003).. Alevnet&AMaeder (2004).. Leger.Woosley&Spruit (2005).," The effect of magnetic fields on stellar structure and evolution has also been investigated by \citet{mm2003}, \citet{mm2004}, \citet{hegermag}."
+. Our code does not include rotation therefore our study. reveals the importance of cluplicity alone., Our code does not include rotation therefore our study reveals the importance of duplicity alone.
+ We discuss how rotation may alfect our results in Section 6., We discuss how rotation may affect our results in Section 6.
+ In this paper we start by providing details of our stellar model construction., In this paper we start by providing details of our stellar model construction.
+ We then specify our definitions of the dilferent. phases of stellar evolution and discuss how binary evolution alfects. the relative lifetimes of stars in. various phases of stellar evolution., We then specify our definitions of the different phases of stellar evolution and discuss how binary evolution affects the relative lifetimes of stars in various phases of stellar evolution.
+ Using these lifetimes we predict how the BSG to RSCG. the WR to O supergiant. the RSG to Wlt and. WC to WN ratios vary with metallicitv.," Using these lifetimes we predict how the BSG to RSG, the WR to O supergiant, the RSG to WR and WC to WN ratios vary with metallicity."
+ We then predict how the relative rates of twpe Ib/c and type LL SNe vary with metallicity., We then predict how the relative rates of type Ib/c and type II SNe vary with metallicity.
+ With each of these predictions we compare the values to the ratios inferred from observations., With each of these predictions we compare the values to the ratios inferred from observations.
+ Finally we discuss the implications of our results., Finally we discuss the implications of our results.
+ We use a detailed stellar evolution code to construct sets of single and binary star models rather than a rapid population svnthesis code. because we require accurate treatment of complex stages of evolution. such as when the envelope is close to being stripped. to obtain more accurate pre-SN models.," We use a detailed stellar evolution code to construct sets of single and binary star models rather than a rapid population synthesis code, because we require accurate treatment of complex stages of evolution, such as when the envelope is close to being stripped, to obtain more accurate pre-SN models."
+ We model the binary parameters with a reasonable resolution owing to the great amount of computer power that is now available., We model the binary parameters with a reasonable resolution owing to the great amount of computer power that is now available.
+ We have compared our binary models to the predictions of a rapid. population synthesis. code (lzzard&Tout2003:lzzarcd.Ramirez-Ruiz2004:Lz-gare&Tout2005) and found reasonable agreement for most of our predictions.," We have compared our binary models to the predictions of a rapid population synthesis code \citep{izzy,izzysne,izzynew} and found reasonable agreement for most of our predictions."
+ However clilliculty in assigning WR. types to population svnthesis helium stars meant the agreement was poor for the WC7WN ratio., However difficulty in assigning WR types to population synthesis helium stars meant the agreement was poor for the WC/WN ratio.
+ Our detailed. stellar models. were calculatedwith the Cambridee stellar evolution code. STARS. originally. developed by Eeeleton(LOT1) ancl updated: most recently by Polsetal.(1995). and Eldridge&Tout(2004a).," Our detailed stellar models were calculatedwith the Cambridge stellar evolution code, STARS, originally developed by \citet{E71} and updated most recently by \citet{P95} and \citet{E03}."
+. Further details can be found at the code's webpages?., Further details can be found at the code's web.
+. Ourmodels are available from the same location for download. without restriction., Ourmodels are available from the same location for download without restriction.
+ Phey are similar to those described by Eldridge&Tout(2004b) but we use 46 rather than 21 zero-age main-sequence models., They are similar to those described by \citet{ETsne} but we use 46 rather than 21 zero-age main-sequence models.
+ Every integer massfrom 5 to 40M. is modelled along with 45. 50. 55. 60. 70. SO. LOO and 120M..," Every integer massfrom $5$ to $40{\rm M}_{\odot}$ is modelled along with 45, 50, 55, 60, 70, 80, 100 and $120{\rm M}_{\odot}$."
+ The initially uniform. composition is VY=0.75—2.5Z and j—-035]L5Z. where NV is the mass fraction of hydrogen. Y that of helium ancl Z the initial metallicity taking the values 0.001. 0.004. 0.008. 0.02 and. 0.04.," The initially uniform composition is $X=0.75-2.5Z$ and $Y=0.25+1.5Z$, where $X$ is the mass fraction of hydrogen, $Y$ that of helium and $Z$ the initial metallicity taking the values 0.001, 0.004, 0.008, 0.02 and 0.04."
+ The composition mixture is scaled solar and Z.= 0.02., The composition mixture is scaled solar and ${\rm Z}_{\odot}=0.02$ .
+ Our mocdels end after core carbon-burning or at the onset of core neon ignition., Our models end after core carbon-burning or at the onset of core neon ignition.
+ During these burning stages. the envelope is only allected to a small degree because FueSEpesas," During these burning stages, the envelope is only affected to a small degree because $t_{\rm nuclear} \ll t_{\rm thermal}$."
+ Furthermore the late stages of evolution are rapid and have a negligible elfect on the final observable state of the progenitor. (Woosley.lleger&Weaver 2002)., Furthermore the late stages of evolution are rapid and have a negligible effect on the final observable state of the progenitor \citep{WHW02}.
+. Our opacity tables include the latest low temperature molecular opacities of Fergusonctal.(2005)., Our opacity tables include the latest low temperature molecular opacities of \citet{FA05}.
+. Phis update leads to tiny changes in the radii ancl surface teniperatures of red supergiant models because the altered opacity values only slightly modify the surface boundary conditions., This update leads to tiny changes in the radii and surface temperatures of red supergiant models because the altered opacity values only slightly modify the surface boundary conditions.
+ All the models employ our standard mass-Ioss prescription because it agrees best with observations— of SN progenitors (Ildridge&Tout 2004b)., All the models employ our standard mass-loss prescription because it agrees best with observations of SN progenitors \citep{ETsne}. .
+. The prescription is based. upon that of Dray&Tout.(2003). but. modified. in several wavs., The prescription is based upon that of \citet{DT03} but modified in several ways.
+ For pre-WI mass loss we use the rates of deexcept. for OB stars for which we use the theoretical rates οἱ Vink.de (2001).., For pre-WR mass loss we use the rates of \citet{dJ} except for OB stars for which we use the theoretical rates of \citet{VKL2001}. .
+ When the star becomes a WRstar CXuec< 0.4. οσοfs)2 4.0) we use the rates of Nugis (2000)..," When the star becomes a WRstar $X_{\rm surface}<0.4$ , $\log (T_{\rm eff}/{\rm K})>4.0$ ) we use the rates of \citet{NL00}. ."
+" We scale all rates with the standard factor GL)"" νάΚΙ.Pauldrach 2003).. except for the rates of Vink.deIxoter&which include their own metallicity scaling."," We scale all rates with the standard factor $(Z/$ $_{\odot})^{0.5}$ \citep{K87,H03}, , except for the rates of \citet{VKL2001} which include their own metallicity scaling."
+ , 
+" (3)), (6)) and (7)).","\ref{eq.3}) ), \ref{eq.6}) ) and \ref{eq.7}) )."
+ This comparison is presented in a form of scatter plot in Figure 7.., This comparison is presented in a form of scatter plot in Figure \ref{figure.7}.
+" The ordinate (vertical) axis represents o calculated from Equation (1)) with Wso calculated from Equations (3)), (6)) and (7)), where o, B and P were taken from Monte Carlo simulations, while the abscissa (horizontal) axis represents directly simulated values of o."," The ordinate (vertical) axis represents $\alpha$ calculated from Equation \ref{eq.1}) ) with $W_{50}$ calculated from Equations\ref{eq.3}) ), \ref{eq.6}) ) and \ref{eq.7}) ), where $\alpha$, $\beta$ and $P$ were taken from Monte Carlo simulations, while the abscissa (horizontal) axis represents directly simulated values of $\alpha$."
+ Red and green colours correspond to positive and negative values of 6 (for details of simulations and detection conditions see Appendix A))., Red and green colours correspond to positive and negative values of $\beta$ (for details of simulations and detection conditions see Appendix \ref{appendix.a}) ).
+" The blue line represents the maximum relative difference between scattered points (in the same scale as for o axes but expressed in percents), thus it describes an accuracy of determination of a."," The blue line represents the maximum relative difference between scattered points (in the same scale as for $\alpha$ axes but expressed in percents), thus it describes an accuracy of determination of $\alpha$."
+" As one can see, for almost orthogonal rotators the difference between the two methods is about 10 percent."," As one can see, for almost orthogonal rotators the difference between the two methods is about 10 percent."
+" However, for almost aligned rotators the error can reach or even exceed 100 percent."," However, for almost aligned rotators the error can reach or even exceed 100 percent."
+ For intermediate values of α the accuracy is in the range of 20 — 30 percent., For intermediate values of $\alpha$ the accuracy is in the range of 20 – 30 percent.
+" This means that Rankin's method of estimating the inclination angle from the width of the core component is quite good an approximation, provided that these values are not too small."," This means that Rankin's method of estimating the inclination angle from the width of the core component is quite good an approximation, provided that these values are not too small."
+" In particular, this method should not be applied to extremely broad profile pulsars for which one can expect very small inclination angles a<5°."," In particular, this method should not be applied to extremely broad profile pulsars for which one can expect very small inclination angles $\alpha\le5^{\circ}$."
+ The most important result of our paper is a clear confirmation of the lower limit in the distribution of half-power pulse-width Wso in core components of pulsar profiles when presented versus the pulsar period P using much bigger database than that used originally by R90.., The most important result of our paper is a clear confirmation of the lower limit in the distribution of half-power pulse–width $W_{50}$ in core components of pulsar profiles when presented versus the pulsar period $P$ using much bigger database than that used originally by \citet{r90}.
+" As illustrated in Figures 1, 2, 3 and 8 this limit is represented by the line Wso=2°.5P~°° at 1 GHz and Wso=29.37P-9? at 1.4 GHz."," As illustrated in Figures \ref{figure.1}, \ref{figure.2}, \ref{figure.3} and \ref{figure.8} this limit is represented by the line $W_{50}=2^{\circ}.5\: P^{-0.5}$ at 1 GHz and $W_{50}=2^{\circ}.37\: P^{-0.5}$ at 1.4 GHz."
+" It seems that more appropriate description of the lower boundary limit is the lower boundary belt (LBB) Wso—(25.375153)Ρ05, which accounts for scatter and errors in W5o measurements."," It seems that more appropriate description of the lower boundary limit is the lower boundary belt (LBB) $W_{50}=(2^{\circ}.37^{+1^{\circ}.43}_{-0^{\circ}.87})\: P^{-0.5}$, which accounts for scatter and errors in $W_{50}$ measurements."
+" Strictly speaking, the widths of the LBB was determined by the scatter and errors of DP-IP cases marked as red symbols in Figure 3."," Strictly speaking, the widths of the LBB was determined by the scatter and errors of DP-IP cases marked as red symbols in Figure 3."
+" However, the normal data points seem to populate the belt marked in Figure 8 in similar manner as the DP-IP cases do in Figure 3."," However, the normal data points seem to populate the belt marked in Figure 8 in similar manner as the DP-IP cases do in Figure 3."
+" The concept of LBB is based not only on core widths of DP-IP pulsars (like in R90)) but here it is based on all Wso measurements of normal pulsars available in the ATNF database, without distinctions into core and conal components, of course."," The concept of LBB is based not only on core widths of DP–IP pulsars (like in \citet{r90}) ) but here it is based on all $W_{50}$ measurements of normal pulsars available in the ATNF database, without distinctions into core and conal components, of course."
+ This is dramatically illustrated in Figure 8 containing 1495 data points., This is dramatically illustrated in Figure \ref{figure.8} containing 1495 data points.
+ This figure is a reproduction of Figure 3 with all special points described in the legend omitted., This figure is a reproduction of Figure \ref{figure.3} with all special points described in the legend omitted.
+" Apparently, the existence of the lower boundary limit in the pulse-width measurements discovered by Rankin (1990) does not depend neither on the method nor on the errors of measurements of the core widths in DP-IP pulsars."," Apparently, the existence of the lower boundary limit in the pulse–width measurements discovered by Rankin (1990) does not depend neither on the method nor on the errors of measurements of the core widths in DP-IP pulsars."
+ The unweighted formal fit to the points within LBB (between dashed lines) is Wiso=(2°.86c0°.07)P~0-5140.02, The unweighted formal fit to the points within LBB (between dashed lines) is $W_{50}=(2^{\circ}.86\pm0^{\circ}.07)\: P^{-0.51\pm0.02}$.
+ The existence of the lower bound in the distribution of Wso versus pulsar period P is very intriguing and requires solid and convincing explanation., The existence of the lower bound in the distribution of $W_{50}$ versus pulsar period $P$ is very intriguing and requires solid and convincing explanation.
+" The exponent —0.5 is evidently related to the geometry of dipolar field lines in the radio emission region, while the value of the numerical factor about 2°.5 is a little more mysterious."," The exponent $-0.5$ is evidently related to the geometry of dipolar field lines in the radio emission region, while the value of the numerical factor about $2^{\circ}.5$ is a little more mysterious."
+" R90 argued that the core component widths are intimately related to the polar cap geometry at the stellar surface, whose bundle of the last open dipolar field lines have an angular extent (opening angle 2p) very close to 2°.45 P~/?."," \citet{r90} argued that the core component widths are intimately related to the polar cap geometry at the stellar surface, whose bundle of the last open dipolar field lines have an angular extent (opening angle $2\rho$ ) very close to $2^{\circ}.45 \: P^{-1/2}$ ."
+" Although this natural and very appealing explanation is commonly accepted, it is not free from theoretical and interpretational problems."," Although this natural and very appealing explanation is commonly accepted, it is not free from theoretical and interpretational problems."
+" First of all, the coherent radio emission cannot originate at or very near the polar cap, because there are"," First of all, the coherent radio emission cannot originate at or very near the polar cap, because there are"
+uodel of the physical situation. thus they fail to xoperlv account for the actual inaccuracy with respect to the seueral relativistic theory.,"model of the physical situation, thus they fail to properly account for the actual inaccuracy with respect to the general relativistic theory."
+ The size of the truncation error bv itself within the lin-lens scheme has little meaning if the actual error introduced by the thin-lens assumption is iot known., The size of the truncation error by itself within the thin-lens scheme has little meaning if the actual error introduced by the thin-lens assumption is not known.
+ Tn order to give an idea of the sizes of both OCYTOYTS (thin-leus aud lack of truncation) with respect to the relativistic theory. we also compare our inteeration of the null eeodesies with thiu-lens models truncated along the line of sight.," In order to give an idea of the sizes of both errors (thin-lens and lack of truncation) with respect to the relativistic theory, we also compare our integration of the null geodesics with thin-lens models truncated along the line of sight."
+ In this case. the inaccuracy of the mocel is eutirely due to a single assumption: that of thin lenses.," In this case, the inaccuracy of the model is entirely due to a single assumption: that of thin lenses."
+ Our moechanisin of triuucation is as simple as possible: a liuc eut off ofthe mass density after an arbitrarily chosen radius., Our mechanism of truncation is as simple as possible: a hard cut off of the mass density after an arbitrarily chosen radius.
+ Iu Section 2.. we outline the wavs the Rohertsou-Walker (RW) backeround ecometiy enters our calculation. aud in Section 3.. we outline the equationsof motion.," In Section \ref{back:sec}, we outline the ways the Robertson-Walker (RW) background geometry enters our calculation, and in Section \ref{eqns:sec}, we outline the equations of motion."
+ The SIS aud NEW eravitational potentials and iu-leus models are developed iu section L., The SIS and NFW gravitational potentials and thin-lens models are developed in section \ref{models:sec}.
+ Section 5 describes our experimental set-up. nmnerical approach. and Ίσα. accuracy.," Section \ref{sim:sec} describes our experimental set-up, numerical approach, and numerical accuracy."
+ Qur nuiadu results comparing the accuracv of he thin-lens models are presented iu section 7.., Our main results comparing the accuracy of the thin-lens models are presented in section \ref{compare:sec}.
+ Section 8 preseuts the variation of thin-lens accuracy as a function of cosmologv parameters or a range of flat cosmologies.," Section \ref{cosmo:sec}
+ presents the variation of thin-lens accuracy as a function of cosmology parameters for a range of flat cosmologies."
+ We discuss and explain our results iu section 9 before couchiding with some conuueuts about our results and the heir duplications for lensing studies in section 10.., We discuss and explain our results in section \ref{discussion:sec} before concluding with some comments about our results and the their implications for lensing studies in section \ref{conclusion:sec}.
+" We take as the backeround a fiat RW. metric with the eurreut accepted values of the IHTubble constant [fy=10 km/s/Mpe. the matter density Q,,=0.3 and the cosmological coustaut density QO,=O07=1Q,,."," We take as the background a flat RW metric with the current accepted values of the Hubble constant $H_0 = 70$ km/s/Mpc, the matter density $\Omega_m = 0.3$ and the cosmological constant density $\Omega_\Lambda
+= 0.7 = 1- \Omega_m$."
+" The metric is with a scale factor e(f) for our cosimological model as in Ryden(2003) The Hubble parameter Z7=àfa is We will be cousiceriug photous enmütted by a source at enüssion fine f. aud arriving at the observer at the observation time ¢,,."," The metric is with a scale factor $a(t)$ for our cosmological model as in \cite{ryden}
+ The Hubble parameter $H\equiv \dot{a}/a$ is We will be considering photons emitted by a source at emission time $t_e$, and arriving at the observer at the observation time $t_o$ ."
+" For unuuerical convenience. we fix the scale factor to the value 1 at the observation time. a(t)=1. which vields t,=L218«lots."," For numerical convenience, we fix the scale factor to the value 1 at the observation time, $a(t_o) = 1$, which yields $t_o = 4.248 \times
+10^{17}\,s $."
+ A natural wav to give the relative locations between the lens. observer. aud source is to place the lens at the spatial origin. aud to use the redshifts to determine the positioning of the lens and observer.," A natural way to give the relative locations between the lens, observer, and source is to place the lens at the spatial origin, and to use the redshifts to determine the positioning of the lens and observer."
+" Using the redshift relation. 1att)""ITE and setting e(f,)=l. we can solve for the value of t. with Eq. 2.."," Using the redshift relation, $ 1+\tilde z =
+\frac{a(t_o)}{a(t_e)}, $ and setting $a(t_o)=1$, we can solve for the value of $t_e$ with Eq. \ref{a1}."
+ To obtain the radial positious of the source aux observer. we orient our coordinates in such a wav that a ight rav travels racially iu the backerouxk spacetime and assume that the observer. leus iux source are at least nearly co-lnear.," To obtain the radial positions of the source and observer, we orient our coordinates in such a way that a light ray travels radially in the background spacetime and assume that the observer, lens and source are at least nearly co-linear."
+ We can then integrate racial null ecoclesics of the metric Eq., We can then integrate radial null geodesics of the metric Eq.
+ 1 to determine the radial positions of the source au observer (by ignoringthe perturbation introduce by the cus}., \ref{FRW1} to determine the radial positions of the source and observer (by ignoringthe perturbation introduced by the lens).
+" The leusing scenario is described in terms of the following weakly perturbed RW metric This metric is accurate to first order iu y. which is a Newtouian potential determined by the proper mass clensity of tle Ileus py via where the derivatives in the Laplacian operator are taken with respect to the proper distance ry,= atte."," The lensing scenario is described in terms of the following weakly perturbed RW metric This metric is accurate to first order in $\varphi$, which is a Newtonian potential determined by the proper mass density of the lens $\rho_p$ via where the derivatives in the Laplacian operator are taken with respect to the proper distance $r_p = a(t) r$ ."
+" Commonly. the proper mass density is taken as à function p,(r,} which depeuds explicitly on the proper distance +, but not ou thotime f."," Commonly, the proper mass density is taken as a function $\rho_p(r_p)$ which depends explicitly on the proper distance $r_p$ but not on thetime $t$ ."
+ Under conventional lensing conditions. a lightrav speuds a verv sinall amount of time iu the area of iuflucuce of the leus (where ο= 0.," Under conventional lensing conditions, a lightray spends a very small amount of time in the area of influence of the lens (where $\varphi\neq 0$ )."
+The sources are saturated at A « 12.,The sources are saturated at $K$ $<$ 12.
+ For such bright sources. 2\LASS Point Source Catalog cata were used.," For such bright sources, 2MASS Point Source Catalog data were used."
+ We also extracted source magnitudes from the UINIDSS images using the methods described above. although our UFTI images (rather than 2MASS data) were used to calibrate the UIXIDSS data.," We also extracted source magnitudes from the UKIDSS images using the methods described above, although our UFTI images (rather than 2MASS data) were used to calibrate the UKIDSS data."
+ The UIXIDSS observations were used to create a density map of the region (see Figure 2). and to obtain photometry over a larger area (han was observed with UFTI.," The UKIDSS observations were used to create a density map of the region (see Figure 2), and to obtain photometry over a larger area than was observed with UFTI."
+ NIB. HIN band spectra of several stars in the (target field were obtained using the UIXIRT 1-5 yan Imager Spectrometer (UIST) (RamsayHowatetal.2004). on the night of 2p July 2005., NIR $HK$ band spectra of several stars in the target field were obtained using the UKIRT 1 - 5 $\mu$ m Imager Spectrometer (UIST) \citep{rh04} on the night of $^{st}$ July 2005.
+ UIST has a 1024 x 1024 InSb array and a (0.12 arcsec per pixel plate scale., UIST has a 1024 $\times$ 1024 InSb array and a 0.12 arcsec per pixel plate scale.
+ The observations were made using a 120 arcsec long slit wilh a width of 4 pixels., The observations were made using a 120 arcsec long slit with a width of 4 pixels.
+ The [iv erism was used. which allowedcomplete wavelength coverage from 1.395 jan to 2.506 jun wilh a spectral resolution of Ro c 500.," The $HK$ grism was used, which allowedcomplete wavelength coverage from 1.395 $\mu$ m to 2.506 $\mu$ m with a spectral resolution of R $\simeq$ 500."
+ The exposure time per [rame was set to 60 sec; 5 exposures were obtained per slit position. resulting in a total integration time of 8 minutes.," The exposure time per frame was set to 60 sec; 8 exposures were obtained per slit position, resulting in a total integration time of 8 minutes."
+ The long slit was aligned carefully to cover multiple sources in a single position., The long slit was aligned carefully to cover multiple sources in a single position.
+ This configuration and integration time provided a 5e point source sensitivity of  16 mag in the ILI bands., This configuration and integration time provided a $\sigma$ point source sensitivity of $\sim$ 16 mag in the $HK$ bands.
+ The target was nodded up aud down the slit in an ABBA fashion: B exposures were subtracted [from adjacent A exposures to give skv-subtractecl 2-dimensional spectral images., The target was nodded up and down the slit in an ABBA fashion; B exposures were subtracted from adjacent A exposures to give sky-subtracted 2-dimensional spectral images.
+ The spectral images were averaged ancl the positive and negative beams of stellar speclra were (hen optimallv exiracted., The spectral images were averaged and the positive and negative beams of stellar spectra were then optimally extracted.
+" These resulting ""group"" spectra were corrected Lor tellure absorption and flux calibrated by division with a similarlv-observed. stanclard star spectrum.", These resulting “group” spectra were corrected for telluric absorption and flux calibrated by division with a similarly-observed standard star spectrum.
+ The G2V source HIPLO2189 was used for this purpose. (, The G2V source HIP102189 was used for this purpose. (
+Nolte that the stancarel star spectrum was first. divided bv a black body template of appropriate temperature {ο preserve (he intrinsic slopes of the calibrated target spectra.),Note that the standard star spectrum was first divided by a black body template of appropriate temperature to preserve the intrinsic slopes of the calibrated target spectra.)
+ The initial reduction steps (ie. dark ancl flat correction) were performed. using the UIXIRT/Starlink data reduction pipeline ORAC-DR (Cavanaghetal.2003): the extraction and analvsis of the spectra were done with ΗΑΕ and IDL., The initial reduction steps (i.e. dark and flat correction) were performed using the UKIRT/Starlink data reduction pipeline ORAC-DR \citep{cav08}; the extraction and analysis of the spectra were done with IRAF and IDL.
+ BessellPVR images of the region. associated. with the IRAS 19343+2026 region were obtained with the 2-m Himalavan Telescope (ICT) of the Indian Astronomical, Bessell images of the region associated with the IRAS 19343+2026 region were obtained with the 2-m Himalayan Telescope (HCT) of the Indian Astronomical
+confused by the asymmetric surface brightness typical of late-type stars (2).. especially so for a star as extended as L7? Pup having a (predicted) apparent angular radius of 9.6 mas (2)..,"confused by the asymmetric surface brightness typical of late-type stars \citep{Ragland-2006}, especially so for a star as extended as $^2$ Pup having a (predicted) apparent angular radius of 9.6 mas \citep{DummSchild98}."
+ A further argument the binary nature of L7 Pup is the detection of all three masers SiO. H»O and OH (?).. since a binary companion. when orbiting in the layers where the maser activity should originate. prevents ifs operation (222).," A further argument the binary nature of $^2$ Pup is the detection of all three masers SiO, $_2$ O and OH \citep{Jura-2002}, since a binary companion, when orbiting in the layers where the maser activity should originate, prevents its operation \citep{HermanHabing85,Lewis-1987,Schwarz-1995}."
+ax=0.9 for the linear mass fluctuation in a sphere of radius Sh! Mpe extrapolated to z=0. and n=1 for the slope of the primordial Huctuation spectrum.,"$\sigma_8=0.9$ for the linear mass fluctuation in a sphere of radius $8 h^{-1}$ Mpc extrapolated to $z=0$, and $n=1$ for the slope of the primordial fluctuation spectrum."
+ Ehe simulation ollowed 21607 dark matter particles from 2=127 to the oesent-day within a cubic region 500f!Mpe on a side., The simulation followed $2160^3$ dark matter particles from $z=127$ to the present-day within a cubic region $500 h^{-1}$ Mpc on a side.
+ ‘The individual particle mass is thus S.6«107.4M.. and he gravitational force had a Plummoer-equivalent comoving softening of 5. tkpe.," The individual particle mass is thus $8.6\times 10^{8}h^{-1}{\rm M_\odot}$, and the gravitational force had a Plummer-equivalent comoving softening of $5h^{-1}$ kpc."
+ Initial conditions were set using the Boltzmann code (Seljak Zalcarriaga 1996) to generate a realisation of the desired. power spectrum which was then imposed on a glass-like uniform particle load (White 1996)., Initial conditions were set using the Boltzmann code (Seljak Zaldarriaga 1996) to generate a realisation of the desired power spectrum which was then imposed on a glass-like uniform particle load (White 1996).
+ The N-body codeGADGLEL2 (Springel 2005) was used to carry out the simulation anc the full data were stored at 64 times spaced approximately equally in the logarithm of the expansion factor., The N-body code (Springel 2005) was used to carry out the simulation and the full data were stored at $64$ times spaced approximately equally in the logarithm of the expansion factor.
+ This allowed us to build trees which store detailed assembly histories for cach of the 5.7 million dark matter haloes at z=0 that contain at least 64 particles according to a group finder with b=0.2 (Davis ct al 1985)., This allowed us to build trees which store detailed assembly histories for each of the $5.7$ million dark matter haloes at $z=0$ that contain at least $64$ particles according to a group finder with $b=0.2$ (Davis et al 1985).
+ Vhis is the set of haloes we analyse in the remainder of this paper., This is the set of haloes we analyse in the remainder of this paper.
+ Dark matter haloes are biased tracers of the unclerlving mass density field. and models for the strength ofthis bias can be constructed using the excursion set formalism of Bond et al (1991) and Lacey Cole (1993).," Dark matter haloes are biased tracers of the underlying mass density field, and models for the strength of this bias can be constructed using the excursion set formalism of Bond et al (1991) and Lacey Cole (1993)."
+ Within this framework Mo White (1996) and Sheth. Mo Vormen (2000) derived analytic expressions for halo bias ancl tested them against :-body simulations ofa variety of cosmogonies (see also Cole Ixaiser (1989)).," Within this framework Mo White (1996) and Sheth, Mo Tormen (2000) derived analytic expressions for halo bias and tested them against N-body simulations of a variety of cosmogonies (see also Cole Kaiser (1989))."
+ The linear density field can be represented xw oCM.zz). the linear mass Iluctuation (extrapolated to redshift 2) within a sphere which on average contains mass AM.," The linear density field can be represented by $\sigma(M, z)$, the linear mass fluctuation (extrapolated to redshift $z$ ) within a sphere which on average contains mass $M$."
+ A characteristic mass for clustering Al.(2) can then be defined through e(M.;.z)=ὃς21.69.," A characteristic mass for clustering $M_*(z)$ can then be defined through $\sigma(M_*, z) = \delta_c\approx 1.69$."
+" Haloes more massive than M, are predicted: to be positively biased. (i.c. more stronely clustered) relative to the underlying mass. while the opposite is true for less massive haloes."," Haloes more massive than $M_*$ are predicted to be positively biased (i.e. more strongly clustered) relative to the underlying mass, while the opposite is true for less massive haloes."
+" The tests in Mo White (1996) showed the autocorrelation function for dark matter haloes of mass ÀJ to be approximately parallel to that of the mass. where the “bias factor” ος.z) is given by Llere 5, is the critical linear overdensity at collapse and depends slightly on cosmology: we use values from Eke et al. ("," The tests in Mo White (1996) showed the autocorrelation function for dark matter haloes of mass $M$ to be approximately parallel to that of the mass, where the “bias factor” $b(\nu,z)$ is given by Here $\delta_c$ is the critical linear overdensity at collapse and depends slightly on cosmology; we use values from Eke et al. ("
+1996).,1996).
+ 720.f/o(M.z) is the dimensionless amplitude of Iluctuations that produce haloes of mass M at redshift z.," $\nu=\delta_c/\sigma(M,z)$ is the dimensionless amplitude of fluctuations that produce haloes of mass $M$ at redshift $z$."
+ Reeent high. resolution AN-bocly simulations have qualitatively confirmed this model and the improved version of Sheth. Mo Tormen (2000). but quantitative fits show some deviations (Jing 1998: Governato et al 1998: Colberg et al 2000: Ixravtsov IxIvpin 1999: Seljak Warren 2004: Alandelbaum et al.," Recent high resolution $N$ -body simulations have qualitatively confirmed this model and the improved version of Sheth, Mo Tormen (2000), but quantitative fits show some deviations (Jing 1998; Governato et al 1998; Colberg et al 2000; Kravtsov Klypin 1999; Seljak Warren 2004; Mandelbaum et al."
+ 2005)., 2005).
+ Lere we use the unprecedented dynamic range and statistics of the Millennium Simulation to study this bias relation further., Here we use the unprecedented dynamic range and statistics of the Millennium Simulation to study this bias relation further.
+ For haloes of given mass and redshift. we derive a bias [actor bv comparing their two point correlation functionto that of the mass.," For haloes of given mass and redshift, we derive a bias factor by comparing their two point correlation functionto that of the mass."
+ More specifically. we estimate b(M.z)as the relative normalisation factor which minimizes the mean square cdillerence in log£ for four equal width bins in logr spanning the separation range Gh.IMpe<r<255.!Mpe.," More specifically, we estimate $b^2(M, z)$ as the relative normalisation factor which minimizes the mean square difference in $\log \xi$ for four equal width bins in $\log r$ spanning the separation range $6h^{-1}{\rm Mpc} < r < 25 h^{-1}{\rm Mpc}$."
+ In this range the measured correlations are all in the quasilincar regime., In this range the measured correlations are all in the quasilinear regime.
+ Since b(A.z) is a steep function of AL for Al»M. itis important to select halo samples in relatively narrow mass ranges in order to determine 6 accurately at high mass.," Since $b(M, z)$ is a steep function of $M$ for $M>M_*$, it is important to select halo samples in relatively narrow mass ranges in order to determine $b$ accurately at high mass."
+ Here we choose mass bins of width M=0.3M.," Here we choose mass bins of width $\Delta
+M=0.3M$."
+ The measured bias factors for the Millennium. Simulation are shown by the symbols in Fig. 1.., The measured bias factors for the Millennium Simulation are shown by the symbols in Fig. \ref{fig:bias}.
+ Note that we plot bias as a function of peak height in order that the predicted relations [for dillerent redshifts coincide., Note that we plot bias as a function of peak height in order that the predicted relations for different redshifts coincide.
+ Dillerent. svibols denote the measured. bias factors for. dilferent: redshitts., Different symbols denote the measured bias factors for different redshifts.
+ For comparison. we overplot the analytic expressions which other authors have derived. theoretically or given as fits to wir own numerical data.," For comparison, we overplot the analytic expressions which other authors have derived theoretically or given as fits to their own numerical data."
+ The first noticeable feature is that when plotted in this wav b(7) depends at most weakly on redshift: at. fixed. p. the bias factors for clifferent redshift are almost identical," The first noticeable feature is that when plotted in this way $b(\nu)$ depends at most weakly on redshift; at fixed $\nu$ , the bias factors for different redshift are almost identical"
+find that all R inodes that satisty ΤCr)=Πές)) precess in the prograde direction with frequency w= 1.5.,"find that all R modes that satisfy $W'(x_{\rm i})=W'(x_{\rm o})=0$ precess in the prograde direction with frequency $\omega =
+0.5$ ."
+ The erowth rates of the four low-order R inodes shown in Figure 1 are a=9.19 for η=IL. which is the owest-order R mode: ¢=5.16 for η=3: 0=0.58 Or)=7: aud a=72 for vn=12.," The growth rates of the four low-order R modes shown in Figure 1 are $\sigma = 9.19$ for $n=1$, which is the lowest-order R mode; $\sigma
+= 5.16$ for $n=3$; $\sigma = 0.58$ for $n=7$; and $\sigma = -4.72$ for $n=12$."
+ The erowth rate falls steeply with increasing mode muuber. aud nodes with »zLO are very strouglv damped.," The growth rate falls steeply with increasing mode number, and modes with $n \gta 10$ are very strongly damped."
+" For fixed mode nuuber ». the erowth rate increases with increasing wy. Lo. dofd,>0."," For fixed mode number $n$, the growth rate increases with increasing $x_{\rm o}$, i.e., $d\sigma/dx_{\rm o} > 0$."
+ We also find that all the Ro aodes that satisfv Woe)ο=0 precess in the prograde direction with frequency w=l., We also find that all the R modes that satisfy $W(x_{\rm i})=W(x_{\rm o})=0$ precess in the prograde direction with frequency $\omega = 1$.
+ This is the same precessio- frequency fotnd by Maloney et ((1996) for all the Ro anodes that satisfy 471/60)>0 as 5Q dustead of μα)=0., This is the same precession frequency found by Maloney et (1996) for all the R modes that satisfy $x^2 W'(x) \rightarrow 0$ as $x\rightarrow 0$ instead of $W'(x_{\rm i}) = 0$.
+ This degeneracy of R-mode precession frequencies is peculiar to isothermal disks: for more general temperature profiles. different B modes precess at different rates (sec Malonev ct 11998).," This degeneracy of R-mode precession frequencies is peculiar to isothermal disks: for more general temperature profiles, different R modes precess at different rates (see Maloney et 1998)."
+ On the other haud. for R modes that satisfy “unixed™ boundary conditions. cither H(5)=Wor)= 00r H'CGr)=Wr)0. the higher-order Loes precess at slelthy different frequencies from he lowest-orcey modes.," On the other hand, for R modes that satisfy “mixed” boundary conditions, either $W(x_{\rm i})=W'(x_{\rm o})=0$ or $W'(x_{\rm
+i})=W(x_{\rm o})=0$, the higher-order modes precess at slightly different frequencies from the lowest-order modes."
+" For Wey)=W'Ge;)0. wos modes precess iu the progradoe direction. but he orque acting on the inner edee of the disk causes the ilt uuction of oue of the low-order modes to increase uonotonicallv with increasing we. 1.6.. LlCr) is positive everywhere,"," For $W(x_{\rm
+i})=W'(x_{\rm o})=0$, most modes precess in the prograde direction, but the torque acting on the inner edge of the disk causes the tilt function of one of the low-order modes to increase monotonically with increasing $x$, i.e., $\beta'(x)$ is positive everywhere."
+ This mode precesses very slowly. either in the prograde or the retrograde seuse. depeudiug ou he value of ος aud is either weakly growing or weakly damped (see Alarkovie Lah. in preparation).," This mode precesses very slowly, either in the prograde or the retrograde sense, depending on the value of $x_{\rm o}$, and is either weakly growing or weakly damped (see Marković Lamb, in preparation)."
+ Consider now the low-frequency normal modes of the disk when both the gravitomaguctic torque aud the radiation warping torque are included., Consider now the low-frequency normal modes of the disk when both the gravitomagnetic torque and the radiation warping torque are included.
+ These are the LFCAIR odes., These are the LFGMR modes.
+ These modes can be found by inteerating tje warp equation (7)) outward from the inner edge o the disk. using the same approach that was used to find the LPCAL modes.," These modes can be found by integrating the warp equation \ref{WarpEqn}) ) outward from the inner edge of the disk, using the same approach that was used to find the LFGM modes."
+" The character of the LFCAIR modes changes with increasing mode μπασον,", The character of the LFGMR modes changes with increasing mode number.
+" Like the tilt functions ofthe low-order R uodes. the tilt functions of the low-order LEGAM uioles are appreciable ouly at large radii axd are affected x owlüch boundary condition is imposed at c, and by the value of wv."," Like the tilt functions of the low-order R modes, the tilt functions of the low-order LFGMR modes are appreciable only at large radii and are affected by which boundary condition is imposed at $x_{\rm o}$ and by the value of $x_{\rm o}$."
+" Like the precession frequencies o fall the R inodes. the precession frequeicles of the low-orider LFOAIR modes depend ou which boundary condition is posed at the outer boundarx. but not on the value of ντο,"," Like the precession frequencies of all the R modes, the precession frequencies of the low-order LFGMR modes depend on which boundary condition is imposed at the outer boundary, but not on the value of $x_{\rm
+o}$."
+ Uulike the R modes but like the LFCAL modes. the frequencies aud camping rates of the LFOAIR aoces are almost completely independent of the iuner boundary couditiou. for the sale reason that this is true of the LFCAL modes.," Unlike the R modes but like the LFGM modes, the frequencies and damping rates of the LFGMR modes are almost completely independent of the inner boundary condition, for the same reason that this is true of the LFGM modes."
+ The properties of the high-order LEGMB modes are ouly weakly dependent ou the boundary condition at ss and the value of we., The properties of the high-order LFGMR modes are only weakly dependent on the boundary condition at $x_{\rm o}$ and the value of $x_{\rm o}$.
+ Oulv. the properties of the owest-order LFCAIR nodes depend on the accretion efficiency €., Only the properties of the lowest-order LFGMR modes depend on the accretion efficiency $\epsilon$.
+ For the parameter values considered i his report. he first six LFCAIR nodes are appreciadle ouly at arge radii aud hence are virtually unaffected by the eravitoniaenetie torque.," For the parameter values considered in this report, the first six LFGMR modes are appreciable only at large radii and hence are virtually unaffected by the gravitomagnetic torque."
+ Their tilt fuuctious. precession ryequencies. aud erowthrates are essentially ideutical o those of the first six R modes (see Figs.," Their tilt functions, precession frequencies, and growth rates are essentially identical to those of the first six R modes (see Figs."
+ ] aud 2). except that the very Ποιος xut of he disk is orced iuto the rotatio1 equator of the central object w the combined action of the eravitomaguetic aud internal viscous torques. as first discussed by Bardecu Petterson (1975).," 1 and 2), except that the very innermost part of the disk is forced into the rotation equator of the central object by the combined action of the gravitomagnetic and internal viscous torques, as first discussed by Bardeen Petterson (1975)."
+ This is the reasou that the propertics of the lowes-order LFCAIR modes do not depend ou the boundary coucdition at the ΠΙΟΣ οςee of the disk., This is the reason that the properties of the lowest-order LFGMR modes do not depend on the boundary condition at the inner edge of the disk.
+ Like the low-order radiation warping modes. the lowest-order LFCOMB. modes precess in the prograde direction with precession frequencies of either 0.5 or 1.0. cpending ou whether the iuner aud outer edges of the disk are free or piuued iu the disk plane.," Like the low-order radiation warping modes, the lowest-order LFGMR modes precess in the prograde direction with precession frequencies of either 0.5 or 1.0, depending on whether the inner and outer edges of the disk are free or pinned in the disk plane."
+ For a compact object of solar mass. these precession Tequeneles are ~1| CIE.," For a compact object of solar mass, these precession frequencies are $\sim
+10^{-5}\,$ Hz."
+ Also like the low-order radiation warping modes. the lowest-order LFOAIR nodes are all growing modes.," Also like the low-order radiation warping modes, the lowest-order LFGMR modes are all growing modes."
+ There are typically two or three LFCAIR modeS hat are hwbrid eravitomaguctic raciatiou-warping nodes., There are typically two or three LFGMR modes that are hybrid gravitomagnetic radiation-warping modes.
+ The shapes of these modes are very simular to he radiation-warping modes in the outer disk. except hat the combined action of the eravitomaguetic aud internal viscous torques again forces the very innerniost vat of the disk iuo the rotation equator of f1C central object.," The shapes of these modes are very similar to the radiation-warping modes in the outer disk, except that the combined action of the gravitomagnetic and internal viscous torques again forces the very innermost part of the disk into the rotation equator of the central object."
+ The hybrid eravitomiagnueic racliation-warping modes generally (it not always: see below) xecess In the prograde direction wit1 precession LOCuencies between ~10.© Iz. for a coupact object of solar mass (~0.08 im our cinensionless mits). and he ~10? Uz frequeuev of the radialon-Wwarplus nodes.," The hybrid gravitomagnetic radiation-warping modes generally (but not always; see below) precess in the prograde direction with precession frequencies between $\sim10^{-6}$ Hz, for a compact object of solar mass $\sim$ 0.08 in our dimensionless units), and the $\sim
+10^{-5}$ Hz frequency of the radiation-warping modes."
+ These modes are eiher weakly damped or weakly erowius., These modes are either weakly damped or weakly growing.
+ Thus. when the radialon-Wwarplus orque is taken iuto account. there is generally no ince-inudepeudenut mode: the Bardeen-Petterson mode," Thus, when the radiation-warping torque is taken into account, there is generally no time-independent mode: the Bardeen-Petterson mode"
+hot halo or a sudden supply of cold. eas through a eas-rich merger event.,hot halo or a sudden supply of cold gas through a gas-rich merger event.
+ The former scenario is more likely in massive earlv-tvpe galaxies at the centres of clusters., The former scenario is more likely in massive early-type galaxies at the centres of clusters.
+ Field/LIow-nass early types are less likely to be fuelled by this mechanism. with eas-rich mergers feeding the black hole in GN hosts.," Field/low-mass early types are less likely to be fuelled by this mechanism, with gas-rich mergers feeding the black hole in AGN hosts."
+ Since in the local Universe such gas-rich merger events are much more rare than cooling episodes. cluster/high-mass earlv-tvpe galaxies are more likely to host radio AGN than their field/Iow-niass counterparts.," Since in the local Universe such gas-rich merger events are much more rare than cooling episodes, cluster/high-mass early-type galaxies are more likely to host radio AGN than their field/low-mass counterparts."
+ The dust lane early-type population. both in the field and in clusters. shows the same distribution of radio properties as the feld control population.," The dust lane early-type population, both in the field and in clusters, shows the same distribution of radio properties as the field control population."
+" ""TFherefore. if merecrs are the major trigeer of radio AGN activity in all Ποιά carly twpes. this is also the case for dust lane carly- galaxies. regardless of whether they are located in rich or poor environments."," Therefore, if mergers are the major trigger of radio AGN activity in all field early types, this is also the case for dust lane early-type galaxies, regardless of whether they are located in rich or poor environments."
+ In other words. it appears as though a dust lane feature serves as a proxy for a recent. gas-rich merecr.," In other words, it appears as though a dust lane feature serves as a proxy for a recent gas-rich merger."
+ This paradigm is supported by the finding in Paper I that almost all dust lance galaxies. exhibit: morphological disturbances and recent starbursts., This paradigm is supported by the finding in Paper I that almost all dust lane galaxies exhibit morphological disturbances and recent starbursts.
+ According to this scenario. all dust lane early. tvpes should have undergone a recent gas-rich merger (and hence presumably host a detectable AGN). while only some of the carly types in the control sample would have done the same.," According to this scenario, all dust lane early types should have undergone a recent gas-rich merger (and hence presumably host a detectable AGN), while only some of the early types in the control sample would have done the same."
+ Figs 5 and 6.. however. show that the racio-loud fractions for these objects are similar.," Figs \ref{fig:RLF_byEnvironment_dustlanesVscontrol} and \ref{fig:RLF_byEnvironment_separate}, however, show that the radio-loud fractions for these objects are similar."
+ This is most likely due to our conservative (1.50 above emission. due to star formation) racio luminosity cut adopted for GN identification., This is most likely due to our conservative $\sigma$ above emission due to star formation) radio luminosity cut adopted for AGN identification.
+ One way the above paradigm can be tested is by comparing radio and emission-line ACN activity in dust lane carly-vpe galaxies and the control sample., One way the above paradigm can be tested is by comparing radio and emission-line AGN activity in dust lane early-type galaxies and the control sample.
+ Lf gas-rich mergers are indeed. responsible for fuelling the αν in dust lane early vpes. one would expect a much higher fraction of radio AGN in these objects to also exhibit emission-line CIN activity. compared to objects in which gas cooling is the major AGN rigger.," If gas-rich mergers are indeed responsible for fuelling the AGN in dust lane early types, one would expect a much higher fraction of radio AGN in these objects to also exhibit emission-line AGN activity, compared to objects in which gas cooling is the major AGN trigger."
+ In Table 2. we present the fraction of radio ACGN josts Chat are also identified as emission-line ACN., In Table \ref{tab:fracOpticalAGN_dustlanes} we present the fraction of radio AGN hosts that are also identified as emission-line AGN.
+ Almost all objects identified as racio XCGN also show up as ine AGN in dust lane galaxies. but this is [ar [rom the case in he control sample. consistent. with findings of the previous section.," Almost all objects identified as radio AGN also show up as emission-line AGN in dust lane galaxies, but this is far from the case in the control sample, consistent with findings of the previous section."
+ ‘Table 3. shows the fraction of various AGN types. as defined in Paper L in dust lane early-type galaxies ancl the control sample.," Table \ref{tab:fracAllAGN} shows the fraction of various AGN types, as defined in Paper I, in dust lane early-type galaxies and the control sample."
+ The control sample shows a much lower level of emission-line ACIN. activity (<20 per cent) than dust lane carly tvpes (270 per cent)., The control sample shows a much lower level of emission-line AGN activity $\leq 20$ per cent) than dust lane early types $>70$ per cent).
+ Taken together with the above result that only dust lane galaxies (but not the control sample) show a significant correlation between emission-line ancl radio ACN activity. our interpretation is that dust lane earlv-tvpe galaxies correspond to a phase in which radio and optical AGN activity is triggered.," Taken together with the above result that only dust lane galaxies (but not the control sample) show a significant correlation between emission-line and radio AGN activity, our interpretation is that dust lane early-type galaxies correspond to a phase in which radio and optical AGN activity is triggered."
+ If dust. lane early types are indeed: associated: with eas-rich mergers. an obvious question concerns their place in the evolution of a post-merger early-type galaxy.," If dust lane early types are indeed associated with gas-rich mergers, an obvious question concerns their place in the evolution of a post-merger early-type galaxy."
+ Following we estimate the age of the last starburst using the double. colour (NUY )(g2) which is relatively dust insensitive.," Following we estimate the age of the last starburst using the double colour $($ $-u)-(g-z)$, which is relatively dust insensitive."
+ Fig., Fig.
+ 7 shows the starburst age distributions for the dust lane sample and a matched control sample of all earlv-type galaxies., \ref{fig:t2control} shows the starburst age distributions for the dust lane sample and a matched control sample of all early-type galaxies.
+ The starburst ages of dust lane early-type galaxies are vounger than the average carly-type galaxy., The starburst ages of dust lane early-type galaxies are younger than the average early-type galaxy.
+ As shown in the previous section. dust lane early-tvpe ealaxies have stellar populations that are vounger than the overall earlv-tvpe population.," As shown in the previous section, dust lane early-type galaxies have stellar populations that are younger than the overall early-type population."
+ Thus. these objects are likely to be starburst or post-starburst galaxies. observed. during the transition between the blue ancl rec-ancl-cdeacd phases.," Thus, these objects are likely to be starburst or post-starburst galaxies, observed during the transition between the blue and red-and-dead phases."
+ The correlation between optical and. radio AGN activity suggests gas-rich. minor mergers (since major mergers are rare in the low-redshift Universe) may drive star formation and AGN activity., The correlation between optical and radio AGN activity suggests gas-rich minor mergers (since major mergers are rare in the low-redshift Universe) may drive star formation and AGN activity.
+ In this scenario. the central engine of the AGN is endowed. with so much fuel that. the accretion disc settles into a raciatively efficient. optically," In this scenario, the central engine of the AGN is endowed with so much fuel that the accretion disc settles into a radiatively efficient, optically"
+Another important characteristic of this pleriou nebula is that its N-rav extent appears larecr then its radio extension. au effect also observed in 21.5-0.9 (273) and 3€ 58 (?7)).,"Another important characteristic of this plerion nebula is that its X-ray extent appears larger then its radio extension, an effect also observed in G21.5-0.9 \cite{wbb01}) ) and 3C 58 \cite{bwm01}) )."
+ This is unexpoectec if both N-rayv aud radioi are just the svuchrotron cussion of the nebula., This is unexpected if both X-ray and radio are just the synchrotron emission of the nebula.
+ However. the inverse Couptou aud uou-thermal brenisstrahluug colmpouncuts of the N-ray emission. as well as other tliermal components related to a possible shell. nav become Huportant in the outer halo ofthe nebula.," However, the inverse Compton and non-thermal bremsstrahlung components of the X-ray emission, as well as other thermal components related to a possible shell, may become important in the outer halo of the nebula."
+ Unfortunately. there is no wav to verify if. apart from the contamination of the IC £13 thermal emission. the spectrum of the outer nebula is purely nou-therma (like iu €:21.5-0.9. 273). or contains an additional theoria component due to the expansion of the nebula iuto he ejecta (as in 3€ 58. ?7)).," Unfortunately, there is no way to verify if, apart from the contamination of the IC 443 thermal emission, the spectrum of the outer nebula is purely non-thermal (like in G21.5-0.9, \cite{wbb01}) ), or contains an additional thermal component due to the expansion of the nebula into the ejecta (as in 3C 58, \cite{bwm01}) )."
+ On the other haud. such a πια radio jiebula iav be he result of a low seusitivitv of the VLA observations reported by 77..," On the other hand, such a small radio nebula may be the result of a low sensitivity of the VLA observations reported by \cite*{ocw01}."
+ The X-ray to radio huuinositv ratio. defined for dustance as in ?7.. is only a factor of 2 less hen the Crab's aud 100 nore than 3€ 58.," The X-ray to radio luminosity ratio, defined for instance as in \cite*{wsp97}, is only a factor of 2 less then the Crab's and 100 more than 3C 58."
+ But if deeper radio observations reveal a nebula comparable with. or arecr then. the Nay. nebula the real value of the ratio nav be reduced.," But if deeper radio observations reveal a nebula comparable with, or larger then, the X-ray nebula the real value of the ratio may be reduced."
+ By extrapolating back the sum of the spectra of he different nebula reeious in the radio regime. aud comparing them with the radio spectrum reported by ??.. we have found a spectral break at 100 GUIIz. instead of 10.000 CITz as reported by ??..," By extrapolating back the sum of the spectra of the different nebula regions in the radio regime, and comparing them with the radio spectrum reported by \cite*{ocw01}, we have found a spectral break at 100 GHz, instead of 10,000 GHz as reported by \cite*{ocw01}."
+ The difference is due to the contribution of the hard iuner nebula regions. which were not properly inodole« using ASCA data. on which the TT estimation relics upon.," The difference is due to the contribution of the hard inner nebula regions, which were not properly modeled using ASCA data, on which the \cite*{ocw01}
+ estimation relies upom."
+" Such a low frequency break is also found in 3€ 58 aud (210-00, as well as in other examples of other ""nou Crab-like plerionreported by ??.."," Such a low frequency break is also found in 3C 58 and G21.5-0.9, as well as in other examples of other “non Crab-like"" plerion reported by \cite*{wsp97}."
+ It should be noted. that there are some differences )etwoeen the IC. L138 plerion aud 3€ 58 (and. 621.5-0.9).," It should be noted, that there are some differences between the IC 443 plerion and 3C 58 (and G21.5-0.9)."
+ Tn particular. the spectral iudex steepening iu 3€ 58 and G21.5-0.9 is higher than 0.5. the value predicted w the one-zone meddle of ??..," In particular, the spectral index steepening in 3C 58 and G21.5-0.9 is higher than 0.5, the value predicted by the one-zone model of \cite*{che00}."
+ Another key difference is he ceutral source in the center. which is very promincut in case of the IC 113 pleriou.," Another key difference is the central source in the center, which is very prominent in case of the IC 443 plerion."
+ However. on the basis of the published Chandra data on G21.5-0.9 (27)). the xeseuce of central compact sources In C21.5-0.9 cannot )o excluded.," However, on the basis of the published Chandra data on G21.5-0.9 \cite{scs00}) ), the presence of central compact sources in G21.5-0.9 cannot be excluded."
+ As for the NS. a setter interpretation of the its cluperature constraint requires a careful analysis of the NS atinosphere effect (e.g. 77)) and eeucral relativitistic corrections.," As for the NS, a better interpretation of the its temperature constraint requires a careful analysis of the NS atmosphere effect (e.g., \cite{psz95}) ) and general relativitistic corrections."
+ The effect of the hydrogen atinosphliere could reduce the effective temperature typically by a factor 1.3 to 2 and thus could make the effective NS radius consisteut with the standard NS radius of —10 kin., The effect of the hydrogen atmosphere could reduce the effective temperature typically by a factor 1.3 to 2 and thus could make the effective NS radius consistent with the standard NS radius of $\sim$ 10 km.
+ It is worth noting that the NS temperature of KT.— 0.1 keV is nicely consistent with the IC £13 age of 30.000 vears (27)). if the standard NS cooling curve of superfiuid. NS is used (e.e.. ?7..- see their Fig.," It is worth noting that the NS temperature of $kT \sim$ 0.1 keV is nicely consistent with the IC 443 age of 30,000 years \cite{che99}) ), if the standard NS cooling curve of superfluid NS is used (e.g., \cite{ykl99}, - see their Fig."
+ 7). but a lower age cannot be excluded.," 7), but a lower age cannot be excluded."
+" The discrepaucy between our best guess for the NS surface temperature aud the temperature derived by ον by fitting the point source CNOU J061705.5]|222127 mav be due to the low value of Ny assumed by them (1.3«10?! 3}, ", The discrepancy between our best guess for the NS surface temperature and the temperature derived by \cite*{ocw01} by fitting the point source CXOU J061705.3+222127 may be due to the low value of $N_{\rm H}$ assumed by them $1.3\times 10^{21}$ $^{-2}$ ).
+We verified that this Ay value is rejected at confidence level by all our fits., We verified that this $N_{\rm H}$ value is rejected at confidence level by all our fits.
+ A deeper oobservation will better constrain the plivsical properties ofthe compact source., A deeper observation will better constrain the physical properties of the compact source.
+ If the hard power-law spect of the uchula core region ids extrapolated up to GeV reguue. 1 would provide a flux of 2.0 15.23410.7 ph sol. which is consistent with theEGRET fux of 3EG JOGL7|2238 (5.000.1«10* ph ? 1).," If the hard power-law spectrum of the nebula core region is extrapolated up to GeV regime, it would provide a flux of 2.0 $\times 10^{-7}$ ph $^{-2}$ $^{-1}$, which is consistent with the flux of 3EG J0617+2238 $5.0\pm 0.4\times 10^{-7}$ ph $^{-2}$ $^{-1}$ )."
+ The nebula is placed outside the uncertainty circle of the 3EC 0617|2238. but the systematic uncertaimties of theEGRET circle allows us to speculate abou possible association of the two sources.," The nebula is placed outside the uncertainty circle of the 3EG J0617+2238, but the systematic uncertainties of the circle allows us to speculate about possible association of the two sources."
+ ΕΠ. however. the GeV. regineEGRET photons indecd originate from svuchrotrou enuüssion of relativistic clectrous. this would seriously constrain the uchula model of 27. where + are accelerated by a relativistic wind termination shock aud then radiate in the downstream region.," If, however, the GeV regime photons indeed originate from synchrotron emission of relativistic electrons, this would seriously constrain the nebula model of \cite*{kc84} where $^{\pm}$ are accelerated by a relativistic wind termination shock and then radiate in the downstream region."
+" The maxima photon CLOYSV Qucasured in CV) produced by au accelerated: electron of maximal Loreutz factor Piya, iu the maguetic feld of the emitting region D, (measured iu GC) can be obtained from 8.6«10197=DI, C. The et svuchrotrou loss time. £42:5*107D,2D Possunt exceed the + acceleration tine by relativistic shock. t4 Qvhere B, is the maeuetic field in the acceleration region)."," The maximum photon energy $\varepsilon$ (measured in GeV) produced by an accelerated electron of maximal Lorentz factor ${\rm \Gamma_{\rm max}}$ in the magnetic field of the emitting region ${\rm _e}$ (measured in G) can be obtained from $8.6 \times 10^{16} \varepsilon \approx B_{\rm e}
+\Gamma_{\rm max}^2$ G. The $^{\pm}$ synchrotron loss time, ${\rm t_s \approx
+5\cdot 10^8 B_{\rm a}^{-2} \Gamma^{-1}}$ s, must exceed the $^{\pm}$ acceleration time by relativistic shock, ${\rm _a}$ (where ${\rm B_a}$ is the magnetic field in the acceleration region)."
+" Since f, cannot be less then the particle evroperiod. one"," Since ${\rm t_a}$ cannot be less then the particle gyroperiod, one"
+The meaning of the expansion of (he universe. or the ‘expansion of space. is explored using two phenomena: the motion of a test particle against a homogeneous background and (he cosmological redshift.,"The meaning of the expansion of the universe, or the `expansion of space,' is explored using two phenomena: the motion of a test particle against a homogeneous background and the cosmological redshift."
+ Contrary to some expectations. a particle removed [rom the ]Iubble flow never returns to il.," Contrary to some expectations, a particle removed from the Hubble flow never returns to it."
+ The cosmological redshift is not an ‘expansion effect: in special cases il can be separated. into a kinematic (special relativistic) part ancl a static (gravitational redshift) part. but in general must be thought of as the effect on light rays of their propagation through curved space-time.," The cosmological redshift is not an `expansion effect;' in special cases it can be separated into a kinematic (special relativistic) part and a static (gravitational redshift) part, but in general must be thought of as the effect on light rays of their propagation through curved space-time."
+ Space as such does not affect matter by 'expanding. but only through its curvature.," Space as such does not affect matter by `expanding,' but only through its curvature."
+ The acceptance of General Relativity (GB) in place of the work of Newton marked a great. chauge in the concepts and language of gravitational pliysics., The acceptance of General Relativity (GR) in place of the work of Newton marked a great change in the concepts and language of gravitational physics.
+ Space itself was taken from its previous role as the unuobstrusive stage on which evervthing happened. to become a major plaver in the game: it was measured and characterised by the metric tensor. aud it could curve and vibrate.," Space itself was taken from its previous role as the unobstrusive stage on which everything happened, to become a major player in the game: it was measured and characterised by the metric tensor, and it could curve and vibrate."
+ Iu the particular case of cosmology. space cefined by a network of comoving coordinates became an expanding rubber sheet. or perhaps a river. carving matter along with it iu its motion.," In the particular case of cosmology, space defined by a network of comoving coordinates became an expanding rubber sheet, or perhaps a river, carrying matter along with it in its motion."
+ At the same time Newtouian plivsies is still an extremely accurate approximation of CR as loug as local speeds are siíall and eravitational fields (space curvature) are weak., At the same time Newtonian physics is still an extremely accurate approximation of GR as long as local speeds are small and gravitational fields (space curvature) are weak.
+ In these reeious it should eive similar. if not ideutical. results to a full relativistic treatiuent.," In these regions it should give similar, if not identical, results to a full relativistic treatment."
+ But the idea of expanding space is hard to fit into the Newtonian picture., But the idea of expanding space is hard to fit into the Newtonian picture.
+ Do galaxies iu the ITubble flow move apart because they have a certain initial velocity aud an inertia. or because they are carried along iu expanding space?," Do galaxies in the Hubble flow move apart because they have a certain initial velocity and an inertia, or because they are carried along in expanding space?"
+ This is not au idle question., This is not an idle question.
+ If one thinks of a particle which is part of the backeround cosmiolosv. the answer oue ects for its qualitative motion will depend ou how one thinks of space.," If one thinks of a particle which is part of the background cosmology, the answer one gets for its qualitative motion will depend on how one thinks of space."
+ The first part of the following study is a mathematical iuvestigation of this problem: using frec-particle motion im au attempt to show just what is happening as the universe expands., The first part of the following study is a mathematical investigation of this problem: using free-particle motion in an attempt to show just what is happening as the universe expands.
+ The second part deals with the cosmological redshift., The second part deals with the cosmological redshift.
+ It is normally portrayed as a direct consequence of cosinological expansion. aud thus might be taken to show muuediatelv the reality of the ‘expansion of space.," It is normally portrayed as a direct consequence of cosmological expansion, and thus might be taken to show immediately the reality of the `expansion of space.'"
+ Towever. kinciatical derivations of it also appear. resulting iu the same forumlac.," However, kinematical derivations of it also appear, resulting in the same formulae."
+ We can expect. then. that clarifvine the plysical source (assmmineg we can eive that plirase a ieauiug) of the cosmological redshift will also tell us something about the nature of the expansion of space.," We can expect, then, that clarifying the physical source (assuming we can give that phrase a meaning) of the cosmological redshift will also tell us something about the nature of the expansion of space."
+with Gp.ijo)=oreOnin)— Le jr) the spherical Bessel functions. e;=K(jy—iji) and the correlation in the quadrupole at unequal times ΟΡ(ΙΡ)=fdpp?PyNdGn.pop.p). Le. Qplyo.to)xCp».,"with $I(\eta_1,\eta_2)=\xeb^{-1}(\eta_{\rm min})-1$ , $j_l(x)$ the spherical Bessel functions, $x_i=k(\eta_0-\eta_i)$ and the correlation in the quadrupole at unequal times $Q_{\rm P}(\eta_1,\eta_2) \equiv \int
+dp p^2 P_i(p) \dtqt(\eta_1,p) \dtqt(\eta_2,p)$, i.e. $Q_{\rm
+P}(\eta_0,\eta_0) \propto C_{T,2}$."
+ We should notice that the limits of the time integrals are the start aud end of the reionization process. Le. we Integrate over the time during which the inhomogeneities appear.," We should notice that the limits of the time integrals are the start and end of the reionization process, i.e. we integrate over the time during which the inhomogeneities appear."
+ For the backgrouud reiouization history we assume that the mean ionization fraction το grows frou. zero to unity between the redshifts z*—02*/2 aud z*+62°/2. which we will choose appropriately.," For the background reionization history we assume that the mean ionization fraction $\xeb$ grows from zero to unity between the redshifts $z^*-\delta z^*/2$ and $z^*+\delta
+z^*/2$, which we will choose appropriately."
+ We have calculated (2)) with a modified. version of CMBFAST (Seljak&Zaldarriaga 1996).., We have calculated \ref{cel}) ) with a modified version of CMBFAST \citep{seljak:96}. .
+ The COBE normalized polarization power spectrum is shown iufigs.l and 2 for ciffereut parameters., The COBE normalized polarization power spectrum is shown in \ref{fig:1} and \ref{fig:2} for different parameters.
+ Iu fig.l the backgrouud reionizatiou is given by z*=50 aud ὃς”=20.," In \ref{fig:1}
+ the background reionization is given by $z^*=50$ and $\delta z^*=20$."
+ The dotted line refers to he order contribution., The dotted line refers to the order contribution.
+ We recognize the leature ou large scales which appears because of the lomogeneous reionization background., We recognize the feature on large scales which appears because of the homogeneous reionization background.
+ The thick long-dashed. sdid. aud short-dashed lines refer to the seconcl order contribution due o inhomogeneous reionization.," The thick long-dashed, solid and short-dashed lines refer to the second order contribution due to inhomogeneous reionization."
+ The long-dashed line is a model with a patch size of 2=20h!Mpc. he solid line for 4?—10/7.tMpe aud the short-dashed line for 2=5htMpe. To uderstand the behavior of the second order effect we will perform two approximatious to he integral in (2)).," The long-dashed line is a model with a patch size of $R=20 h^{-1} {\rm Mpc}$, the solid line for $R=10 h^{-1} {\rm Mpc}$ and the short-dashed line for $R=5 h^{-1} {\rm
+Mpc}$ To understand the behavior of the second order effect we will perform two approximations to the integral in \ref{cel}) )."
+ First we realize that the expression gGjo.Ην.)Gg.ye). for reasonable j»arameters of the reionizatiou history. is a 2-dimeusional function with a narrow peak.," First we realize that the expression $g(\eta_0,\eta_1)g(\eta_0,\eta_2)I(\eta_1,\eta_2)$, for reasonable parameters of the reionization history, is a $2$ -dimensional function with a narrow peak."
+ This allows is to perform the two time integrations in (2)) and calculate the integrand at a certain time i uultiplied by ai area factor (557., This allows us to perform the two time integrations in \ref{cel}) ) and calculate the integrand at a certain time $\eta^{*}$ multiplied by an area factor $(\delta\eta)^2$.
+ The second feature we exploit is that the spherical Bessel function wr) has a tight peak at /=e for large multipole moments /. so we can approximate it with a à-[uuction.," The second feature we exploit is that the spherical Bessel function $j_l(x)$ has a tight peak at $l=x$ for large multipole moments $l$, so we can approximate it with a $\delta$ -function."
+ Therefore we get Os(0)with Ou=Ευ—4) as theangular size of a patch as seen today.," Therefore we get _0^3,with $\Theta_0=R/(\eta_0-\eta^*)$ as theangular size of a patch as seen today."
+ The only unknown quantity in this expression is the ellective time period 1., The only unknown quantity in this expression is the effective time period $\delta\eta$ .
+ For thecase with 2=10htMpe we get, For thecase with $R=10 h^{-1} {\rm Mpc}$ we get
+The bright X-ray nova AOG20-00 (V616 Mon) was discovered when it erupted in 1975 1975).,The bright X-ray nova A0620–00 (V616 Mon) was discovered when it erupted in 1975 \citep{elvis1975}.
+. Observations obtained alter its return to quiescence revealed that AQG20-00 is an, Observations obtained after its return to quiescence revealed that A0620–00 is an
+will ullimately allow for detailed fitting of data to simultaneously. determine atmospheric structure and abundance profiles for a given exoplanel atinosplhere.,will ultimately allow for detailed fitting of data to simultaneously determine atmospheric structure and abundance profiles for a given exoplanet atmosphere.
+ However. the fitting procedure is currently challenged by the fact that the models are under-constrained by the available data.," However, the fitting procedure is currently challenged by the fact that the models are under-constrained by the available data."
+ Backward modeling will become particularly. useful in the future once full spectra. of exoplanet atmospheres become available., Backward modeling will become particularly useful in the future once full spectra of exoplanet atmospheres become available.
+ In the meantime. given the minimal quantity of available data for GJ 1214b. we use a forward modeling approach in this paper. similar to the work of ?.. to predict the abundance profiles and spectral signature for this transiting super-Earth.," In the meantime, given the minimal quantity of available data for GJ 1214b, we use a forward modeling approach in this paper, similar to the work of \citet{mos11}, to predict the abundance profiles and spectral signature for this transiting super-Earth."
+ We calculate alinospheric abundances for the transiting super-Earth GJ Lith that result. from non-equilibrium chemical processes (photolysis and vertical mixing) using the photochenmistry model developed in ? and ? lor hot Jupiter exoplanets., We calculate atmospheric abundances for the transiting super-Earth GJ 1214b that result from non-equilibrium chemical processes (photolysis and vertical mixing) using the photochemistry model developed in \citet{zah09} and \citet{zah11} for hot Jupiter exoplanets.
+ This model is based on the 1-D photochemical kinetics code initiallydescribed in ? and ?.., This model is based on the 1-D photochemical kinetics code initiallydescribed in \citet{kas89} and \citet{kas90}.
+ While GJ 1214b is not a hot Jupiter. this code is generally applicable to exoplanets with hydrogen-rich atmosplieres.," While GJ 1214b is not a hot Jupiter, this code is generally applicable to exoplanets with hydrogen-rich atmospheres."
+ since ? (he code has been completely. re-written in (he C programming language., Since \citet{zah11} the code has been completely re-written in the C programming language.
+ Updates have been made to further generalize (he code for calculating the composition of any given exoplanet atmosphere., Updates have been made to further generalize the code for calculating the composition of any given exoplanet atmosphere.
+ Specific improvements include functionality for using an arbitrary temperature-pressure profile ancl increased flexibility in the altitude step size. along with a eeneralized treatment of UV photolvsis.," Specific improvements include functionality for using an arbitrary temperature-pressure profile and increased flexibility in the altitude step size, along with a generalized treatment of UV photolysis."
+ We use the photochemical kinetics eode to calculate chemical abundances for 61 atomic and molecular species in GJ 1214b's atmosphere as a [unction of height. (parametrized within the code in terms of atmospheric pressure) by simultaneously solving (he equations of continuity. and flux., We use the photochemical kinetics code to calculate chemical abundances for 61 atomic and molecular species in GJ 1214b's atmosphere as a function of height (parametrized within the code in terms of atmospheric pressure) by simultaneously solving the equations of continuity and flux.
+ In this study. we solve for 698 chemical reactions including 33 photolvsis reactions.," In this study, we solve for 698 chemical reactions including 33 photolysis reactions."
+ The reaction rates and phlotolvsis cross sections emploved in this work are listed in 2.theirTables2and3.., The reaction rates and photolysis cross sections employed in this work are listed in \citet[][their Tables 2 and 3]{zah11}.
+ Reverse reaction rales are calculated using the method outlined in?) as outlined below., Reverse reaction rates are calculated using the method outlined in \citet{vis11} as outlined below.
+" For reactions of the form A+B—C+D with forward reaction rate hy. the reverse rale (i.e.. the rate for C+D— A+B) is kh,=hypexp(-—AG/RT). where AG. the Gibbs free energv. is obtained [rom enthalpies and entropies of the of (he reactants and products. AG=(διονο- Sp)."," For reactions of the form ${\rm A} + {\rm B} \rightarrow {\rm C} + {\rm D}$ with forward reaction rate $k_f$ , the reverse rate (i.e., the rate for $ {\rm C} + {\rm D} \rightarrow {\rm A} + {\rm B}$ ) is $k_r=k_f\exp{\left(-\Delta G/RT\right)}$, where $\Delta G$, the Gibbs free energy, is obtained from enthalpies and entropies of the of the reactants and products, $\Delta G=H_A+H_B-H_C-H_D - T\left(S_A+S_B-S_C-S_D\right)$ ."
+" Associative reactions of the lorm A+B—AD. the rate for the reverse reaction (dissociation of AB) is hy=ATP)exp(-AG/RT). where 2,=10"" dvnes/em? is one atmosphere."," Associative reactions of the form ${\rm A} + {\rm B} \rightarrow {\rm AB}$, the rate for the reverse reaction (dissociation of AB) is $k_r=k_f\left(kT/P_{\circ}\right) \exp{\left(-\Delta G/RT\right)}$, where $P_{\circ}=10^6$ $^2$ is one atmosphere."
+ This must be done in both the low pressure and high pressure limits., This must be done in both the low pressure and high pressure limits.
+ Similarly. (he associative reverse of a dissocialive reaction is given by ky.=keG/KT)expC-AG/RT).," Similarly, the associative reverse of a dissociative reaction is given by $k_r=k_f\left(P_{\circ}/kT\right) \exp{\left(-\Delta G/RT\right)}$."
+ For this code we havefit the reverse reaction rates to the standard. Arrhenius form fh= AZexp(—T/C)., For this code we havefit the reverse reaction rates to the standard Arrhenius form $k=AT^b\exp{\left(-T/C\right)}$ .
+ This maximizes code flexibility and, This maximizes code flexibility and
+"original images before sky subtraction and registration; for the 2005 and 2006 datasets, star subtraction was done after sky subtraction and registration.","original images before sky subtraction and registration; for the 2005 and 2006 datasets, star subtraction was done after sky subtraction and registration."
+" For the M87 dataset, which encompasses data taken both in 2004 and 2005, we have re-reduced the 2004 images for consistency; however, this test showed that differences between the two methods are small as long as sufficient area is available far away from bright stars to set the sky background."," For the M87 dataset, which encompasses data taken both in 2004 and 2005, we have re-reduced the 2004 images for consistency; however, this test showed that differences between the two methods are small as long as sufficient area is available far away from bright stars to set the sky background."
+ Sky subtraction is problematic because the size of the Virgo Cluster is much bigger than the imaging field of view., Sky subtraction is problematic because the size of the Virgo Cluster is much bigger than the imaging field of view.
+" First, we mask the images of all sources brighter than 2.50 above sky, then spatially rebin the images by calculating the mode of the remaining unmasked pixels in 32x32 pixel “superpixels.”"," First, we mask the images of all sources brighter than $\sigma$ above sky, then spatially rebin the images by calculating the mode of the remaining unmasked pixels in 32x32 pixel “superpixels.”"
+ Regions of images far away from bright stars and galaxies are then selected in which to measure a preliminary sky value to subtract., Regions of images far away from bright stars and galaxies are then selected in which to measure a preliminary sky value to subtract.
+" At this point, the binned images are registered and an iterative plane fitting process is employed to fit residual sky planes to the binned images, constrained to minimize frame-to-frame variations in spatially-overlapping regions."," At this point, the binned images are registered and an iterative plane fitting process is employed to fit residual sky planes to the binned images, constrained to minimize frame-to-frame variations in spatially-overlapping regions."
+" The resulting edge-to-edge sky gradients which are removed from the images are typically about 5 ADU, or «0.596 of the absolute sky level (1100-1500 ADU, or —221.6-21.3)."," The resulting edge-to-edge sky gradients which are removed from the images are typically about 5 ADU, or $<$ of the absolute sky level (1100-1500 ADU, or 21.6-21.3)."
+" Once sky subtraction is complete, we register and median combine the individual images to create the final mosaic."," Once sky subtraction is complete, we register and median combine the individual images to create the final mosaic."
+" We quantify photometric uncertainty on two scales: small scale uncertainty due to read noise, photon noise, and small-scale flatfielding errors, and large scale uncertainties due to uncertainty in large scale flat fielding and sky subtraction."," We quantify photometric uncertainty on two scales: small scale uncertainty due to read noise, photon noise, and small-scale flatfielding errors, and large scale uncertainties due to uncertainty in large scale flat fielding and sky subtraction."
+" Small-scale uncertainties can be “root-N reduced"" by combining many individual images together into the final mosaic, and also by spatially binning the mosaic."," Small-scale uncertainties can be “root-N reduced” by combining many individual images together into the final mosaic, and also by spatially binning the mosaic."
+" Large-scale uncertainties do not reduce so easily, and we characterize them as a floor to our uncertainty, as described below."," Large-scale uncertainties do not reduce so easily, and we characterize them as a floor to our uncertainty, as described below."
+" We refer the reader to Morrison (1997), Feldmeier (2002), Mihos (2009), and Rudick (2010) for more details of our error model as applied to the deep surface photometry."," We refer the reader to Morrison (1997), Feldmeier (2002), Mihos (2009), and Rudick (2010) for more details of our error model as applied to the deep surface photometry."
+" For the CCD used in these observations, the gain was 26 /ADU and readnoise was 126” or 6 ADU."," For the CCD used in these observations, the gain was $^-$ /ADU and readnoise was $^-$ or 6 ADU."
+" The typical sky background was 1250 ADU, which produces photon noise of 24.5 ADU per pixel."," The typical sky background was 1250 ADU, which produces photon noise of 24.5 ADU per pixel."
+ Photon noise in the blank skies used to construct the flat field will also result in small-scale flatfielding uncertainty (which will be mediated significantly by our dithering of object images)., Photon noise in the blank skies used to construct the flat field will also result in small-scale flatfielding uncertainty (which will be mediated significantly by our dithering of object images).
+ The resulting small-scale flat fielding error is 2.7 ADU (see Mihos 2008)., The resulting small-scale flat fielding error is 2.7 ADU (see Mihos 2008).
+ These errors add in quadrature to give an uncertainty of 25 ADU/pixel in a single Virgo image., These errors add in quadrature to give an uncertainty of 25 ADU/pixel in a single Virgo image.
+" When median combining N individual images together, the noise goes down as 1.22/N; for a typical dataset analyzed here, N~30, giving a per pixel noise in the final mosaic of 0.6 ADU."," When median combining $N$ individual images together, the noise goes down as $1.22\sqrt{N}$; for a typical dataset analyzed here, $N \approx 30$, giving a per pixel noise in the final mosaic of 0.6 ADU."
+" Furthermore, when examining residual images, we spatially rebin the images in blocks of 9x9 pixels; these rebinned images then have a per pixel uncertainty of «0.1 ADU due to these small-scale effects."," Furthermore, when examining residual images, we spatially rebin the images in blocks of 9x9 pixels; these rebinned images then have a per pixel uncertainty of $<$ 0.1 ADU due to these small-scale effects."
+" Over larger scales, however, the dominant uncertainty lies in the large-scale flat fielding and sky subtraction."," Over larger scales, however, the dominant uncertainty lies in the large-scale flat fielding and sky subtraction."
+" Flat fielding errors can come from a variety of imperfections in the blank sky images used to construct the flat, such as spatial variations in the sky brightness, the unsubtracted extended wings of bright stars, and diffuse low surface brightness structures such as galactic cirrus."," Flat fielding errors can come from a variety of imperfections in the blank sky images used to construct the flat, such as spatial variations in the sky brightness, the unsubtracted extended wings of bright stars, and diffuse low surface brightness structures such as galactic cirrus."
+" Night sky variations in the object images also lead to uncertainty in the sky subtraction, as does the lack of substantial areas of pure sky in which to define the sky levels."," Night sky variations in the object images also lead to uncertainty in the sky subtraction, as does the lack of substantial areas of pure sky in which to define the sky levels."
+" Variability in the night sky happens both randomly and systematically (with, for example, hour angle, time of night, or airmass)."," Variability in the night sky happens both randomly and systematically (with, for example, hour angle, time of night, or airmass)."
+" To quantify these combined uncertainties, we measure the residual flux in 100x100 pixel boxes in blank sky regions in the flattened, sky-subtracted mosaic."," To quantify these combined uncertainties, we measure the residual flux in 100x100 pixel boxes in blank sky regions in the flattened, sky-subtracted mosaic."
+" These boxes show a mean residual sky background of 0.1 +/- 0.5 ADU, and so we adopt a conservative 1 ADU uncertainty as the floor to our error model."," These boxes show a mean residual sky background of 0.1 +/- 0.5 ADU, and so we adopt a conservative 1 ADU uncertainty as the floor to our error model."
+" Given our photometric solution, 1 ADU corresponds to a surface brightness of wy=229.2."," Given our photometric solution, 1 ADU corresponds to a surface brightness of 29.2."
+" Our approach to removing the smooth light profile from these galaxies involves fitting nested elliptical annuli over the galaxy's entire extent, and then subtracting the best-fit model from the image."," Our approach to removing the smooth light profile from these galaxies involves fitting nested elliptical annuli over the galaxy's entire extent, and then subtracting the best-fit model from the image."
+" As this fitting process is affected by all of the light in each isophote, it is imperative to remove any extra sources of non-galaxy light, lest they be interpreted as parts of the galaxy."," As this fitting process is affected by all of the light in each isophote, it is imperative to remove any extra sources of non-galaxy light, lest they be interpreted as parts of the galaxy."
+" 'The star mask and subtraction from the image reduction eliminates most of the light from the stars in each field, but other sources still remain."," The star mask and subtraction from the image reduction eliminates most of the light from the stars in each field, but other sources still remain."
+" We have masked, by hand, all light in each image from nearby galaxies as well as the bright wings, diffraction spikes, and column bleeds around bright stars."," We have masked, by hand, all light in each image from nearby galaxies as well as the bright wings, diffraction spikes, and column bleeds around bright stars."
+" To mask all small-scale sources of light, we use an automated source detection algorithm on an image with large scale structures removed."," To mask all small-scale sources of light, we use an automated source detection algorithm on an image with large scale structures removed."
+ The large- structure is removed by smoothing the original, The large-scale structure is removed by smoothing the original
+Since crowcling was not an issue. DAOPIIOT aperture photometry was used (ο measure point source brightnesses in each coadded image.,"Since crowding was not an issue, DAOPHOT aperture photometry was used to measure point source brightnesses in each coadded image."
+ Variable seeing was accounted for by setting the measurement aperture radius in each image to 1/2 the EWIIM of the stellar profile to maximize the signal-to-noise ratio., Variable seeing was accounted for by setting the measurement aperture radius in each image to 1/2 the FWHM of the stellar profile to maximize the signal-to-noise ratio.
+ The FWIIM was measured from a set of well exposed. isolated stars in each image.," The FWHM was measured from a set of well exposed, isolated stars in each image."
+ Calibration required the use of a diverse set of references., Calibration required the use of a diverse set of references.
+ For M32. we used the study of Magnierοἱal.(1992) which presents BYRI CCD photometry for 361.231 objects in the Ποια of M31 and also includes M32.," For M32, we used the study of \citet{mag92} which presents BVRI CCD photometry for 361,281 objects in the field of M31 and also includes M32."
+ For NGC 205. the DVRI photometry published in was used.," For NGC 205, the BVRI photometry published in \citet{lee96} was used."
+ NGC 147 and NGC 185 were both calibrated using V-banud photometry presented in the study of Nowotnyelal.(20023)., NGC 147 and NGC 185 were both calibrated using V-band photometry presented in the study of \citet{now03}.
+. These studies all made corrections for ealactic extinction., These studies all made corrections for galactic extinction.
+ In all cases. the epoch with the best photometric conditions was calibrated using stus in common wilh (hie references above and then all other epochs were calibrated to the reference epoch.," In all cases, the epoch with the best photometric conditions was calibrated using stars in common with the references above and then all other epochs were calibrated to the reference epoch."
+ The number of objects used in the calibrations were as follows: over LOO for M32. 8 for NGC 205. 21 for NGC 147. and 19 for NGC 185.," The number of objects used in the calibrations were as follows: over 100 for M32, 8 for NGC 205, 21 for NGC 147, and 19 for NGC 185."
+ In all cases we achieved a photometric accuracy of 0.1 magnitude or better in all filters., In all cases we achieved a photometric accuracy of 0.1 magnitude or better in all filters.
+ We measured stars in (he(ST images using very small apertures (1.6px for the WEDPC and 1.08px for the ACS)., We measured stars in the images using very small apertures (1.6px for the WFPC and 1.08px for the ACS).
+ This avoided contamination from cosmic ravs (in the ACS image) and contamination due to the crowding at aint magnitudes for both of these images., This avoided contamination from cosmic rays (in the ACS image) and contamination due to the crowding at faint magnitudes for both of these images.
+ Instead of determining aperture corrections and using the standard photometric calibration. we chose to use calibrated reference stars to bootstrap the calibration of our observed magnitudes.," Instead of determining aperture corrections and using the standard photometric calibration, we chose to use calibrated reference stars to bootstrap the calibration of our observed magnitudes."
+ We used the same references we used to calibrate the Tenagra images. with considerably fewer objects. due to their smaller fields of view.," We used the same references we used to calibrate the Tenagra images, with considerably fewer objects, due to their smaller fields of view."
+ For M32. only one of the relerence stars was available for photometric calibration. while lor NGC 205 six stars were available.," For M32, only one of the reference stars was available for photometric calibration, while for NGC 205 six stars were available."
+ The NGC 205 calibration has an RAIS scatter of 0.2 magnitudes in the V-band., The NGC 205 calibration has an RMS scatter of 0.2 magnitudes in the V-band.
+"where taking for the e1ipty sets (D;j aud [Cj the iudices ds)(T) aud ds,(7) equal zero.",where taking for the empty sets $B_j$ and $C_l$ the indices $ds_k^+(l)$ and $ds_k^-(l)$ equal zero.
+ The index dsj() describes a percentage of the contribution of the case when ος ando costaahx) are »ositive to the positive part of the hird term of taem (1)., The index $ds_k^+(l)$ describes a percentage of the contribution of the case when $c_\tau$ and $cos(\frac{\pi k \tau}{M})$ are positive to the positive part of the third term of the sum (1).
+ The index dss).(0) describes a sa.uar contriution. mmt for the case when the both ος and οςονά are snultaueouslv uceative.," The index $ds_k^-(l)$ describes a similar contribution, but for the case when the both $c_\tau$ and $cos(\frac{\pi k \tau}{M})$ are simultaneously negative."
+ Thauks to these oue can decide which the positive or the weeative values of he autocorrelation function have a larger contribution o the positive values of the estimator S$)., Thanks to these one can decide which the positive or the negative values of the autocorrelation function have a larger contribution to the positive values of the estimator $\hat{S_k}$.
+" When the difference es](1)ws),CI) is positive. the statement ‘the &-th peak really exists’ can not be rejected."," When the difference $ws_k^+(l)-ws_k^-(j)$ is positive, the statement 'the $k$ -th peak really exists' can not be rejected."
+ Thus. he following formula should be satisfied: Because the A-th peak could exist as a result of the echo-effect. it is necessary tfo verify the secoud condition: To verity the tuplication (8) firstly it is necessary to evaluate the sets €; for /CIe of the values of 7 for which the autocorrelation function iud the cosine function with the period d are positive aud the sets Dj.jCEg of values of 7 for which the autocorrelation function aud the cosine function with the period [24]k are negative.," Thus, the following formula should be satisfied: Because the $k$ -th peak could exist as a result of the echo-effect, it is necessary to verify the second condition: To verify the implication (8) firstly it is necessary to evaluate the sets $C_l$ for $l\in I_C$ of the values of $\tau$ for which the autocorrelation function and the cosine function with the period $\Big [\frac{2M}{k}\Big ]$ are positive and the sets $B_j, j\in I_B$ of values of $\tau$ for which the autocorrelation function and the cosine function with the period $\Big [\frac{2M}{k}\Big]$ are negative."
+ Secondly. a rcentaee of the contribution of the sum of products of positive values of ez aud costSF) to the sui of positive products of the values of ez and cos{ae) should be evaluated.," Secondly, a percentage of the contribution of the sum of products of positive values of $c_\tau$ and $cos(\frac{\pi k \tau}{M})$ to the sum of positive products of the values of $c_\tau$ and $cos(\frac{\pi k \tau}{M})$ should be evaluated."
+ As a result the indexes ds)(7) for each set C5 where 7 is the iudex frou the set Ze: are obtained., As a result the indexes $ds_k^+(l)$ for each set $C_l$ where $l$ is the index from the set $I_C$ are obtained.
+ Thirdly. from all sets €; such that /CFe - or which the index ds.(I) is the ereatest should be «ποσο.," Thirdly, from all sets $C_l$ such that $l\in I_C$ the set $C_m$ for which the index $ds_k^+(l)$ is the greatest should be chosen."
+planets (section [.1.2)).,planets (section \ref{sec:late_Z_supply}) ).
+ The relatively small amount of hydrogen aud helium present iu HD119026b is also to be explained by the relatively slow formation of the planet in a low-deusity euvironment (section L2.1))., The relatively small amount of hydrogen and helium present in HD149026b is also to be explained by the relatively slow formation of the planet in a low-density environment (section \ref{sec:limited_gas_supply}) ).
+ Another possibility is related to erosion of the envelope during giant impacts., Another possibility is related to erosion of the envelope during giant impacts.
+ A reasonable impact velocity of 2-2.5 times the two-bods surface escape velocity was found to result in a substantial loss of the envelope gas. while solid cores are probably merged (section 1.2.2)).," A reasonable impact velocity of 2–2.5 times the two-body surface escape velocity was found to result in a substantial loss of the envelope gas, while solid cores are probably merged (section \ref{sec:envelope_loss}) )."
+ Iu sumanary. we Can propose at least two scenarios for the origin of HD1(00560.," In summary, we can propose at least two scenarios for the origin of HD149026b."
+ Above two scenarios are different from each other in that the former leaves at least one outer eiaut planet while the latter does not uecessarily need auy outer giant. planets., Above two scenarios are different from each other in that the former leaves at least one outer giant planet while the latter does not necessarily need any outer giant planets.
+ Loue-term racial velocity observatious of HDI19026) are thus needed to cletermine the presence of other plauets in the system., Long-term radial velocity observations of HD149026b are thus needed to determine the presence of other planets in the system.
+ We thank the anonymous referee for fruitful comments., We thank the anonymous referee for fruitful comments.
+ This research was supported. by Ministry of Education. Culture. Sports. Science aud Technology of Japan. Grant-in-Aid for Scientific Research on Priority Areas. “Development of Extra-solar Planetary Science”.," This research was supported by Ministry of Education, Culture, Sports, Science and Technology of Japan, Grant-in-Aid for Scientific Research on Priority Areas, “Development of Extra-solar Planetary Science”."
+also triggered MCAL and SuperAGILE and is the subject of a dedicated paper (?)..,also triggered MCAL and SuperAGILE and is the subject of a dedicated paper \citep{Giuliani2008b}.
+ This GRB is further considered in subsection 3.3 regarding spectral fitting of MCAL data., This GRB is further considered in subsection \ref{ecal} regarding spectral fitting of MCAL data.
+ Figure 3 shows the MCAL sensitivity to GRBs., Figure \ref{FigMCAL_sensitivity_theta40} shows the MCAL sensitivity to GRBs.
+ The sensitivity has been calculated according to the procedure described in ? and reports thepeak flux in the 100—1000keV energy band at the detector's threshold (So significance above background in | second). as a function of the peak energy. for three sample spectral shapes modeled according to the Band Model (?)..," The sensitivity has been calculated according to the procedure described in \citet{Band2003} and reports thepeak flux in the $100-1000~\mathrm{keV}$ energy band at the detector's threshold $5\sigma$ significance above background in 1 second), as a function of the peak energy, for three sample spectral shapes modeled according to the Band Model \citep{Band1993}."
+ The sensitivity curves have been calculated using the MCAL efficiency at a 40° off-axis angle resulting from Monte Carlo simulations and the in-flight background level (about 210 counts/s over all the detector in the 330-700 keV trigger band)., The sensitivity curves have been calculated using the MCAL efficiency at a $40^\circ$ off-axis angle resulting from Monte Carlo simulations and the in-flight background level (about 210 counts/s over all the detector in the 330-700 keV trigger band).
+ Superimposed on the sensitivity curves are data points corresponding to GRBs with public spectral parameters. published as GRB Coordinates Network circulars either by the Suzaku-WAM or Konus-Wind teams.," Superimposed on the sensitivity curves are data points corresponding to GRBs with public spectral parameters, published as GRB Coordinates Network circulars either by the Suzaku-WAM or Konus-Wind teams."
+ Solid triangles refer to GRBs also detected by MCAL at an incident angle lower than 90°: GRB 071125 (2?).. GRB 080204 (??).. GRB 080319C (??)..," Solid triangles refer to GRBs also detected by MCAL at an incident angle lower than $90^\circ$: GRB 071125 \citep{OhnoGCN2007,Golenetskii2007GCN7137}, GRB 080204 \citep{TeradaGCN2008,Golenetskii2008GCN7263}, GRB 080319C \citep{OndaGCN2008,Golenetskii2008GCN7487}."
+ Solid circles refer to GRBs detected by MCAL at incident angles greater than 90°: GRB 071227 (??).. GRB 080122 (?2).. GRB 080328 (2?).," Solid circles refer to GRBs detected by MCAL at incident angles greater than $90^\circ$: GRB 071227 \citep{OndaGCN2008b,GolenetskiiGCN2007}, GRB 080122 \citep{UeharaGCN2008,Golenetskii2008GCN7219}, , GRB 080328 \citep{KodakaGCN2008,GolenetskiiGCN2008}."
+.For this plot. the Suzaku-WAM peak flux values have been used because they," For this plot, the Suzaku-WAM peak flux values have been used because they"
+Eve (IBRes) has claimed proton-douinated composition even above 1029 eV2010).. There are different arguments aud the UITECT. composition has not oen settled experimentally.,Eye (HiRes) has claimed proton-dominated composition even above $10^{19}$ eV. There are different arguments and the UHECR composition has not been settled experimentally.
+ Proton composition seenis »ossible at present., Proton composition seems possible at present.
+ If UIIECTR sources are transient. that is. the source activity is shorter than the dispersion of the arrival nue produced dy cosmic magnetic. fields during xopasation. the direct ideutification of VITECR sources o» UIIECR observations is a more difficult task han that for steady sources due to the delav of he arrival tine between UIIECRs and other neutral nesseugers (photons. neutrinos aud gravitational waves) enütted by the same source activitv.," If UHECR sources are transient, that is, the source activity is shorter than the dispersion of the arrival time produced by cosmic magnetic fields during propagation, the direct identification of UHECR sources by UHECR observations is a more difficult task than that for steady sources due to the delay of the arrival time between UHECRs and other neutral messengers (photons, neutrinos and gravitational waves) emitted by the same source activity."
+ Multiaunesseuser approaches are defuitelv powerful but it is important ο extract as ΙΟ information as possible from UITEC'R observations as one of the messcuecrs.," Multi-messenger approaches are definitely powerful, but it is important to extract as much information as possible from UHECR observations as one of the messengers."
+ For trausicut sources. relatious between observed quantities and the xoperties of CIECR sources lave been discussed. considering intervene cosnuüc magnetic fields2009).," For transient sources, relations between observed quantities and the properties of UHECR sources have been discussed, considering intervening cosmic magnetic fields."
+. Auy cauclidate of primary UIIECR sources has to possess enough ΟΠΟΙΟΥ: budget to reproduce the observed flux. which is the product of the energv input per activity aud the rate of bursts or flares.," Any candidate of primary UHECR sources has to possess enough energy budget to reproduce the observed flux, which is the product of the energy input per activity and the rate of bursts or flares."
+ The rate of bursts or flares is related to the apparent UIITECT source nuuber deusity. which can be determined from anisotropy in UIIEC'R arrival distribution.," The rate of bursts or flares is related to the apparent UHECR source number density, which can be determined from anisotropy in UHECR arrival distribution."
+ The relation between the rate and the apparent source nuniber deusity is less obvious. depending ou the Galactic magnetic field (CAIF) and poorly kuown extragalactic magnetic fields (ECGAIFs).," The relation between the rate and the apparent source number density is less obvious, depending on the Galactic magnetic field (GMF) and poorly known extragalactic magnetic fields (EGMFs)."
+ The CAIF unavoidably affects UITECTiS arriving at the Earth., The GMF unavoidably affects UHECRs arriving at the Earth.
+ Considering the GMPE allows us to evaluate the above relation aud to give a constraint on candidates of transient UIIECT. sources through comparing between the rate interred from ΕΠΕΟΤ observations aud that of known trausieut phenomena2009)., Considering the GMF allows us to evaluate the above relation and to give a constraint on candidates of transient UHECR sources through comparing between the rate inferred from UHECR observations and that of known transient phenomena.
+. Despite current uncertainty in EGMESR. it has been believed. that the imhomogencity of the EGAIFs in the Universe is crucial for the propagation of VITECRs2008).," Despite current uncertainty in EGMFs, it has been believed that the inhomogeneity of the EGMFs in the Universe is crucial for the propagation of UHECRs."
+". The Universe indeed has structures, which consist of clusters of galaxies. filaments. sheets and voids."," The Universe indeed has structures, which consist of clusters of galaxies, filaments, sheets and voids."
+ It has been sugeested that magnetic fields in the structured regions were amplified via cosimological structure formation. and. various nunerical sinulatious of the structure formation have shown that the EGAIF distribution follows the matter distribution2008).," It has been suggested that magnetic fields in the structured regions were amplified via cosmological structure formation, and various numerical simulations of the structure formation have shown that the EGMF distribution follows the matter distribution."
+ So. astroplivsical nerallyobjects. iucludiug astrophysical UITECR sources are ec enibedded iu the structured regious.," So, astrophysical objects including astrophysical UHECR sources are generally embedded in the structured regions."
+ Thus. these magnetic structures also unavoidably affect the propagation of UIIECTi's.," Thus, these magnetic structures also unavoidably affect the propagation of UHECRs."
+ The structured ECGAIFs play. essential roles iu the time-delay and time-profile spread of UITECTs as well as their deflections., The structured EGMFs play essential roles in the time-delay and time-profile spread of UHECRs as well as their deflections.
+ Iu this paper. we study the roles of structured ECGMES ou the propagation of UIIECTRs. and see how the properties of transient UIIECT sources. e.g.. the rate of UIIECR bursts or flares py. can be constrained.," In this paper, we study the roles of structured EGMFs on the propagation of UHECRs, and see how the properties of transient UHECR sources, e.g., the rate of UHECR bursts or flares $\rho_s$, can be constrained."
+ Iu section 2.. we describe basic relations between the CHECR burst rate aud observational quantities that can be obtained by VITECR experiments.," In section \ref{idea}, we describe basic relations between the UHECR burst rate and observational quantities that can be obtained by UHECR experiments."
+ Possible signatures 6| transient UIIECTR sources that may be seen in the arrival distribution of UIIECTis are also ciscussed., Possible signatures of transient UHECR sources that may be seen in the arrival distribution of UHECRs are also discussed.
+" Iu section ον,Ah. we nunericallv calculate the propagation of CHECRs in the magnetized structures for various source positions. aud evaluate the time spread due to the structured fields."," In section \ref{prop}, we numerically calculate the propagation of UHECRs in the magnetized structures for various source positions, and evaluate the time spread due to the structured fields."
+ Then. in section L. we discuss possible constraints on the ΤΙΠΟΤΑ burst rate and cosmic-ray enerev input burst. taking mto account large uucertaimtyv in the void EGME.," Then, in section \ref{results}, we discuss possible constraints on the UHECR burst rate and cosmic-ray energy input burst, taking into account large uncertainty in the void EGMF."
+ We present REA in section 5 and stnunarize this study in section 6.., We present discussions in section \ref{conclusion} and summarize this study in section \ref{summary}. .
+" Throughout the paper. the protou-conrposition is assuied. aud AC DM cosmology with Hj=11 kun | +. Q,,—033aud Q4—0,7 is adopted."," Throughout the paper, the proton-composition is assumed, and $\Lambda$ CDM cosmology with $H_0 = 71$ km $^{-1}$ $^{-1}$, $\Omega_m = 0.3$ and $\Omega_{\Lambda} = 0.7$ is adopted."
+ UTECRs above 5«101? eV. cannot avoid enerey-loss due to photonieson production with cosmic microwave backeround (CMD) plotous1966)., UHECRs above $5 \times {10}^{19}$ eV cannot avoid energy-loss due to photomeson production with cosmic microwave background (CMB) photons.
+". Παισο, nearby ΕΠΟ sources within the so-called Caciscu-Zatsepin-huzgun (GZRK) radius mainly contribute to the observed fiw.which is typically the enerev-loss leugth of ~100 Ape at ος1015 eV, For practical purpose. we here introduce D,(E) that is the maxima distance of sources of UITECRS observed with the enerey £ at the Earth."," Hence, nearby UHECR sources within the so-called Greisen-Zatsepin-Kuz'min (GZK) radius mainly contribute to the observed flux,which is typically the energy-loss length of $\sim 100$ Mpc at $6 \times {10}^{19}$ eV. For practical purpose, we here introduce $D_{\rm max}(E)$ that is the maximum distance of sources of UHECRs observed with the energy $E$ at the Earth."
+ This is defined as the distance within which sources are responsible for 99 of the total observed flux., This is defined as the distance within which sources are responsible for 99 of the total observed flux.
+ This definition is more appropriate than what we used in(2009).. where the energv-loss leneth of the photomeson xoduction is simply adopted. because sources outside ie energv-loss leneth can still significantly contribute o the total observed flux due to a volume effect.," This definition is more appropriate than what we used in, where the energy-loss length of the photomeson production is simply adopted, because sources outside the energy-loss length can still significantly contribute to the total observed flux due to a volume effect."
+ We estimate. Dyas(£) through a backtracking method of 1ο propagation of UITECRs with a continuous enerev-oss approNimation on the photomeson production aud Bethe-Heitler pair creation with €MD photons, We estimate $D_{\rm max}(E)$ through a backtracking method of the propagation of UHECRs with a continuous energy-loss approximation on the photomeson production and Bethe-Heitler pair creation with CMB photons.
+2006).. Figure 1 shows Diyas(EE) caleulated ou je assinption that ΕΠΟ ciuissivity is proportional o (11QS and the spectral iudex of -2.6. where τ is cosinological redshift.," Figure \ref{fig:dmax} shows $D_{\rm max}(E)$ calculated on the assmption that UHECR emissivity is proportional to ${(1 + z)}^3$ and the spectral index of -2.6, where $z$ is cosmological redshift."
+" Note that D,4(E) is not scusitive ou these assinaptionus since it is esseutiallv determined by he cnerey-loss of the plotomeson production.", Note that $D_{\rm max}(E)$ is not sensitive on these assumptions since it is essentially determined by the energy-loss of the photomeson production.
+ One sees hat DG) rapidly decreases owing to the photomeson xoduction with CAIB photons with increasing the ΟΠΟΙΟΥ of protons., One sees that $D_{\rm max}(E)$ rapidly decreases owing to the photomeson production with CMB photons with increasing the energy of protons.
+ DAGE) is typically 200 Alpe aud 5 Ape at E=6<10! eV and 1079 eV for Ej=107! eV. respectively.," $D_{\rm max}(E)$ is typically 200 Mpc and 75 Mpc at $E = 6 \times 10^{19}$ eV and $10^{20}$ eV for $E_{\rm max} = 10^{21}$ eV, respectively."
+ Trajectories of UIIECTRs are deflected im magnetizedextragalactic space. whose imaenetic fields are still unecrtain review).," Trajectories of UHECRs are deflected in magnetizedextragalactic space, whose magnetic fields are still uncertain ."
+. Asstumine that ECAIFs are incoherent (AL.« D). atypical deflection anele of UMECRs during propagation in extragalactic space is estimated tobe," Assuming that EGMFs are incoherent $\lambda \ll D$ ), a typical deflection angle of UHECRs during propagation in extragalactic space is estimated to be"
+Our existing models required the presence of hot. small ervstalline grains in the region of 1. AU with no evidence for material on more distant. orbits.,"Our existing models required the presence of hot, small crystalline grains in the region of 1 AU with no evidence for material on more distant orbits."
+ The new observations confirm and extend these results., The new observations confirm and extend these results.
+ The physical model used to characterize (he mineralogy of the LD 69330. dust. was developed [or Spitzer observations of the Tempel 1/Deep Impact data 2006)., The physical model used to characterize the mineralogy of the HD 69830 dust was developed for Spitzer observations of the Tempel 1/Deep Impact data .
+. In our case. optically (hin emission originates in material located in a narrow annulus around the star.," In our case, optically thin emission originates in material located in a narrow annulus around the star."
+ A wide variety of minerals ancl ices are successively added. into the model until the 4? of the fit stops changing by a significant amount., A wide variety of minerals and ices are successively added into the model until the $\chi^2$ of the fit stops changing by a significant amount.
+ Additional variables include the size distribution of the dust. grains and distance to the host star., Additional variables include the size distribution of the dust grains and distance to the host star.
+ Our model of the 2004 data revealed a population of small. super-thermal grains at about. 425 lx composed of silicates (both olivine and pyroxene. but mainly forsterite. [avalite. and ferrosilite). a smattering of sullides and carbonates. and surprisingly some water ice.," Our model of the 2004 data revealed a population of small, super-thermal grains at about 425 K composed of silicates (both olivine and pyroxene, but mainly forsterite, fayalite, and ferrosilite), a smattering of sulfides and carbonates, and surprisingly some water ice."
+ Using the higher SNR Loltes spectra rellores2)). we now reevaluate the mineralogy of the ILD69330 svstem.," Using the higher SNR LoRes spectra \\ref{lores2}) ), we now reevaluate the mineralogy of the HD69830 system."
+ Figure G shows the best fit model broken clown into its various species components while Table 6 lists the relative abundances of species identified in the LoRes spectrum. as well as signilicant," Figure \ref{fit} shows the best fit model broken down into its various species components while Table \ref{caseytable} lists the relative abundances of species identified in the LoRes spectrum, as well as significant non-detections."
+ The final reduced. 47? for the model fit to the data is 1.01., The final reduced $\chi^2$ for the model fit to the data is 1.01.
+ For the most part. our results are the same as in our previous work. with values and upper limits improved due to the much hieher «quality aud SNR. of the 2007 data.," For the most part, our results are the same as in our previous work, with values and upper limits improved due to the much higher quality and SNR of the 2007 data."
+ Each of the three 2007 spectra included in Figure ο had much longer integration times than did the 2004 data. e.g. 70 and 480 seconds in the SLI and LL1 modules in 2007 vs. 12 seconds ancl 28 seconds in 2004.," Each of the three 2007 spectra included in Figure \ref{HD69830LoRes} had much longer integration times than did the 2004 data, e.g. 70 and 480 seconds in the SL1 and LL1 modules in 2007 vs. 12 seconds and 28 seconds in 2004."
+ Differences from the earlier analvsis are flagged in the table and include the new (albeit weak) detection of phiyllosilicate tale aud iron oxide aud the failure to confirm a number of minor carbonates and pyroxene species., Differences from the earlier analysis are flagged in the table and include the new (albeit weak) detection of phyllosilicate talc and iron oxide and the failure to confirm a number of minor carbonates and pyroxene species.
+ The one important difference in (he new analysis is the lack of convincing evidence for the broad feature due to water ice (11-15 yam). a spectral region now understood to have serious problems with the tear-drop artifact in the data rellores)).," The one important difference in the new analysis is the lack of convincing evidence for the broad feature due to water ice (11-15 $\mu$ m), a spectral region now understood to have serious problems with the tear-drop artifact in the data \\ref{lores}) )."
+ The new spectra show no obvious emission from water ice. which is significantly different (han the model findings reported in(2007).," The new spectra show no obvious emission from water ice, which is significantly different than the model findings reported in."
+. The newer data. when compared to the old. show shifts in the 10-35 jan region where waler ice emission is most pronounced.," The newer data, when compared to the old, show shifts in the 10-35 $\mu$ m region where water ice emission is most pronounced."
+ As the other mineralogical markers seem to show that the dust in the svstem has, As the other mineralogical markers seem to show that the dust in the system has
+draw a complete picture of the of the iun-situ formation and evolution scenario for the cold classicals.,draw a complete picture of the of the in-situ formation and evolution scenario for the cold classicals.
+ However. the considerable improvement of the model for carly dvuamical evolution of the wiper belt prescuted here supports the overall validity of the We thauk Alessandro Morbidelli. Wal Levisou. Darin Ragozzine aud Peter Goldreich for useful conversations.," However, the considerable improvement of the model for early dynamical evolution of the Kuiper belt presented here supports the overall validity of the We thank Alessandro Morbidelli, Hal Levison, Darin Ragozzine and Peter Goldreich for useful conversations."
+sampling. iDrizzle could be used to recover the PSF if the telescope has either temporal or spatial stability.,"sampling, iDrizzle could be used to recover the PSF if the telescope has either temporal or spatial stability."
+ With temporal stability. images of dillerent stars with similar colors taken near (he same location on the detector could be combined to form a single PSF.," With temporal stability, images of different stars with similar colors taken near the same location on the detector could be combined to form a single PSF."
+ Similarly. if the PSF is stable over a sufficiently. wide field. different stars of similar color from a single image mav be combined as a substitute for an increased. number of dithers.," Similarly, if the PSF is stable over a sufficiently wide field, different stars of similar color from a single image may be combined as a substitute for an increased number of dithers."
+ While in both cases one will have to fit lor (he magnitudes ancl relative positions of the stars. as magnitudes and positions are identical (up to a known shift) for all dither sets of (he same star. the small nunmber of parameters fits should tvpicallv add little noise. particularly if the fitting is done itlerativelv while solving for the shape of the PSF.," While in both cases one will have to fit for the magnitudes and relative positions of the stars, as magnitudes and positions are identical (up to a known shift) for all dither sets of the same star, the small number of parameters fits should typically add little noise, particularly if the fitting is done iteratively while solving for the shape of the PSF."
+ The new method presented here. iDrizzle. may find use in a number of applications where high image lidelityv is required.," The new method presented here, iDrizzle, may find use in a number of applications where high image fidelity is required."
+ In order that the community mar test this algorithm and expand upon its use. the Pyraf script used to create (he UDF image. along with the UDF data used in this test. is posted on the web for download by interested readers ad [ruchter/iDrizzle/UDE-iDrizzle.," In order that the community may test this algorithm and expand upon its use, the Pyraf script used to create the UDF image, along with the UDF data used in this test, is posted on the web for download by interested readers at fruchter/iDrizzle/UDF-iDrizzle."
+ ] wish to thank Robert Lupton lor helpful and amusing discussions. the Institute of Pure and Applied Mathematics at UCLA for hosting a workshop on astronomical imaging where I was first able to discuss some of these ideas. and the Space Telescope Science Institute. which supported this work.," I wish to thank Robert Lupton for helpful and amusing discussions, the Institute of Pure and Applied Mathematics at UCLA for hosting a workshop on astronomical imaging where I was first able to discuss some of these ideas, and the Space Telescope Science Institute, which supported this work."
+ ILans-Maurtin. Adorf and Thomas Strolimer gave very. useful thoughts on analytical methods related to irregular sampling., Hans-Martin Adorf and Thomas Strohmer gave very useful thoughts on analytical methods related to irregular sampling.
+ Anton IXoekemoer provided the aligned UDF data. and helped with the first Python implementation of the iDrizzle algoritlim.," Anton Koekemoer provided the aligned UDF data, and helped with the first Python implementation of the iDrizzle algorithm."
+ I also thank the releree. Tod Lauer. as well as Stefano Casertano. who both provided comments that lead to a clearer final manuscript.," I also thank the referee, Tod Lauer, as well as Stefano Casertano, who both provided comments that lead to a clearer final manuscript."
+since Shaplev's pioneering work (Shaplev1918).. globular clusters (GCs) have plaved a kev-role in our understanding of ihe Universe and of the manner in which our Galaxy,"Since Shapley's pioneering work \citep{shapley}, globular clusters (GCs) have played a key-role in our understanding of the Universe and of the manner in which our Galaxy"
+ f=74+(2.751+0.011)(0.080.01)xIX(0.0450.008)(V—I) The finalr.m.s of the fitting was 0.030. 0.015. 0.010. and 0.010 in U. D. V and J. Global photometric errors were estimated using the scheme developed by Patat Carraro (2001. Appendix Al). which takes into account the errors resulting from the PSF [itting procedure (i.e.. from ALLSTAR). and the calibration errors (corresponding to the zero point. color terms. and extinction errors).,"$ V = v + (1.760\pm0.007) + (0.15\pm0.01) \times X + (0.028\pm0.007) \times (B-V)$ $ I = i + (2.751\pm0.011) + (0.08\pm0.01) \times X + (0.045\pm0.008) \times (V-I)$ The final of the fitting was 0.030, 0.015, 0.010, and 0.010 in $U$ , $B$, $V$ and $I$, Global photometric errors were estimated using the scheme developed by Patat Carraro (2001, Appendix A1), which takes into account the errors resulting from the PSF fitting procedure (i.e., from ALLSTAR), and the calibration errors (corresponding to the zero point, color terms, and extinction errors)."
+ In Fie., In Fig.
+ 2 we present our global photometric errors in V. (D— V). (U—B). and (V£) plotted as a function of V magnitude.," 2 we present our global photometric errors in $V$, $(B-V)$ , $(U-B)$, and $(V-I)$ plotted as a function of $V$ magnitude."
+ Quick inspection shows that stars brighter (han V.2220 mag have errors lower than 0.05 mag in magnitude and lower (han ~0.10 mag in (2—V) and (V—I)., Quick inspection shows that stars brighter than $V \approx 20$ mag have errors lower than $\sim0.05$ mag in magnitude and lower than $\sim0.10$ mag in $(B-V)$ and $(V-I)$.
+ Iligher errors are seen in (U—D). Our final optical photometric catalog consists of 13038. entries having. measurements down to V.—20. and 43471 entries having measures down to V.—22.," Higher errors are seen in $(U-B)$ Our final optical photometric catalog consists of 13038 entries having measurements down to $V\sim 20$, and 43471 entries having measures down to $V\sim 22$."
+ Completeness corrections were determined by running artificial star experiments on the data., Completeness corrections were determined by running artificial star experiments on the data.
+" Basically. we created several artificial images by adding artificial stars to the original frames,"," Basically, we created several artificial images by adding artificial stars to the original frames."
+ These stars were added at random positions. aud had (he same color ancl Iumninositydistribution of the true sample.," These stars were added at random positions, and had the same color and luminositydistribution of the true sample."
+To avoid generating overcrowding. in each experiment we,"To avoid generating overcrowding, in each experiment we"
+We thank all the member of the SCA team for the operation of the satellite and maintenance of the soltware. John Alulehaey for information on his unpublished CCD optical image and. useful comments. ancl Yasushi Fukazawa for providing the ASCA results on the cluster. sample in his thesis.,"We thank all the member of the ASCA team for the operation of the satellite and maintenance of the software, John Mulchaey for information on his unpublished CCD optical image and useful comments, and Yasushi Fukazawa for providing the ASCA results on the cluster sample in his thesis."
+ This research has made use of data obtained through the Ligh LEnerey Astrophysics Science Archive Research Center Online Service. provided. hy the NASA's Goddard Space Flight Center.," This research has made use of data obtained through the High Energy Astrophysics Science Archive Research Center Online Service, provided by the NASA's Goddard Space Flight Center."
+ Phe NASA/IPAC Extragalactic Database (NED) is operated by the Jet Propulsion Laboratory. California Institute of Technology. under contract with the National Acronautics ancl Space Administration.," The NASA/IPAC Extragalactic Database (NED) is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration."
+ Roval Society (ACE.ACESE) and PPARC (INL) are thanked for support.," Royal Society (ACE,ACF,SE) and PPARC (KI) are thanked for support."
+(Ie). and (Mg=—18.4£0.8 mag (le-BL).,"(Ic), and $\langle M_{R,\rm peak}\rangle =-18.4\pm 0.8$ mag (Ic-BL)."
+ We next comparepea? the light curve shapes of the literature sample after shifting each curve relative to the epoch and magnitude of maximum light (Figure 5))., We next compare the light curve shapes of the literature sample after shifting each curve relative to the epoch and magnitude of maximum light (Figure \ref{fig:VR_curves}) ).
+ This figure highlights the significant dispersion in the light curve widths (see also Cloechiatti&Wheeler 1997a..Clocchiatti&Wheeler 1997b)).," This figure highlights the significant dispersion in the light curve widths (see also \citealt{cw97a}, \citealt{cw97b}) )."
+ We build template light curves in the V— and R-bands by interpolating over these normalized light curves and extracting the weighted mean flux density for time intervals spanning —20 to +40 days with respect to the epoch of maximum light., We build template light curves in the $V-$ and $R-$ bands by interpolating over these normalized light curves and extracting the weighted mean flux density for time intervals spanning $-20$ to $+40$ days with respect to the epoch of maximum light.
+ The uncertainties in the resulting templates are derived from the standard deviation in the interpolated curves at each time interval and represent the intrinsic dispersion in. light curve widths., The uncertainties in the resulting templates are derived from the standard deviation in the interpolated curves at each time interval and represent the intrinsic dispersion in light curve widths.
+ Our resulting template curves are shown in Figure 5.., Our resulting template curves are shown in Figure \ref{fig:VR_curves}.
+ We measure the post-maximum decay rate for each of the literature SNe using the Ais diagnostic which represents the decay in magnitudes during the 15 days following the light curve maximum (Phillips.1993)., We measure the post-maximum decay rate for each of the literature SNe using the $\Delta m_{15}$ diagnostic which represents the decay in magnitudes during the 15 days following the light curve maximum \citep{phi93}.
+. By interpolating over the well-sampled V— and R-band light curves. we derive Ans and Anis. estimates. respectively (Table 4)).," By interpolating over the well-sampled $V-$ and $R-$ band light curves, we derive $\Delta m_{15,V}$ and $\Delta
+m_{15,R}$ estimates, respectively (Table \ref{tab:templates_lc}) )."
+ We note that this is à variation on the standard definition for Az which assumes a time baseline relative to the B—band light curve maximum.," We note that this is a variation on the standard definition for $\Delta
+m_{15}$ which assumes a time baseline relative to the $B-$ band light curve maximum."
+ As shown in Figure 6.. the literature sample reveals a broad dispersion in Amys with mean values of (Anijsy)=0.94+0.31 mag and (Aus=0.74+0.27 mag.," As shown in Figure \ref{fig:template_delta_m15}, the literature sample reveals a broad dispersion in $\Delta m_{15}$ with mean values of $\langle\Delta m_{15,V}\rangle=0.94\pm 0.31$ mag and $\langle\Delta m_{15,R}\rangle =0.74\pm 0.27$ mag."
+ Thus. the R-band light curves decay more slowly than those in the V-band.," Thus, the $R-$ band light curves decay more slowly than those in the $V-$ band."
+ After dividing the literature sample into the SN sub-classes. we find that the Ausg values are statistically consistent: (Anis)=0.70.1 mag (Ib). (Anise)=0.90.5 mag (le). and (Ausp=0.7£0.1 mag (Ic-BL).," After dividing the literature sample into the SN sub-classes, we find that the $\Delta m_{15,R}$ values are statistically consistent: $\langle \Delta m_{15,R}\rangle =0.7\pm 0.1$ mag (Ib), $\langle \Delta
+m_{15,R}\rangle =0.9\pm 0.5$ mag (Ic), and $\langle \Delta
+m_{15,R}\rangle =0.7\pm 0.1$ mag (Ic-BL)."
+ We also note that the fast-fading Type Ile 11994I is a 2o outlier of the overall SN Ibe distribution., We also note that the fast-fading Type Ic 1994I is a $2\sigma$ outlier of the overall SN Ibc distribution.
+ Since the light curves for some of the SNe in our P60 sample are not as densely sampled as those of the literature sample. we use the V— and R—-band template light curves described above to fit the epoch and magnitude of the light curve maxima for all Gold and Silver P60 SNe.," Since the light curves for some of the SNe in our P60 sample are not as densely sampled as those of the literature sample, we use the $V-$ and $R-$ band template light curves described above to fit the epoch and magnitude of the light curve maxima for all Gold and Silver P60 SNe."
+ We use a \* minimization technique for the template fitting analysis and extract typical uncertainties on the epoch and apparent magnitude at maximum of about 0.5 (E1.0) d and +0.1 (40.2) mag. respectively for the Gold (Silver) SNe.," We use a $\chi^2$ minimization technique for the template fitting analysis and extract typical uncertainties on the epoch and apparent magnitude at maximum of about $\pm 0.5$ $\pm 1.0$ ) d and $\pm 0.1$ $\pm 0.2$ ) mag, respectively for the Gold (Silver) SNe."
+ For the Silver SNe. we fit only the R-band light curves since the V—band data are insufficient to estimate the parameters of the light curve maxima.," For the Silver SNe, we fit only the $R-$ band light curves since the $V-$ band data are insufficient to estimate the parameters of the light curve maxima."
+ The Gold and Silver P60 light curves are compared with the individual template fits in Figures 2. and 3.. respectively. and the estimated epoch and magnitude of maximum light for our sample are listed in Table 5..," The Gold and Silver P60 light curves are compared with the individual template fits in Figures \ref{fig:postage1} and \ref{fig:postage2}, respectively, and the estimated epoch and magnitude of maximum light for our sample are listed in Table \ref{tab:p60_lc}."
+ The Bronze sample do not have sufficient early data to reasonably constrain the epoch and magnitude of the light curve peak through template fitting., The Bronze sample do not have sufficient early data to reasonably constrain the epoch and magnitude of the light curve peak through template fitting.
+ In Figures 7 and 8. we display the absolute magnitude light curves for the P60 Gold and Silver samples after shifting them to the epoch of maximum light and correcting for Galactic extinction. (E(B—Wega;z0.09 mag for the P60 sample)., In Figures \ref{fig:moal_V} and \ref{fig:moal_R} we display the absolute magnitude light curves for the P60 Gold and Silver samples after shifting them to the epoch of maximum light and correcting for Galactic extinction $\langle E(B-V)_{\rm Gal}\rangle\approx 0.09$ mag for the P60 sample).
+ Similar to the case of the literature sample. we find a broad distribution in. My and Values that span ~2-3 magnitudes.," Similar to the case of the literature sample, we find a broad distribution in $M_{V,\rm peak}$ and $M_{R,\rm peak}$ values that span $\sim 2-3$ magnitudes."
+" Dividing pou,the Gold and Mg,Silver samples into the SNe Ib. Ic. and Ie-BL (8. 7. and 2 SNe. respectively). we report mean values of (Mj;=—17.420.5 mag (Ib). (Mepoak)=—]17.540.4 mag (Ic). and (Mgpea?=18.2+0.5 mag tIc-BL)."," Dividing the Gold and Silver samples into the SNe Ib, Ic, and Ic-BL (8, 7, and 2 SNe, respectively), we report mean values of $\langle M_{R,\rm peak}\rangle =-17.4\pm 0.5$ mag (Ib), $\langle M_{R,\rm peak}\rangle =-17.5\pm 0.4$ mag (Ic), and $\langle M_{R,\rm peak}\rangle =-18.2\pm 0.5$ mag (Ic-BL)."
+ These values are overall consistent with those of the literature sample., These values are overall consistent with those of the literature sample.
+ Adopting the fitted parameters for the light curve maxima. we next extract Aniyjsy and Anis values for our Gold and Silver SNe following the same procedure adopted for the literature sample refsee:templates)).," Adopting the fitted parameters for the light curve maxima, we next extract $\Delta m_{15,V}$ and $\Delta m_{15,R}$ values for our Gold and Silver SNe following the same procedure adopted for the literature sample \\ref{sec:templates}) )."
+ In Table 5. we list the resulting As values for the V— and R-bands. and in Figure 9. we compare their distributions with those from the literature sample.," In Table \ref{tab:p60_lc} we list the resulting $\Delta m_{15}$ values for the $V-$ and $R-$ bands, and in Figure \ref{fig:delta_m15} we compare their distributions with those from the literature sample."
+ We find mean values of (Anisv?20.78£0.11 mag and (Anise)=0.6130.13 mag. somewhat lower than those derived for the literature sample but consistent within the Io uncertainties.," We find mean values of $\langle\Delta
+m_{15,V}\rangle=0.78\pm 0.11$ mag and $\langle\Delta
+m_{15,R}\rangle=0.61\pm 0.13$ mag, somewhat lower than those derived for the literature sample but consistent within the $1\sigma$ uncertainties."
+ Similar to the case of the literature sample. we find no evidence for a statistically significant difference between the decay rates of the SNe sub-classes.," Similar to the case of the literature sample, we find no evidence for a statistically significant difference between the decay rates of the SNe sub-classes."
+ Finally. we comment on the rise time of the SNe Ibe in this sample. re. the time interval between the epoch of first detection and the epoch of maximum light.," Finally, we comment on the rise time of the SNe Ibc in this sample, i.e. the time interval between the epoch of first detection and the epoch of maximum light."
+ As shown in Figures 7 and 8.. SNe [be with broader light curves tend to have correspondingly longer rise times.," As shown in Figures \ref{fig:moal_V} and \ref{fig:moal_R}, SNe Ibc with broader light curves tend to have correspondingly longer rise times."
+ We find that several SNe in this P60 sample have R—band rise times exceeding 20 days and thus are comparable (even longer) than those observed for SNe la (Conleyetafl..2006:Haydeneral.2010).," We find that several SNe in this P60 sample have $R-$ band rise times exceeding 20 days and thus are comparable (even longer) than those observed for SNe Ia \citep{chh+06,hgk+10}."
+. This study emphasizes the dispersion in rise times of SNe Ibe light curves through dense. multi-band photometry of young explosions.," This study emphasizes the dispersion in rise times of SNe Ibc light curves through dense, multi-band photometry of young explosions."
+ By combining the Gold and Silver P60 samples with the literature sample. we derive mean values for the peak absolute magnitudes and light curve decline rates for a larger sample of SNe Ibe to search for statistical differences.," By combining the Gold and Silver P60 samples with the literature sample, we derive mean values for the peak absolute magnitudes and light curve decline rates for a larger sample of SNe Ibc to search for statistical differences."
+" Correcting for Galactic extinction only. we find Mypea,=—17.020.9 mag and Me=—17.3£0.7 mag."," Correcting for Galactic extinction only, we find $M_{V,\rm peak}=-17.0\pm 0.9$ mag and $M_{R,\rm peak}=-17.4\pm
+0.7$ mag."
+" Forcomparison. Lief recentlypea, reported that the typical unfiltered peak absolute magnitudes for SNe Ibe within dz:60 Mpe are My=—16.141.1 mag."," Forcomparison, \citet{llc+10} recently reported that the typical unfiltered peak absolute magnitudes for SNe Ibc within $d\approx 60$ Mpc are $M_{\rm peak}=-16.1\pm 1.1$ mag."
+" Restricting our sample to this same distance. we find a slightly higher value of Mx=—16.9-Ε0.6 mag i- the R-band. although we note that thesepea, values are consister= within the 1o uncertainties."," Restricting our sample to this same distance, we find a slightly higher value of $M_{R,\rm peak}=-16.9\pm 0.6$ mag in the $R-$ band, although we note that these values are consistent within the $1\sigma$ uncertainties."
+ Dividing our extended sample into the three spectroscoptο sub-classes (12 Ib. 10 Ie. and 5 Ie-BL). we report meaαυ] peak magnitudes with higher statistical confidence: (Mxja;= mag (Ib). (Mis?=—17.43:0.4 mag (Ic). and (Mg)=—18.30.6 mag (Ic-BL).," Dividing our extended sample into the three spectroscopic sub-classes (12 Ib, 10 Ic, and 5 Ic-BL), we report mean peak magnitudes with higher statistical confidence: $\langle M_{R,\rm peak}\rangle
+=-17.0\pm 0.7$ mag (Ib), $\langle M_{R,\rm peak}\rangle=-17.4\pm 0.4$ mag (Ic), and $\langle M_{R,\rm peak}\rangle =-18.3\pm 0.6$ mag (Ic-BL)."
+ This is in contrast. to the report of Lietal...(2010). that SNe Ib are marginally brighter than SNe Ie: we attribute this apparent difference to small number statistics., This is in contrast to the report of \citet{llc+10} that SNe Ib are marginally brighter than SNe Ic; we attribute this apparent difference to small number statistics.
+ While the light curve maxima for our sample of Bronze SNe are not well constrained. we are able to place strict lower limits on the peak absolute magnitudes by adopting the brightest data point the V— and R-bands and correcting for Galactic extinction.," While the light curve maxima for our sample of Bronze SNe are not well constrained, we are able to place strict lower limits on the peak absolute magnitudes by adopting the brightest data point the $V-$ and $R-$ bands and correcting for Galactic extinction."
+" These lower limits range from My,=—16.3 to 18.0 mag and Myi47—17.1 to 718.5 mag."," These lower limits range from $M_{V,\rm
+peak}=-16.3$ to $-18.0$ mag and $M_{R,\rm peak}=-17.1$ to $-18.5$ mag."
+ In Figure 10. we present the differential distribution of peak absolute magnitudes for our combined sample. including the lower limits derived for the Bronze P60 SNe.," In Figure \ref{fig:absmags_g} we present the differential distribution of peak absolute magnitudes for our combined sample, including the lower limits derived for the Bronze P60 SNe."
+ We next calculate mean decay rates for the combined sample of (Antsy)20.870.25 mag and (Anisjj=0.66£0.20 mag.," We next calculate mean decay rates for the combined sample of $\langle
+\Delta m_{15,V}\rangle=0.87\pm 0.25$ mag and $\langle\Delta
+m_{15,R}\rangle=0.66\pm 0.20$ mag."
+ Dividing into the sub-classes. we find (Ausp?=0.62£0.14 mag (Ib). (Ausg=0.73£0.27 mag (Ie). and (Anise)= mag (Ic-BL).," Dividing into the sub-classes, we find $\langle\Delta m_{15,R}\rangle=0.62\pm 0.14$ mag (Ib), $\langle\Delta m_{15,R}\rangle=0.73\pm 0.27$ mag (Ic), and $\langle\Delta m_{15,R}\rangle=0.60\pm 0.14$ mag (Ic-BL)."
+ A Kolmogorov-Smirnov (K-S) test on these distributions reveals an probability that SNe Ib and Ie are drawn from the same population. and the probability that SNe Ie and Ie-BL are the same is 86%..," A Kolmogorov-Smirnov (K-S) test on these distributions reveals an probability that SNe Ib and Ic are drawn from the same population, and the probability that SNe Ic and Ic-BL are the same is ."
+ Summing the SNe Ib and Ic distributions, Summing the SNe Ib and Ic distributions
+Preprint n=2>24 1+z1.35 Πα p=0.1—1gem? Ho Ho ὃς1-2«10°G for Zeeman splitting to play a role in the formation of the Fe Ha line (Loeb 2003).," $n=2\rightarrow 3$ $1+z\approx 1.35$ $\alpha$ $\rho \approx 0.1-1\,{\rm g\,cm}^{-3}$ $\alpha$ $\alpha$ $B<1-2\times 10^9\,{\rm G}$ for Zeeman splitting to play a role in the formation of the Fe $\alpha$ line (Loeb 2003)."
+ CBW calculated the equivalent width (EW) as a function of Fe column. Λι. for varying effective temperatures and pressures in the thin line-forming region.," CBW calculated the equivalent width (EW) as a function of Fe column, $N_{\rm Fe}$, for varying effective temperatures and pressures in the thin line-forming region."
+ To match the observed EW required. a total tron column of Nym1-3«10°’em was needed. higher than predicted by BCP by a factor of a few.," To match the observed EW required, a total iron column of $N_{\rm Fe} \approx 1-3 \times 10^{20}\,{\rm cm}^{-2}$ was needed, higher than predicted by BCP by a factor of a few."
+ In this Letter. we rotationally broaden the energy and angle dependent line profiles of CBW and compare them to the observations.," In this Letter, we rotationally broaden the energy and angle dependent line profiles of CBW and compare them to the observations."
+ We explain our method of calculating the rotational broadening in $2 and show that the Schwarzschild+Doppler (S+D) approximation (Miller Lamb 1998: Braje. Romani. Rauch 2000) is extremely accurate in reproducing the line profiles from the full general relativistic calculation. even up to 300 Hz.," We explain our method of calculating the rotational broadening in 2 and show that the Schwarzschild+Doppler (S+D) approximation (Miller Lamb 1998; Braje, Romani, Rauch 2000) is extremely accurate in reproducing the line profiles from the full general relativistic calculation, even up to 300 Hz."
+ We fit the data in $3 and show that the intrinsically broad line profile prohibits any meaningful constraint on the NS radius if the 44.7 Hz burst oscillation seen by Villarreal Strohmayer (2004) is the spin frequency., We fit the data in 3 and show that the intrinsically broad line profile prohibits any meaningful constraint on the NS radius if the 44.7 Hz burst oscillation seen by Villarreal Strohmayer (2004) is the spin frequency.
+ At slow spins or at low inclinations. the line profile shape ts set by both rotational broadening and fine structure splitting. making 1t absolutely necessary to use realistic line profiles when modeling the lines in oor any other favorably oriented. (and hence detectable) rapidly rotating NS.," At slow spins or at low inclinations, the line profile shape is set by both rotational broadening and fine structure splitting, making it absolutely necessary to use realistic line profiles when modeling the lines in or any other favorably oriented (and hence detectable) rapidly rotating NS."
+ We confirm the importance of the intrinsic line. profile in. determining. NS. parameters. by constrasting it with a simple line profile., We confirm the importance of the intrinsic line profile in determining NS parameters by constrasting it with a simple line profile.
+ We close 33 by accurately determining the redshift and the amount of Fe in the upper atmosphere., We close 3 by accurately determining the redshift and the amount of Fe in the upper atmosphere.
+ We conclude in $4 by summarizing our findings and discussing the likelihood that lines like that seen in wwill be found in the more rapidly rotating NSs., We conclude in 4 by summarizing our findings and discussing the likelihood that lines like that seen in will be found in the more rapidly rotating NSs.
+ We compute the impact of rotational broadening on CBW's intrinsic line profiles using two independent methods: a fully relativistic method that employs ray-tracing in the spacetime of a rotating NS and an approximate method which retains the most important relativistic effects to lowest order in v/c., We compute the impact of rotational broadening on CBW's intrinsic line profiles using two independent methods: a fully relativistic method that employs ray-tracing in the spacetime of a rotating NS and an approximate method which retains the most important relativistic effects to lowest order in $v/c$.
+ In the fully relativistic approach. the spacetime metric for rotating NSs with various spin frequencies. masses and equations of state are computed numerically asdescribedin Cadeau.LeahyMorsink(20052).," In the fully relativistic approach, the spacetime metric for rotating NSs with various spin frequencies, masses and equations of state are computed numerically asdescribedin Cadeau,LeahyMorsink(2005a)."
+ Once the (non-spherical), Once the (non-spherical)
+"So now the free parameters are reduced to 6, i.e. the density ny and temperature 7, of hot plasma, the angle ας and pitch angle distribution slope N of the loss-cone, the radius ube of one flux tube and the size of the emitting region A.","So now the free parameters are reduced to 6, i.e. the density $n_{\rm h}$ and temperature $T_{\rm h}$ of hot plasma, the angle $\alpha_{\rm c}$ and pitch angle distribution slope $N$ of the loss-cone, the radius $r_{\rm tube}$ of one flux tube and the size of the emitting region $A$."
+" The functions of the parameters are: ny and Τι can control the change of flux density dramatically, «, and N can control the change of flux density gently, A can control both flux density and time duration of the radio pulse which can be constrained by observations, and rupe describes the radius of the flux tube in the active region."," The functions of the parameters are: $n_{\rm h}$ and $T_{\rm h}$ can control the change of flux density dramatically, $\alpha_{\rm c}$ and $N$ can control the change of flux density gently, $A$ can control both flux density and time duration of the radio pulse which can be constrained by observations, and $r_{\rm tube}$ describes the radius of the flux tube in the active region."
+ The free parameters can be constrained by making our simulations match up to observations., The free parameters can be constrained by making our simulations match up to observations.
+ Fig., Fig.
+ 6 shows the comparison of our simulation with observations (see Hallinan et al., \ref{fig_com} shows the comparison of our simulation with observations (see Hallinan et al.
+ 2010 for further details)., 2010 for further details).
+ The observation is taken from 2008 May 19 (UT) at 4725 MHz for the M8.5V dwarf TVLM 513., The observation is taken from 2008 May 19 (UT) at 4725 MHz for the M8.5V dwarf TVLM 513.
+ This means that the magnetic field strength is 1652 G because of the cyclotron frequency., This means that the magnetic field strength is 1652 G because of the cyclotron frequency.
+" We adopt fj/f.=0.1, implying the density of cold plasma in the emitting region is ~2.64x10? cm?."," We adopt $f_{\rm p}/f_{\rm c}=0.1$, implying the density of cold plasma in the emitting region is $\sim$ $\times10^{9}$ $^{-3}$ ."
+" The size of the radio-emitting region is constrained by the time duration of the two pulses (Ax (54.75Aty)), implying 1620x1620 km? for the first pulse and 2060x2060 km? for the second pulse."," The size of the radio-emitting region is constrained by the time duration of the two pulses $A\approx(54.75\Delta t)^{2}$ ), implying $\times$ 1620 $^{2}$ for the first pulse and $\times$ 2060 $^{2}$ for the second pulse."
+ The radius of each flux tube is taken as 180 km arbitrarily because we are not able to get any information on its value from the observed radio light curve., The radius of each flux tube is taken as 180 km arbitrarily because we are not able to get any information on its value from the observed radio light curve.
+" We take ας=30° and N=5 for the first pulse, ας=35? and Ν=2 for the second pulse."," We take $\alpha_{\rm c}=30^{\circ}$ and $N=5$ for the first pulse, $\alpha_{\rm c}=35^{\circ}$ and $N=2$ for the second pulse."
+" 7j and ny are same for the both pulses, equal to 107 K and 1.25x10? cm? respectively."," $T_{\rm h}$ and $n_{\rm h}$ are same for the both pulses, equal to $10^{7}$ K and $1.25\times10^{5}$ $^{-3}$ respectively."
+ The simulated light curve reproduces the observation., The simulated light curve reproduces the observation.
+" However, this does not mean that the above values for the different parameters are unique."," However, this does not mean that the above values for the different parameters are unique."
+" In order to compare our simulation with the observed two pulses, we take a flat distribution for ας and N in the range of 10?—50? and 1—6 for the both pulses, m,: (1.25-5)x105 cm, Ty: (1—4.5)x107 K for the first pulse and nj: (1.25—4.5)x10° cm™, Τι: (1—3)x107 K for the second pulse (see inset panel in Fig. 6))."," In order to compare our simulation with the observed two pulses, we take a flat distribution for $\alpha_{\rm c}$ and $N$ in the range of $10^{\circ}-50^{\circ}$ and $1-6$ for the both pulses, $n_{\rm h}:$ $(1.25-5)\times10^{5}$ $^{-3}$, $T_{\rm h}:$ $(1-4.5)\times10^{7}$ K for the first pulse and $n_{\rm h}:$ $(1.25-4.5)\times10^{5}$ $^{-3}$, $T_{\rm h}:$ $(1-3)\times10^{7}$ K for the second pulse (see inset panel in Fig. \ref{fig_com}) )."
+ We note that the size of the emitting region depends on the latitude 0 in our simulation (see $4.5))., We note that the size of the emitting region depends on the latitude $\theta$ in our simulation (see \ref{sec_environment}) ).
+ Changing the latitude from 30? (the value for the simulation) to 70? gives a value for the size of 640x640 km? for the first pulse and 814x814, Changing the latitude from $30^{\circ}$ (the value for the simulation) to $70^{\circ}$ gives a value for the size of $\times$ 640 $^{2}$ for the first pulse and $\times$ 814
+ Changing the latitude from 30? (the value for the simulation) to 70? gives a value for the size of 640x640 km? for the first pulse and 814x814, Changing the latitude from $30^{\circ}$ (the value for the simulation) to $70^{\circ}$ gives a value for the size of $\times$ 640 $^{2}$ for the first pulse and $\times$ 814$
+ Changing the latitude from 30? (the value for the simulation) to 70? gives a value for the size of 640x640 km? for the first pulse and 814x814, Changing the latitude from $30^{\circ}$ (the value for the simulation) to $70^{\circ}$ gives a value for the size of $\times$ 640 $^{2}$ for the first pulse and $\times$ 814$^
+ Changing the latitude from 30? (the value for the simulation) to 70? gives a value for the size of 640x640 km? for the first pulse and 814x814, Changing the latitude from $30^{\circ}$ (the value for the simulation) to $70^{\circ}$ gives a value for the size of $\times$ 640 $^{2}$ for the first pulse and $\times$ 814$^{
+ Changing the latitude from 30? (the value for the simulation) to 70? gives a value for the size of 640x640 km? for the first pulse and 814x814k, Changing the latitude from $30^{\circ}$ (the value for the simulation) to $70^{\circ}$ gives a value for the size of $\times$ 640 $^{2}$ for the first pulse and $\times$ 814$^{2
+ Changing the latitude from 30? (the value for the simulation) to 70? gives a value for the size of 640x640 km? for the first pulse and 814x814km, Changing the latitude from $30^{\circ}$ (the value for the simulation) to $70^{\circ}$ gives a value for the size of $\times$ 640 $^{2}$ for the first pulse and $\times$ 814$^{2}
+ Changing the latitude from 30? (the value for the simulation) to 70? gives a value for the size of 640x640 km? for the first pulse and 814x814km?, Changing the latitude from $30^{\circ}$ (the value for the simulation) to $70^{\circ}$ gives a value for the size of $\times$ 640 $^{2}$ for the first pulse and $\times$ 814$^{2}$
+On the other hand. ifthe MWefr was ionized internally. different parts of it would have reached the critical collapse fraction of L1-cooling halos necessary for reionization at dilerent times due to their respective overdensities (Barkana&Loeb 2004).,"On the other hand, if the MWgfr was ionized internally, different parts of it would have reached the critical collapse fraction of -cooling halos necessary for reionization at different times due to their respective overdensities \citep{BL04}."
+. Thus. there should be fluctuations in the redshift’ of reionization. within the MW. itself. about. the mean at μα. an elleet that has not been considered. in the literature.," Thus, there should be fluctuations in the redshift of reionization within the MW itself about the mean at $\bar{z}_{\rm rei}$, an effect that has not been considered in the literature."
+ Each subhalo progenitor will photoionize and reheat its own gas content at the time when it has produced Νταν lonizing photons per hydrogen atom., Each subhalo progenitor will photoionize and reheat its own gas content at the time when it has produced $N_{\rm clump}$ ionizing photons per hydrogen atom.
+ This condition is met when where ο) gives the ratio of Nenu in the overdense progenitor to the average Nou ln the universe., This condition is met when where $g(\delta)$ gives the ratio of $N_{\rm clump}$ in the overdense progenitor to the average $N_{\rm clump}$ in the universe.
+ The factor (Lofiie) selects only those photons that do not escape from. the overdense progenitor and. can be used to. photoionize it., The factor $\left(1-f_{esc}\right)$ selects only those photons that do not escape from the overdense progenitor and can be used to photoionize it.
+ Equation 3. tells us how much stellar mass to assign aprogenitor of mass Ady., Equation \ref{eq:MsIn} tells us how much stellar mass to assign aprogenitor of mass $M_{\rm halo}$.
+" Lach progenitor. A. has its own ALAd, progenitors. D. that form stars at. redshift tig;Doce before⋅ system (X collapses at ty."," Each progenitor, A, has its own $M>M_4$ progenitors, B, that form stars at redshift $z'_{\rm rei}>\bar{z}_{\rm rei}$, before system A collapses at $\bar{z}_{\rm rei}$."
+ Darvonic. material from the region that formed the A progenitor not converted into stars al ETel is photoionized before it can be incorporated into A at ty and cannot form stars unless progenitor A eventually reaches a mass of Als. the mass corresponding to a virial temperature of Z5 or a circular velocity of V5=52kms+.," Baryonic material from the region that formed the A progenitor not converted into stars at $z'_{\rm rei}$ is photoionized before it can be incorporated into A at $\bar{z}_{\rm rei}$ and cannot form stars unless progenitor A eventually reaches a mass of $M_5$, the mass corresponding to a virial temperature of $T_5$ or a circular velocity of $V_5=52\,\kms$."
+" However. we expect most of the eas in D. progenitors with AJMj at , ⋅ ∶↓∣⇁↓↿∪⇂∪↓⋅⊔↓⊳∖⇂⋜⊔⋅⊳∖⊳"," However, we expect most of the gas in B progenitors with $M>M_4$ at $z'_{\rm rei}$ to form stars."
+ Accounting For the inside-out morphology of the MWefr rejonization provides a natural explanation for the low values of fgio found by several studies (Macauetal.2008b:]xoposovetal.2009:Busha 2009).. whieh assumed a single value for the reionization redshift of the entire NWefr.," Accounting for the inside-out morphology of the MWgfr reionization provides a natural explanation for the low values of $f_{\rm \HI, ex}$ found by several studies \citep{Madau08b,Koposov09,Busha09}, which assumed a single value for the reionization redshift of the entire MWgfr."
+ Since stellar mass is a linear function of halo mass in both equations 2 and 3.. both morphologies fit the same observed LF when V faigto/(1foe) ds of order unity at σι=11. we find that gis70.01.," Since stellar mass is a linear function of halo mass in both equations \ref{eq:MsEx} and \ref{eq:MsIn}, both morphologies fit the same observed LF when If $f_{\rm \HI}\,g(\delta)/\left(1-f_{\rm esc}\right)$ is of order unity at $\bar{z}_{\rm rei}=11$, we find that $f_{\rm \HI,
+ex}\approx0.01$."
+" We can understand. this physically bv considering how. at the redshift of star formation suppression. a lower cllicicney is equivalent to a smaller amount of matter in collapsed. objects. (smaller. Adj, in Eq. 2))"," We can understand this physically by considering how, at the redshift of star formation suppression, a lower efficiency is equivalent to a smaller amount of matter in collapsed objects (smaller $M_{\rm halo}$ in Eq. \ref{eq:MsEx}) )"
+ in terms of calculating how much stellar mass is produced., in terms of calculating how much stellar mass is produced.
+ We see from this that. for a given average redshift of reionization in the Al\Welr. the insicle-out ane outside-in morphologics fit the luminosity function equally well but result in cillerent interpretations of the fitting parameters.," We see from this that, for a given average redshift of reionization in the MWgfr, the inside-out and outside-in morphologies fit the luminosity function equally well but result in different interpretations of the fitting parameters."
+ We assume that cosmic reionization anc the photoheating of the MWefr. completes rapidly alter 2j., We assume that cosmic reionization and the photoheating of the MWgfr completes rapidly after $\bar{z}_{\rm rei}$.
+ Subsequently. only halos that have accumulated a mass of Mzz1 are able to hold onto their gas long enough for it to form further stars via lli--cooling.," Subsequently, only halos that have accumulated a mass of $M_5>1$ are able to hold onto their gas long enough for it to form further stars via -cooling."
+ We assume that this star formation is quenched when the subhalo begins to interact with the MW. host., We assume that this star formation is quenched when the subhalo begins to interact with the MW host.
+" The merger histories of cach surviving MM. subhalo at 2=0 show only eight. subhalos that achieved a maximum circular velocity greater than \=, at some point in their histories.", The merger histories of each surviving MW subhalo at $z=0$ show only eight subhalos that achieved a maximum circular velocity greater than $V_5$ at some point in their histories.
+ We determine the redshift ab which the maximum circular. velocity. reaches its peak., We determine the redshift at which the maximum circular velocity reaches its peak.
+ The subsequent reduction is due to tidal interactions during infall into the AIW host with both neighboring chwarls and the host itself (Ixravtsovctal.2004). that we assume coincide with the quenching of star formation through ram pressure., The subsequent reduction is due to tidal interactions during infall into the MW host with both neighboring dwarfs and the host itself \citep{Kravtsov04} that we assume coincide with the quenching of star formation through ram pressure.
+ In agreement with Ixoposovetal.(2009).. we find that this mass loss tvpicallv occurs around +=24 (in five of our eight subhalos) but can happen as carly as 2—8 and as late as z=1.6.," In agreement with \citet{Koposov09}, we find that this mass loss typically occurs around $z=2-4$ (in five of our eight subhalos) but can happen as early as $z=8$ and as late as $z=1.6$."
+ We add an additional mass of fsfiMiis worth of stars minus the mass of any stars formed in either of the first two epochs to each subhalo whose peak circular velocity exceeds 3$," We add an additional mass of $f_5\,f_b\,M_{\rm max}$ worth of stars minus the mass of any stars formed in either of the first two epochs to each subhalo whose peak circular velocity exceeds $V_5$."
+ λογο. fsods an elliciency parameter for star formation during this epoch that takes into account how much of the hot gas is cooled to form stars and how much of the eas mass is removed during infall.," Here, $f_5$ is an efficiency parameter for star formation during this epoch that takes into account how much of the hot gas is cooled to form stars and how much of the gas mass is removed during infall."
+ Ady... is the virial mass corresponding to the maximum value of the circular velocity for each subhalo., $M_{max}$ is the virial mass corresponding to the maximum value of the circular velocity for each subhalo.
+ We calculate the age of these new stars from the redshift at which the subhalo attains its maximum circular velocity and use this age to calculate the luminosity of cach solar mass of stars from Druzual&Charlot(2003)., We calculate the age of these new stars from the redshift at which the subhalo attains its maximum circular velocity and use this age to calculate the luminosity of each solar mass of stars from \citet{BC03}.
+. The final episode of star formation that we include in our mocel occurs in the last 10Gyr since z=1.6.," The final episode of star formation that we include in our model occurs in the last $10\,\rm {Gyr}$ since $z=1.6$."
+ Since we are interested primarily in the physies of the early. universe. we do not attempt to develop here a detailed. model for the production of these stars but presume that metal cooling is involved and use the observations of Orbanetal.(2008) to adel additional voung stars in à stochastic wav.," Since we are interested primarily in the physics of the early universe, we do not attempt to develop here a detailed model for the production of these stars but presume that metal cooling is involved and use the observations of \citet{Orban08} to add additional young stars in a stochastic way."
+ For each of the classical. pre-SDSS MW satellite. Orbanal.(2008). have measured. floes. the fraction of stellar mass produced in the last I0vr. and 7. the masseaveraged stellar age.," For each of the classical, pre-SDSS MW satellite, \citet{Orban08} have measured $f_{10G}$, the fraction of stellar mass produced in the last $10\,\rm{Gyr}$, and $\tau$, the mass-averaged stellar age."
+ They find that the metallicities of the ultra-aint. SDSS satellites are so low that no star formation is expected. for these in the last 10yr.," They find that the metallicities of the ultra-faint, SDSS satellites are so low that no star formation is expected for these in the last $10\,\rm{Gyr}$."
+ As we will show. subhalos that achieve a maximum circular velocity of at east V5 are associated only with the classical MW. satellites stucied by Orbanetal.(2008)... while stars that. formed »ior to reionization populate both the classical ancl SDSS dSphs.," As we will show, subhalos that achieve a maximum circular velocity of at least $V_5$ are associated only with the classical MW satellites studied by \citet{Orban08}, while stars that formed prior to reionization populate both the classical and SDSS dSphs."
+ Therefore. we set. fioc;=0 for subhalo progenitors hat only formed. stars before reionization to be consistent with these metallicity observations.," Therefore, we set $f_{10G}=0$ for subhalo progenitors that only formed stars before reionization to be consistent with these metallicity observations."
+ However. setting this constraint only for progenitors with L»-induced cooling docs not significantly change our results.," However, setting this constraint only for progenitors with $_2$-induced cooling does not significantly change our results."
+ The cumulative probability distribution of fiuc; ds approximately linear with a linear [fit reaching unity at fioc:=0.89 (Orbanetal.2008)., The cumulative probability distribution of $f_{10G}$ is approximately linear with a linear fit reaching unity at $f_{10G}=0.89$ \citep{Orban08}.
+. Thus. for each subhalo with stars that formed after reionization. as outlined above. we select a value lor fioc; at random from a Lat probability distribution within the interval 0.00.0.89].," Thus, for each subhalo with stars that formed after reionization, as outlined above, we select a value for $f_{10G}$ at random from a flat probability distribution within the interval $\left[0.00,0.89\right]$."
+ Phe mass of stars produced since >=1.6 is. therefore. given by fioc/(1—fioc) times the mass of stars produced. from. the mechanisms outlined in the previous three subsections.," The mass of stars produced since $z=1.6$ is, therefore, given by $f_{10G}/(1-f_{10G})$ times the mass of stars produced from the mechanisms outlined in the previous three subsections."
+ If we pretend that all of the stars formed in the ATW satellites prior to 2=1.6 have an age of I4 Gyr. we can use the measured values of fioc; and 7 t0 estimate the mass- average age of those stars formed in the last 10 Gwe.," If we pretend that all of the stars formed in the MW satellites prior to $z=1.6$ have an age of $14\,\rm{Gyr}$ , we can use the measured values of $f_{10G}$ and $\tau$ to estimate the mass-weighted average age of those stars formed in the last $10\,\rm{Gyr}$ ."
+"their previous frames taken on 2008 December 25.559 UT and 2009 August 07.737 UT with limiting magnitudes of 20.0 and 19.4, respectively.","their previous frames taken on 2008 December 25.559 UT and 2009 August 07.737 UT with limiting magnitudes of 20.0 and 19.4, respectively."
+ The source was confirmed by these authors on 2009 August 18.696., The source was confirmed by these authors on 2009 August 18.696.
+ But later they withdraw the interpretation as possible nova as it did not change brightness., But later they withdraw the interpretation as possible nova as it did not change brightness.
+ The source was identified by us with a star number 40671 in the photometric survey of variable stars in 333 carried out by Hartmanetal.(2006)., The source was identified by us with a star number 40671 in the photometric survey of variable stars in 33 carried out by \citet{h1}.
+ We will name it hereafter [HSB2006] 40671., We will name it hereafter [HSB2006] 40671.
+ The survey observations were carried out between 2003 August and 2005 January., The survey observations were carried out between 2003 August and 2005 January.
+ The authors classify the star as a Long Period Variable (LPV)., The authors classify the star as a Long Period Variable (LPV).
+ During their observations (filled circles in Fig., During their observations (filled circles in Fig.
+" 3), the star was red with R—Iz&2.7 and &4.1 mag at maximum and minimum brightness, respectively."," 3), the star was red with $R - I \approx 2.7$ and $\approx 4.1$ mag at maximum and minimum brightness, respectively."
+ The star became extremely red during the registered minimum (R=25 mag)., The star became extremely red during the registered minimum $R \ga 25$ mag).
+ We have carried out spectral observations with the Russian 6-m BTA telescope using the SCORPIO spectrograph (Afanasiev&Moiseev2005)., We have carried out spectral observations with the Russian 6-m BTA telescope using the SCORPIO spectrograph \citep{a1}.
+. The spectrum has been taken on 2009 October 9.874 UT in a range with spectral resolution of 13A., The spectrum has been taken on 2009 October 9.874 UT in a range with spectral resolution of 13.
+. The spectroscopy was followed by images on the same date., The spectroscopy was followed by images on the same date.
+ We found the star at a level of 18.8 mag in the band., We found the star at a level of 18.8 mag in the band.
+ One more photometric observation has been taken with the 1-m Zeiss telescope of the Special Astrophysical Observatory (RAS) on 2009 October 20.9 UT., One more photometric observation has been taken with the 1-m Zeiss telescope of the Special Astrophysical Observatory (RAS) on 2009 October 20.9 UT.
+ We have searched for additional images of this star in Internet archives of different observatories., We have searched for additional images of this star in Internet archives of different observatories.
+" Dozens of images were found, most of them come from medium-size telescopes equipped with wide field mosaic CCD cameras."," Dozens of images were found, most of them come from medium-size telescopes equipped with wide field mosaic CCD cameras."
+" The star position is well placed on the frames of the Calar Alto Observatory taken in 1989 and 1998 with the 80/120-cm Schmidt telescope, on the frames taken in 2001 and 2002 in BVRI bands with the 8.22m SUBARU telescope of the National Astronomical Observatory of Japan."," The star position is well placed on the frames of the Calar Alto Observatory taken in 1989 and 1998 with the 80/120-cm Schmidt telescope, on the frames taken in 2001 and 2002 in $BVRI$ bands with the 8.2-m SUBARU telescope of the National Astronomical Observatory of Japan."
+" One V band observation from the 4.2-m William Herschel Telescope, and three SDSS r’ band observations from the 2.54-m Isaac Newton Telescope were used as well, both telescopes belong to the Isaac Newton Group of Telescopes at the Roque de Los Muchachos Observatory at La Palma, Spain."," One $V$ band observation from the 4.2-m William Herschel Telescope, and three SDSS $r'$ band observations from the 2.54-m Isaac Newton Telescope were used as well, both telescopes belong to the Isaac Newton Group of Telescopes at the Roque de Los Muchachos Observatory at La Palma, Spain."
+ We also apply one J band measurement obtained from the 4.0-m Mayall telescope image taken in 2001 and published by Masseyetal.(2006).., We also apply one $I$ band measurement obtained from the 4.0-m Mayall telescope image taken in 2001 and published by \citet{m3}.
+" Archival images from the 8.2-m SUBARU telescope, the 4.2-m William Herschel Telescope, and the 2.54-m Isaac Newton Telescope were downloaded as raw FITS files and then processed in the same way as our own 6-m BTA and Zeiss telescope images."," Archival images from the 8.2-m SUBARU telescope, the 4.2-m William Herschel Telescope, and the 2.54-m Isaac Newton Telescope were downloaded as raw FITS files and then processed in the same way as our own 6-m BTA and Zeiss telescope images."
+ Standard reduction procedures, Standard reduction procedures
+"Tw decreases with decreasing L,, as in many explosion models, only the early stage of the neutrino-driven windwith L,~10°? erg s~! may be relevant to the high Twt,9=1.5—3 (see,e.g.,Arconesetal.2007) that is needed for efficient vp-processing 4.1).","$T_\mathrm{wt}$ decreases with decreasing $L_\nu$, as in many explosion models, only the early stage of the neutrino-driven windwith $L_\nu
+\sim 10^{52}$ erg $^{-1}$ may be relevant to the high $T_\mathrm{wt, 9}
+= 1.5-3$ \citep[see, e.g.,][]{Arco2007} that is needed for efficient $\nu
+$ p-processing 4.1)."
+" The mass of the proto-neutron star My, can be somewhat different from its canonical value of 1.4Mo, depending on its progenitor mass."," The mass of the proto-neutron star $M_\mathrm{ns}$ can be somewhat different from its canonical value of $1.4\, M_\odot$, depending on its progenitor mass."
+" In this subsection, we examine the nucleosynthesis calculations with My,=1.2,1.4,1.6, 1.8, and 2.0 Mo, while L, and Y,,9 are kept to be their fiducial values of 10°? erg s~! and 0.600."," In this subsection, we examine the nucleosynthesis calculations with $M_\mathrm{ns} = 1.2, 1.4, 1.6, 1.8$ , and $2.0\,
+M_\odot$ , while $L_\nu$ and $Y_\mathrm{e, 9}$ are kept to be their fiducial values of $10^{52}$ erg $^{-1}$ and 0.600."
+" For each case, the fiducial value of Tyt,9=2.19 is obtained by adjusting r«« (from 16th to 20th lines in Table 1) as in$ 4.2."," For each case, the fiducial value of $T_\mathrm{wt, 9} = 2.19$ is obtained by adjusting $r_\mathrm{wt}$ (from 16th to 20th lines in Table 1) as in 4.2."
+ We find clear correlation between an increase of Mj and an increasinga efficiency of vp-processing in Figure 4 and Table 1., We find a clear correlation between an increase of $M_\mathrm{ns}$ and an increasing efficiency of $\nu$ p-processing in Figure 4 and Table 1.
+" This is due to a larger S and a smaller τι for a larger Μις (e.g.,Otsukietal.2000;Wanajoal. 2001),, both of which help to increase Y,/Yp and thus An."," This is due to a larger $S$ and a smaller $\tau_1$ for a larger $M_\mathrm{ns}$ \citep[e.g.,][]{Otsu2000, Wana2001}, both of which help to increase $Y_\mathrm{p}/Y_\mathrm{h}$ and thus $\Delta_\mathrm{n}$."
+" This means that a more massive progenitor (up to Mo, which forms a neutron is favored for the vp-process, given that all the other star)parameters are the same."," This means that a more massive progenitor (up to $\sim 30\, M_\odot$ , which forms a neutron star) is favored for the $\nu$ p-process, given that all the other parameters are the same."
+" In reality, however, other factors, such as the evolutions of Ly, r««, and Y, should be dependent on the progenitor mass (e.g.,Arconesetal.2007),, which prevents us from drawing any firm conclusions."," In reality, however, other factors, such as the evolutions of $L_\nu$, $r_\mathrm{wt}$, and $Y_\mathrm{e}$ should be dependent on the progenitor mass \citep[e.g.,][]{Arco2007}, which prevents us from drawing any firm conclusions."
+" It should be emphasized, however, that the outflow with a typical mass of May,=1.4 can already provide physical conditions sufficient for Moproducing the p-nuclei up to A~110."," It should be emphasized, however, that the outflow with a typical mass of $M_\mathrm{ns} = 1.4\, M_\odot$ can already provide physical conditions sufficient for producing the p-nuclei up to $A \sim 110$."
+" 'The electron fraction Y, is obviously one of the most important ingredients in the rp-process as it controls the proton-richness in the ejecta.", The electron fraction $Y_\mathrm{e}$ is obviously one of the most important ingredients in the $\nu$ p-process as it controls the proton-richness in the ejecta.
+ Recent hydrodynamical studies with elaborate neutrino transport indicate that Y. exceeds 0.5 and increases up to ~0.6 during the neutrino-driven wind phase (Fischeretal.2010;Hüdepohletal.2010).," Recent hydrodynamical studies with elaborate neutrino transport indicate that $Y_\mathrm{e}$ exceeds 0.5 and increases up to $\sim 0.6$ during the neutrino-driven wind phase \citep{Fisc2010, Hued2010}."
+. It should be noted that Y. substantially decreases from its initial value owing to the neutrino effects 4.1)., It should be noted that $Y_\mathrm{e}$ substantially decreases from its initial value owing to the neutrino effects 4.1).
+" In our standard model, the value decreases from Y,9=0.600 (at Το=9) to Y.a=0.550 at the onset of the vp-process (Το=3)."," In our standard model, the value decreases from $Y_\mathrm{e, 9} = 0.600$ (at $T_9 = 9$ ) to $Y_\mathrm{e, 3} = 0.550$ at the onset of the $\nu$ p-process $T_9 = 3$ )."
+" However, these neutrino effects would highly dependent on the neutrino luminosities and energies of electron and anti-electron neutrinos assumed in this study."," However, these neutrino effects would highly dependent on the neutrino luminosities and energies of electron and anti-electron neutrinos assumed in this study."
+" In this subsection, therefore, we take the value at the onset of the vp-process, Yo,3, as a reference, rather than the initial value Y,9. Figure 5 and Table 1 (the last (6 lines) show the nucleosynthetic results for Yo3= and 0.655 (see Table 1 0.523,for their0.550, initial0.576, values Y;,9)."," In this subsection, therefore, we take the value at the onset of the $\nu
+$ p-process, $Y_\mathrm{e, 3}$, as a reference, rather than the initial value $Y_\mathrm{e, 9}$ Figure 5 and Table 1 (the last 6 lines) show the nucleosynthetic results for $Y_\mathrm{e, 3} = 0.523, 0.550, 0.576, 0.603, 0.629$ , and 0.655 (see Table 1 for their initial values $Y_\mathrm{e, 9}$ )."
+" The other parameters My, Ly, and ryt (and thus Ty) are kept to be their fiducial values (1st line in Table 1)."," The other parameters $M_\mathrm{ns}$, $L_\nu$ , and $r_\mathrm{wt}$ (and thus $T_\mathrm{wt}$ ) are kept to be their fiducial values (1st line in Table 1)."
+" We find a great impact of the Y, variation; an increase of only AY.3~0.03 leads to a 10-unit increase of Ajax, while fmax is similar for >0.550."," We find a great impact of the $Y_\mathrm{e}$ variation; an increase of only $\Delta Y_\mathrm{e, 3} \sim 0.03$ leads to a 10-unit increase of $A_\mathrm{max}$, while $f_\mathrm{max}$ is similar for $Y_\mathrm{e, 3} > 0.550$."
+" This is due to the larger Y,/Yn (at Tg=3) Y.for a larger Yo3, leading the larger A,, despite the same np, (Table 1)."," This is due to the larger $Y_\mathrm{p}/Y_\mathrm{h}$ (at $T_9 = 3$ ) for a larger $Y_\mathrm{e,
+ 3}$, leading the larger $\Delta_{n}$ despite the same $n_{\bar{\nu}_\mathrm{e}}$ (Table 1)."
+" In order to elucidate the effect of Y, in more detail, the production factor f for each p-nucleus is drawn in Figure 6 as a function of Ye3, where Mys, Ly, and ryt are kept to be their fiducial values."," In order to elucidate the effect of $Y_\mathrm{e}$ in more detail, the production factor $f$ for each p-nucleus is drawn in Figure 6 as a function of $Y_\mathrm{e, 3}$ , where $M_\mathrm{ns}$, $L_\nu$, and $r_\mathrm{wt}$ are kept to be their fiducial values."
+" Each element is color coded with the solid, dashed, and long-dashed lines for the lightest, second-lightest, and third-lightest (11°Sn is only the case) isotopes, respectively (see 1st column of Table 4 for the list of p-nuclei)."," Each element is color coded with the solid, dashed, and long-dashed lines for the lightest, second-lightest, and third-lightest $^{115}$ Sn is only the case) isotopes, respectively (see 1st column of Table 4 for the list of p-nuclei)."
+ We find in the top panel of Figure 6 that the p-nuclei up to A=108 (19504) take the maximum production factors between =0.53 and 0.60.," We find in the top panel of Figure 6 that the p-nuclei up to $A = 108$ $^{108}$ Cd) take the maximum production factors between $Y_\mathrm{e, 3} = 0.53$ and 0.60."
+" Given the maximum Y,,3 to be ~0.6Y.,3 according to some recent hydrodynamic results (e.g.,Fischeret this implies that the maximum mass number of 2010),the p-nuclei produced by the vp-process is A~110."," Given the maximum $Y_\mathrm{e, 3}$ to be $\sim 0.6$ according to some recent hydrodynamic results \citep[e.g.,][]{Fisc2010, Hued2010}, this implies that the maximum mass number of the p-nuclei produced by the $\nu$ p-process is $A \sim
+110$."
+" In principle, the heavier p-nuclei can be synthesized if the matter is more proton-rich than =0.6."," In principle, the heavier p-nuclei can be synthesized if the matter is more proton-rich than $Y_\mathrm{e, 3} = 0.6$."
+ The middle panel of Figure 6 shows that Y.3the production factors of the p-nuclei from A=113 (!!?In) up to A=138 (19509) are maximal between Y.3=0.61 and 0.63.," The middle panel of Figure 6 shows that the production factors of the p-nuclei from $A =
+113$ $^{113}$ In) up to $A = 138$ $^{138}$ Ce) are maximal between $Y_\mathrm{e, 3} = 0.61$ and 0.63."
+" Furthermore, '44Sm and Gd reach the maximum production factors at Y,4=0.64 and 0.66, respectively (bottom panel in Figure 6)."," Furthermore, $^{144}$ Sm and $^{152}$ Gd reach the maximum production factors at $Y_\mathrm{e, 3} =
+0.64$ and 0.66, respectively (bottom panel in Figure 6)."
+ The end point of the vp- appears to be at A~180 (199 in our explored cases., The end point of the $\nu$ p-process appears to be at $A \sim 180$ $^{180}$ Ta) in our explored cases.
+ It should be noted that the wind Ta)termination also plays a crucial role as explored in 4.1., It should be noted that the wind termination also plays a crucial role as explored in 4.1.
+" This is evident if we compare Figures 6 and 7, where the latter is the result for rw_= oo."," This is evident if we compare Figures 6 and 7, where the latter is the result for $r_\mathrm{wt} = \infty$ ."
+" Without wind termination, more proton-richness (AY.3~ 0.05) is requiredfor a given p- to be produced, but with asubstantially smaller production factor."," Without wind termination, more proton-richness $\Delta Y_\mathrm{e, 3} \sim 0.05$ ) is requiredfor a given p-nucleus to be produced, but with asubstantially smaller production factor."
+" The p-nuclei heavier than A=140 cannot be produced at all without wind termination (bottom panel in Figure We can understand the 7).reason for the above resultfrom Figure 8, which displays the snapshots of nucleosynthesis for selectedcases on the nuclear chart when the temperature drops to Ty=2 (left) and 1 "," The p-nuclei heavier than $A = 140$ cannot be produced at all without wind termination (bottom panel in Figure 7).We can understand the reason for the above resultfrom Figure 8, which displays the snapshots of nucleosynthesis for selectedcases on the nuclear chart when the temperature drops to $T_9 = 2$ (left) and 1 (right)."
+"Top, middle, and bottom panels are for the standard (right).model, that with Y.9 replaced by 0.800 (Το= 0.655), and that with Y,ο and ryt replaced by 0.800 and oo (without wind"," Top, middle, and bottom panels are for the standard model, that with $Y_\mathrm{e, 9}$ replaced by 0.800 $Y_\mathrm{e, 3} = 0.655$ ), and that with $Y_\mathrm{e, 9}$ and $r_\mathrm{wt}$ replaced by 0.800 and $\infty$ (without wind"
+The discovery of the accelerating universe. driven by some nivsterious dark energy. has motivated the conceptualization and design of a number of future surveys that will seek to discover its nature.,"The discovery of the accelerating universe, driven by some mysterious dark energy, has motivated the conceptualization and design of a number of future surveys that will seek to discover its nature."
+ These include. but are not limited to: Wide-field Eiber-fed Multi-Object Spectrograph (WEMOS). the Dark [ποιον Survey (DES). Panoramic Survey Telescope. Μαρια Response System (Pan-STARRS). Barvon Oscillation Spectroscopic Survey (BOSS). Large Sky Arca Alulti-Object Fibre Spectroscopic Telescope (LAAIOS7. Hubble. Sphere. Hydrogen: Survey (LISLIS). Square Wilometre Array (SIVA). Large. Synoptic Survey ‘Telescope (CLSS'T).. Euclid and the Joint Dark Energy Mission CJDISMD.," These include, but are not limited to: Wide-field Fiber-fed Multi-Object Spectrograph (WFMOS), the Dark Energy Survey (DES), Panoramic Survey Telescope Rapid Response System (Pan-STARRS), Baryon Oscillation Spectroscopic Survey (BOSS), Large Sky Area Multi-Object Fibre Spectroscopic Telescope (LAMOST), Hubble Sphere Hydrogen Survey (HSHS), Square Kilometre Array (SKA), Large Synoptic Survey Telescope (LSST), Euclid and the Joint Dark Energy Mission (JDEM)."
+ These will deploy an array of methods to obe the dark energy. such as baryon acoustic oscillations (BAO). weak lensing. cluster number counts. ancl (vpe-La supernovae (SIN-Ia).," These will deploy an array of methods to probe the dark energy, such as baryon acoustic oscillations (BAO), weak lensing, cluster number counts, and type-Ia supernovae (SN-Ia)."
+ In such a crowded marketplace it is important to ve a compelling product by demonstrating cllective use of resources., In such a crowded marketplace it is important to have a compelling product by demonstrating effective use of resources.
+ In previous papers some of the present authors have examined the application of design principles o the construction of new SULVONVS. bv optimizing the surveys to give the best science return. (Bassett 2005:," In previous papers some of the present authors have examined the application of design principles to the construction of new surveys, by optimizing the surveys to give the best science return (Bassett 2005;"
+Ecuation (11)) indicaes that the magnetic flux enclosed within C changes as a result of (1) advection of the mean maejetic field B by the mean flow @ (first term on the r.los.).,"Equation \ref{eq:mean_induction}) ) indicates that the magnetic flux enclosed within $\mathcal{C}$ changes as a result of (1) advection of the mean magnetic field $\overline{\bf B}$ by the mean flow $\overline{v}$ (first term on the r.h.s.),"
+ aud (9) correlations beweentheir fIuctuating components., and (2) correlations betweentheir fluctuating components.
+ Figure 10. slows plots of the mean magnetic field aud of the fluxtransport rates defited iu Eq. (11))., Figure \ref{fig:flux_transport} shows plots of the mean magnetic field and of the fluxtransport rates defined in Eq. \ref{eq:mean_induction}) ).
+ To suxotl out. [Iuctuatior sin time. the azimuthal means aud azimuthally-Iuctuaing quantities were aso averaged. between /9.3 aud /LO hr iu time.," To smooth out fluctuations in time, the azimuthal means and azimuthally-fluctuating quantities were also averaged between $t=9.3$ and $t=10$ hr in time."
+ This remainS ¢‘Ousisteut with the derivatio lla)ove since the oyeratious of averaging temporally and. azimuthalN COLLlle., This remains consistent with the derivation above since the operations of averaging temporally and azimuthally commute.
+ At this stage of he simulation. the mean iuagnetie field withi1 a radius of δ Min lias al'eacly reached kG values alid we focus our atention ou how the weaker ielcl at larger raclii is laisportecd.," At this stage of the simulation, the mean magnetic field within a radius of $8$ Mm has already reached kG values and we focus our attention on how the weaker field at larger radii is transported."
+ Tuspection of the mean velocity field 1'eveals that the contribution [rom the racially directed outflow (σι> 0) dominates such that maguetic {lus is trausporteck {om the spot., Inspection of the mean velocity field reveals that the contribution from the radially directed outflow $\overline{v}_r>0$ ) dominates such that magnetic flux is transported from the spot.
+ Howevend is found that correatious between the velocity aud magnetic fields (i.e. £) provide a meaus for tlje inward ulieration of verti‘al field.," However, it is found that correlations between the velocity and magnetic fields (i.e. $\vec{\bf \mathcal{E}}$ ) provide a means for the inward migration of vertical field."
+" The ampliπες of d; is in fact larger tlan that of $4, and the jet. result Is all accuuulation of flux in the spot.", The amplitude of $\dot{\Phi}_{\rm f}$ is in fact larger than that of $\dot{\Phi}_{\rm m}$ and the net result is an accumulation of flux in the spot.
+Finally. while it is impossible for a constant-time-stepping integrator to be moment sviplectic and (o exactly conserve energy (2)... variational integrators generally have bounded energy error.,"Finally, while it is impossible for a constant-time-stepping integrator to be momentum-conserving, symplectic and to exactly conserve energy \citep{Ge1988}, variational integrators generally have bounded energy error."
+ ? explain that the discretized. trajectory is sampling the continuous trajectory of a Lagrangian svstem. L. which is near L.," \citet{Lew2004} explain that the discretized trajectory is sampling the continuous trajectory of a Lagrangian system, $\tilde{L}$ which is near $L$."
+ L satisfies where the integral is evaluated on the trajectory which satisfies the Euler-Lagrange equations Lor L with εν)=qu and q(fq-h)=qo ancl the positions which are arguments for Lf satisfy equations(2).," $\tilde{L}$ satisfies where the integral is evaluated on the trajectory which satisfies the Euler-Lagrange equations for $\tilde{L}$ with $q\left( t_1
+\right) = q_1$ and $q\left( t_1 + h \right) = q_2$ and the positions which are arguments for $H$ satisfy equations."
+. Equation implies that {1 is the exact generating function for time evolution under L., Equation implies that $H$ is the exact generating function for time evolution under $\tilde{L}$.
+ In general it is only possible to compute a truncation of L to any desired order in fr. but il is possible to prove that L is close to L in the space of possible Lagrangians.," In general it is only possible to compute a truncation of $\tilde{L}$ to any desired order in $h$, but it is possible to prove that $\tilde{L}$ is close to $L$ in the space of possible Lagrangians."
+ Since the trajectory remains on the energy level set of L in phase space. near the energy level-set for L. the energy error remains bounded.," Since the trajectory remains on the energy level set of $\tilde{L}$ in phase space, near the energy level-set for $L$, the energy error remains bounded."
+ To define a variational integrator. we need an f/-(vpe function which approximates the action for a svstem over an interval.," To define a variational integrator, we need an $H$ -type function which approximates the action for a system over an interval."
+ There are many wavs to find such a finction: see ?. [or an extensive discussion of the various (vpes of integrator., There are many ways to find such a function; see \citet{Marsden2001} for an extensive discussion of the various types of integrator.
+ In this paper. we will focus on the so-called Galerkin Gauss-Lobatto (GGL) integrators.," In this paper, we will focus on the so-called Galerkin Gauss-Lobatto (GGL) integrators."
+ These integrators assume a polvnonmial trajectory in Gime and approximate the action integral using a Qauss-Lobatto ceqadrature rule., These integrators assume a polynomial trajectory in time and approximate the action integral using a Gauss-Lobatto quadrature rule.
+ Gauss-Lobatto quadrature is appropriate because it gives the highest-order integration rule lor a given number of points subject to the constraint that (he Lagrangian is evaluated once at the beginning and once at the end of the interval., Gauss-Lobatto quadrature is appropriate because it gives the highest-order integration rule for a given number of points subject to the constraint that the Lagrangian is evaluated once at the beginning and once at the end of the interval.
+" Evaluating at the beginning and end of the interval is important because il preserves (he svinetries of the continuous Lagrangian,", Evaluating at the beginning and end of the interval is important because it preserves the symmetries of the continuous Lagrangian.
+ ? show that all GGL integrators5 can be written as svinplectic partitioned integrators and derive formulas which relate the SPRIX coefficients to the discrete Lagrangian., \citet{Marsden2001} show that all GGL integrators can be written as symplectic partitioned Runge-Kutta integrators and derive formulas which relate the SPRK coefficients to the discrete Lagrangian.
+5 Because variational integrators5 are svmplectic. GGL integrators5 automatically salisly (he constraints on svmplectic Iunge-Ixutta coefficients (see. for example. ???)).," Because variational integrators are symplectic, GGL integrators automatically satisfy the constraints on symplectic Runge-Kutta coefficients (see, for example,\citet{Suris1989,Marsden2001,Stuchi2002}) )."
+ The particular integration aleorithims we consider here belong to a sub-class of partitioned methods often called Runge-hutta-Nvstronun methods., The particular integration algorithms we consider here belong to a sub-class of partitioned Runge-Kutta methods often called Runge-Kutta-Nyströmm methods.
+ We define, We define
+is really an implicit equation for the temperature since the mean opacities involved in the calculation of Ti and 75 depend onthe temperature 7.,is really an implicit equation for the temperature since the mean opacities involved in the calculation of $T_1$ and $T_2$ depend onthe temperature $T$ .
+ As in Calvetetal.(1991).. Calvetοἱal.(1992). and D05. we have assumed that the opacities are constant in the wall atmosphere. Le.. a; and ϱ are independent of τη.," As in \citet{Calvet2}, \citet{Calvet1} and D05, we have assumed that the opacities are constant in the wall atmosphere, i.e., $\alpha_j$ and $q_j$ are independent of $\tau_d$."
+" If the opacily is dominated by dust. and the temperature distribution in the wall atmosphere is such that there is no sublimation of any dust ingredient at a particular optical depth. then ib is valid to assume that a, is constant."," If the opacity is dominated by dust, and the temperature distribution in the wall atmosphere is such that there is no sublimation of any dust ingredient at a particular optical depth, then it is valid to assume that $\alpha_j$ is constant."
+ Also. if the contrast in temperature between the uppermost and the deepest lavers is not very high. y; and q; can be assumed to be constant too.," Also, if the contrast in temperature between the uppermost and the deepest layers is not very high, $\chi_d$ and $q_j$ can be assumed to be constant too."
+ Thus. in order to quantify the mean opacities. we use (he temperature evaluated at Ty=0. Le. (he surface temperature γω.," Thus, in order to quantify the mean opacities, we use the temperature evaluated at $\tau_d=0$, i.e., the surface temperature $T_{wall}$."
+ This is lound solving the implicit equation (2)). for a given wall radius A4. a given configuration of the binary svstem. and the coordinates of each pixel in the wall. a pair (LX.Y).," This is found solving the implicit equation \ref{eq_t}) ), for a given wall radius $R_{cb}$, a given configuration of the binary system, and the coordinates of each pixel in the wall, a pair $(X,Y)$."
+ The monochromatic emergent intensity from each pixel of the wall in the direction of the observer is calculated by integrating the transler equation., The monochromatic emergent intensity from each pixel of the wall in the direction of the observer is calculated by integrating the transfer equation.
+ We consider (wo components: the thermal emission and the scattering of stellar radiation., We consider two components: the thermal emission and the scattering of stellar radiation.
+" Finally. the SED is calculated by multüplving the emereent intensity bv (he solid angle subtended by each (7visible"") pixel. as seen by (he observer."," Finally, the SED is calculated by multiplying the emergent intensity by the solid angle subtended by each (“visible”) pixel, as seen by the observer."
+ As described in 2.. the size of the hole (R4) is given bv the dvnamical interaction between the disk and the stars.," As described in \ref{sec-description}, the size of the hole $R_{cb}$ ) is given by the dynamical interaction between the disk and the stars."
+ However. the Full hvdrodsnamical treatment should include ingredients like the gas viscosity. that allows material from the CD disk to lose angular momentum and fall into the hole.," However, the full hydrodynamical treatment should include ingredients like the gas viscosity, that allows material from the CB disk to lose angular momentum and fall into the hole."
+ This results in a small quantity of material inside the CD disk hole. that can be treated as being optically thin.," This results in a small quantity of material inside the CB disk hole, that can be treated as being optically thin."
+" Avoiding the precise description of the cdenamical evolution of such particles. in this work we simplify the problem considering them distributed in a region located between a radius 2=£B,,;, and 2=A... uniformly distributed in the azimuthal angle in a evlindrical coordinate svstem. centered at {he center of mass of the binary svstem."," Avoiding the precise description of the dynamical evolution of such particles, in this work we simplify the problem considering them distributed in a region located between a radius $R=R_{min}$ and $R=R_{cb}$, uniformly distributed in the azimuthal angle in a cylindrical coordinate system, centered at the center of mass of the binary system."
+ Indeed. simulations of the material inside the hole (Artvmowiez&Lubow1996:Günther&IxXlev2002) indicate that the material is non axisvinmetric distributed €xl spans all radii less (han Jt...," Indeed, simulations of the material inside the hole \citep{Artymowicz2,Gunther1} indicate that the material is non axisymmetric distributed and spans all radii less than $R_{cb}$."
+ However. for a circular binary svstem with undetectable accretion (towards the stars. we can argue that most of the material is restricted to lie at racius larger (han (he outer Lagrangian point for the three body cireular restricted problem (Szebehely 1967).," However, for a circular binary system with undetectable accretion towards the stars, we can argue that most of the material is restricted to lie at radius larger than the outer Lagrangian point for the three body circular restricted problem \citep{Szebehely}. ."
+. The surface density is given by a power law (XoxA.p= 1) such that X increases closer (o the wall.as is expected (Artvimowiez&Lubow 1996)..," The surface density is given by a power law $\Sigma\propto R^p, p=1$ ) such that $\Sigma$ increases closer to the wall,as is expected \citep{Artymowicz2}. ."
+burst data in excess of 2 keV. However. the burst temperatures during our low- and medium-temperature spectra are similar to those observed during the decay and peak of the bursts from 00748-0676.,"burst data in excess of 2 keV. However, the burst temperatures during our low- and medium-temperature spectra are similar to those observed during the decay and peak of the bursts from 0748–676."
+ We might therefore expect to see absorption features in our spectra., We might therefore expect to see absorption features in our spectra.
+ Rotational broadening may explain. the absence of absorption features in bburst spectra as it will modify the line profile (Ozzel Psaltis 2003)., Rotational broadening may explain the absence of absorption features in burst spectra as it will modify the line profile (Özzel Psaltis 2003).
+ Chang et al. (, Chang et al. (
+2005) calculated that for an edge-on neutron star with a rotation rate >200 Hz. the absorption features would become undetectable.,"2005) calculated that for an edge-on neutron star with a rotation rate $> 200$ Hz, the absorption features would become undetectable."
+ Thompson et al. (, Thompson et al. (
+2005) reported a possible burst oscillation in aat ~611 Hz. which. if proven to be true. would make the absorption lines less likely to be detected.,"2005) reported a possible burst oscillation in at $\sim 611$ Hz, which, if proven to be true, would make the absorption lines less likely to be detected."
+ 00748-6706. however. is a very slow rotator with a rotation EXOrate of 44.7 Hz (Villarreal Strohmayer 2004).," 0748–676, however, is a very slow rotator with a rotation rate of 44.7 Hz (Villarreal Strohmayer 2004)."
+ More recently. Chang et al. (," More recently, Chang et al. ("
+2006) fitted the line profiles of 00748-0676 with a theoretical model including the effect of rotational broadening; they reported a gravitational redshift of z=0.345700 confidence).,2006) fitted the line profiles of 0748–676 with a theoretical model including the effect of rotational broadening; they reported a gravitational redshift of $z = 0.345^{+0.005}_{-0.008}$ confidence).
+ However. they showed that rapid rotation does not necessarily make spectral lines undetectable.," However, they showed that rapid rotation does not necessarily make spectral lines undetectable."
+ If a rapidly rotating neutron star 1s seen face-on. the spectral lines will be narrower.," If a rapidly rotating neutron star is seen face-on, the spectral lines will be narrower."
+ 00748-0676 is an edge-on source while hhas an inclination. angle <70° (Homer. Charles. O'Donoghue 1998).," 0748–676 is an edge-on source while has an inclination angle $< 70^{\circ}$ (Homer, Charles, O'Donoghue 1998)."
+ A theoretical study shows that narrow lines can be detected even at spin frequencies as high as 600 Hz for emission near the rotation axis (Bhattacharyya et al., A theoretical study shows that narrow lines can be detected even at spin frequencies as high as 600 Hz for emission near the rotation axis (Bhattacharyya et al.
+ 20006)., 2006).
+ It is therefore uncertain whether a possible high spin frequency of the neutron star could be the cause of the of spectral lines in ssince most acereting neutron stars for which the spin has been determined have a spin frequency <600 Hz (e.g.. van der Klis 2006).," It is therefore uncertain whether a possible high spin frequency of the neutron star could be the cause of the non-detection of spectral lines in since most accreting neutron stars for which the spin has been determined have a spin frequency $< 600$ Hz (e.g., van der Klis 2006)."
+ The chemical composition of the atmosphere of the neutron star is another factor that can influence the line profiles (Chang et al., The chemical composition of the atmosphere of the neutron star is another factor that can influence the line profiles (Chang et al.
+ 2005)., 2005).
+ A more detailed theoretical study of line profiles ina mixed Hydrogen/Helium environment ts required to examine the effect on the ionization balance (see Bildsten 2000)., A more detailed theoretical study of line profiles in a mixed Hydrogen/Helium environment is required to examine the effect on the ionization balance (see Bildsten 2000).
+ We also have to consider the possibility that the lines seen in EXO 0748-676 are not gravitational redshifted iron lines from the surface., We also have to consider the possibility that the lines seen in EXO 0748-676 are not gravitational redshifted iron lines from the surface.
+ 00748-6706 ts a high inclination system and shows large dips in the X-ray light curve due to an absorbing medium in the outer accretion disk. far from the inner part of the system.," 0748–676 is a high inclination system and shows large dips in the X-ray light curve due to an absorbing medium in the outer accretion disk, far from the inner part of the system."
+ Homan. Wijnands. van den Berg (2003) proposed that a large fraction of the bursts used by Cottam et al. (," Homan, Wijnands, van den Berg (2003) proposed that a large fraction of the bursts used by Cottam et al. ("
+2002) occurred during times of significant absorption at low energies (<2 keV) due to obscuring medium in the outer disk.,2002) occurred during times of significant absorption at low energies $<2$ keV) due to obscuring medium in the outer disk.
+ This might have introduced absorption features in the burst spectra which could be due to the material inside the obscuring matter and not due to material on the surface of the neutron star., This might have introduced absorption features in the burst spectra which could be due to the material inside the obscuring matter and not due to material on the surface of the neutron star.
+ However. 1t is difficult to associate the lines seen in the burst spectra of 00748-6706 with He-like or H-like lines from abundant elements unless large velocity shifts are invoked.," However, it is difficult to associate the lines seen in the burst spectra of 0748–676 with He-like or H-like lines from abundant elements unless large velocity shifts are invoked."
+ In contrast. if obscuration in the outer disk drives dipping behavior. small velocity shifts would be expected.," In contrast, if obscuration in the outer disk drives dipping behavior, small velocity shifts would be expected."
+ On balance. it is not clear that obscuration in the outer disk could produce the absorption lines observed.," On balance, it is not clear that obscuration in the outer disk could produce the absorption lines observed."
+ Additional observations of burst. sources. with can help to better understand the conditions in which lines in X-ray bursts are strongest., Additional observations of burst sources with can help to better understand the conditions in which lines in X-ray bursts are strongest.
+ Future X-ray missions with high collecting area and high spectral resolution such asConstellation-X and are ideally suited to detecting absorption lines from the surface of neutron stars., Future X-ray missions with high collecting area and high spectral resolution such as and are ideally suited to detecting absorption lines from the surface of neutron stars.
+ We ultimately look forward to definitive observations with these missions., We ultimately look forward to definitive observations with these missions.
+derived for av=1 powerlaw lens with small ¢ and a source in the origin by Ixassiola Ixovner (1993).,"derived for a $\nu=1$ power–law lens with small $\epsilon$ and a source in the origin by Kassiola Kovner \shortcite
+{Kassiola1993}."
+. In contrast to the non-singular mass models by Kent Faleo (1988) our lens models do not produce a fifth image in the centre due to the central singularity of the mass distribution., In contrast to the non-singular mass models by Kent Falco \shortcite {Kent1988} our lens models do not produce a fifth image in the centre due to the central singularity of the mass distribution.
+ The predicted. source positions are very similar to the ones by Kent Falco., The predicted source positions are very similar to the ones by Kent Falco.
+ The remaining parameter space of the bar mocdels is illustrated. in figure 5.., The remaining parameter space of the bar models is illustrated in figure \ref{kappa0b}.
+" Similar to figure 3.. the contours of the confidence regions of normally distributed: models around the minimum of 47 in the parameter space of A and ὁ are shown for the model with v=1.05. e=0.31 and À—2. The parameter space was scanned with a stepsize of 0.003 in s, and 0.025 aresee in 5."," Similar to figure \ref {enuDegeneracy}, the contours of the confidence regions of normally distributed models around the minimum of $\chi^2$ in the parameter space of $\kappa_{\rm
+c}$ and $b$ are shown for the model with $\nu=1.05$, $e=0.31$ and $\lambda=2$ The parameter space was scanned with a stepsize of $0.003$ in $\kappa_{\rm c}$ and $0.025$ arcsec in $b$."
+ The bar parameters are well constrained at the bottom ofa steep boomerang-shaped valley., The bar parameters are well constrained at the bottom of a steep boomerang-shaped valley.
+ The surface mass density 5 in units of the critical density andthe local shear 7 at the positions of the images are given in table 2.., The surface mass density $\kappa$ in units of the critical density andthe local shear $\gamma$ at the positions of the images are given in table \ref {LocalLensingParameters}.
+ The values for ν=0.95 and 1.05 are similar to what one gets for a circular 7=1 powerlaw mass distribution (singular isothermal sphere) with a mixture of internal ancl external shear1994)., The values for $\nu=0.95$ and $1.05$ are similar to what one gets for a circular $\nu=1$ power–law mass distribution (singular isothermal sphere) with a mixture of internal and external shear.
+. It can be taken from tables 1 and 2. that a change of A does not cause a measurable change of observable quantities except for the semii-minor axis of the bar: the uncertainty of the model. parameters is dominated by the uncertainty of the bulge cllipticity e/exponent. v.," It can be taken from tables \ref
+{ParametersofModels} and \ref {LocalLensingParameters} that a change of $\lambda$ does not cause a measurable change of observable quantities except for the semi-minor axis of the bar; the uncertainty of the model parameters is dominated by the uncertainty of the bulge ellipticity $e$ /exponent $\nu$ ."
+ The width of the bar has not been measured previously., The width of the bar has not been measured previously.
+ In, In
+rreeious on timescales as short as 100 yr.,"regions on timescales as short as $\sim
+100\,$ yr."
+ As long as the Hickeriug contiuues. there is no direct relation between he age of the star and the size of the iregion.," As long as the flickering continues, there is no direct relation between the age of the star and the size of the region."
+ This result appears able to resolve the lifetime problem., This result appears able to resolve the lifetime problem.
+ Furthermore. we have ideutified the structures inside he rregious that produce the continuum enmüssiou.," Furthermore, we have identified the structures inside the regions that produce the continuum emission."
+ We fine hat enüssionu peaks are not necesarilv related to the xositious ofstars., We find that emission peaks are not necesarily related to the positions of stars.
+ Instead. stroug cission is produced by uwtiallv ionized eas from the accretion flow that cuters he ireegion.," Instead, strong emission is produced by partially ionized gas from the accretion flow that enters the region."
+ These filaments can be ionized aud blown away w thermal pressure., These filaments can be ionized and blown away by thermal pressure.
+ Since the continuum enmüsson Is determined solely bv the structure of the flow field around the massive star. the appearance of the rregion depends strongly on the augle from which it is viewed.," Since the continuum emission is determined solely by the structure of the flow field around the massive star, the appearance of the region depends strongly on the angle from which it is viewed."
+ For example. we fiud +that the same rregion can be classified as shell-like or cometary. depending on the position of the observer.," For example, we find that the same region can be classified as shell-like or cometary, depending on the position of the observer."
+ We have evaluated the distribution of apparent rregion morphologies., We have evaluated the distribution of apparent region morphologies.
+ The multiple sink simmlation run D reproduces the observed relative frequencies reported roni surveys. includiue the high fraction of spherical or unresolved regions.," The multiple sink simulation run B reproduces the observed relative frequencies reported from surveys, including the high fraction of spherical or unresolved regions."
+ The single sink simulation run A. rowever. fails to reproduce the morphology statistics.," The single sink simulation run A, however, fails to reproduce the morphology statistics."
+ This is because the sinele star grows so quickly that it is statistically unable to produce a significant fraction of he simallest rregious., This is because the single star grows so quickly that it is statistically unable to produce a significant fraction of the smallest regions.
+ The formation of a whole stellar cluster is recessary to ect the relative frequencies of the smallest rregious right., The formation of a whole stellar cluster is necessary to get the relative frequencies of the smallest regions right.
+ This analysis provides strong evidence against models of massive star formation where all or most hielhi-iuass stars fori in isolation., This analysis provides strong evidence against models of massive star formation where all or most high-mass stars form in isolation.
+ Iu addition to the ireeion morplologics. we can also reproduce the differeut SEDs characteristic of rregious.," In addition to the region morphologies, we can also reproduce the different SEDs characteristic of regions."
+ We find that our initial gas mass of 1000AL. is too small to produce observable dust emission at VLA or ALMA wavelengths.," We find that our initial gas mass of $1000\,M_\odot$ is too small to produce observable dust emission at VLA or ALMA wavelengths."
+ Nevertheless. ultracompact ireeious at different times aud places im our models show SEDs with both regular transitions from optically thick. to optically thin spectral slopes as well as anomalous scaling with a spectral slope around unity.," Nevertheless, ultracompact regions at different times and places in our models show SEDs with both regular transitions from optically thick to optically thin spectral slopes as well as anomalous scaling with a spectral slope around unity."
+ These slopes are entirely produced by inhomogencitics in the deusityv structure. with no contribution from dust.," These slopes are entirely produced by inhomogeneities in the density structure, with no contribution from dust."
+ Towever. follow-up simulations with more massive initial cluuips will be required to completely pin down the role of dust in observed SEDs.," However, follow-up simulations with more massive initial clumps will be required to completely pin down the role of dust in observed SEDs."
+ We thank Roberto Calvaun-Macdrid for helpful cobunents aud stinmlating discussions., We thank Roberto Galvánn-Madrid for helpful comments and stimulating discussions.
+ T.P. is a Fellow of theWürtteinberg funded by. their program Iuteruational Collaboration II (grant P-LS-SPIL/13)., T.P. is a Fellow of the funded by their program International Collaboration II (grant P-LS-SPII/18).
+ Πο also acknowledges support from an Annette Kade Fellowship for his visit to the American Museu of Natural History and a Visiting Scicutist Award of the Sinithsonian Astroplivsical Observatory (SAO)., He also acknowledges support from an Annette Kade Fellowship for his visit to the American Museum of Natural History and a Visiting Scientist Award of the Smithsonian Astrophysical Observatory (SAO).
+ We thank theForschungsgemceinschaft (DEC) for support via the Enuny Nocther Crauts BA 3607/1 aud KL1358/1l. as well as erants KL1358/lL. IL1358/5. KL 1358/10. aud IKL 1358/11. and the U.S. National Science," We thank the (DFG) for support via the Emmy Noether Grants BA 3607/1 and KL1358/1, as well as grants KL1358/4, KL1358/5, KL 1358/10, and KL 1358/11, and the U.S. National Science"
+gradients.,gradients.
+" This not only increases the mean value and the variations in |VP|, but also creates very peaked distributions and distributions with tails skewed towards the right."," This not only increases the mean value and the variations in $|\nabla \textbf{P}|$, but also creates very peaked distributions and distributions with tails skewed towards the right."
+" In terms of the Alfvénnic Mach number, the main contributor is the LOS direction chosen relative to the ordered field."," In terms of the Alfvénnic Mach number, the main contributor is the LOS direction chosen relative to the ordered field."
+" When we look parallel to the mean field as shown in in Figure 7,, the sub-Alfvénnic cases show higher mean and variance while the super-Alfvénnic cases show generally higher skewness and kurtosis."," When we look parallel to the mean field as shown in in Figure \ref{fig:moments}, the sub-Alfvénnic cases show higher mean and variance while the super-Alfvénnic cases show generally higher skewness and kurtosis."
+" We plot the mean and variance of the in Figure 8,, which confirms the trend seen in the gradients."," We plot the mean and variance of the in Figure \ref{fig:moments_rm}, which confirms the trend seen in the gradients."
+" In the case of the LOS parallel to the mean field, the mean value of the polarization angle has a strong Alfvénnic dependency with no sonic Mach number dependency."," In the case of the LOS parallel to the mean field, the mean value of the polarization angle has a strong Alfvénnic dependency with no sonic Mach number dependency."
+ The variance shows strong dependency on both sonic and Alfvénnic Mach numbers., The variance shows strong dependency on both sonic and Alfvénnic Mach numbers.
+" However, these trends are unique for sight-lines parallel to the mean field LOS."," However, these trends are unique for sight-lines parallel to the mean field LOS."
+ When looking perpendicular to the LOS the effects change., When looking perpendicular to the LOS the effects change.
+" For example, the moments of |VP| show stronger Skewness and kurtosis in the case of the sub-Alfvénnic models."," For example, the moments of $|\nabla \textbf{P}|$ show stronger skewness and kurtosis in the case of the sub-Alfvénnic models."
+" Thus to gauge the Alfvénnic Mach number using the distribution, the ordered LOS magnetic field is required."," Thus to gauge the Alfvénnic Mach number using the distribution, the ordered LOS magnetic field is required."
+ An additional effect that must be considered is the issue of the telescope resolution., An additional effect that must be considered is the issue of the telescope resolution.
+" For example, Figure 9 shows model number 6 without smoothing (left) and with Gaussian smoothing with FWHM-6 pixels (right)."," For example, Figure \ref{fig:smoothiimage} shows model number 6 without smoothing (left) and with Gaussian smoothing with FWHM=6 pixels (right)."
+ It is clear that smoothing changesthe distribution of maps of [VP|., It is clear that smoothing changesthe distribution of maps of $|\nabla \textbf{P}|$.
+ We plot the moments vs. smoothing FWHM for four models in Figure 10.., We plot the moments vs. smoothing FWHM for four models in Figure \ref{fig:momentsm}.
+" As the resolution of maps of Q and U decrease, so do the moments."," As the resolution of maps of Q and U decrease, so do the moments."
+ Thus one needs to take the smoothing of the data into account when comparing PDFs of |VP|., Thus one needs to take the smoothing of the data into account when comparing PDFs of $|\nabla \textbf{P}|$.
+ We investigate the moments of regions of the SGPS data with the most emission removing the bad pixels from our investigation., We investigate the moments of regions of the SGPS data with the most emission removing the bad pixels from our investigation.
+" We find that the SGPS data has average skewness of 0.825 and kurtosis of 0.928 and mean value and variance of 0.004 and 3.38110, respectively."," We find that the SGPS data has average skewness of 0.825 and kurtosis of 0.928 and mean value and variance of 0.004 and $3.38x10^{-6}$, respectively."
+" Comparing with our numerical set up, this most closely matches subsonic to transonic values of the moments, with telescope smoothing taken into account."," Comparing with our numerical set up, this most closely matches subsonic to transonic values of the moments, with telescope smoothing taken into account."
+" From section 4.1 it is clear that, for a given map of as the sonic Mach number increases the skewness and |VP|,kurtosis also increase."," From section 4.1 it is clear that, for a given map of $|\nabla \textbf{P}|$, as the sonic Mach number increases the skewness and kurtosis also increase."
+" However, these were globally averaged values of both the moments (i.e. the PDF of the entire image) and the sonic Mach number."," However, these were globally averaged values of both the moments (i.e. the PDF of the entire image) and the sonic Mach number."
+" If we want to look at smaller scale variations, we can calculate the PDFs of smaller portions of the image."," If we want to look at smaller scale variations, we can calculate the PDFs of smaller portions of the image."
+" Using a moving kernel (box, circle etc.),"," Using a moving kernel (box, circle etc.),"
+ we can create a “moment map” which is essentially a smoothed map that calculates the moments at every point with a given box size.," we can create a “moment map"" which is essentially a smoothed map that calculates the moments at every point with a given box size."
+" Because we are dealing with gradient quantities, we expect these maps to particularly trace regions where shocks are interacting and along shock fronts, thus tracing areas where the sonic Mach number is changing."," Because we are dealing with gradient quantities, we expect these maps to particularly trace regions where shocks are interacting and along shock fronts, thus tracing areas where the sonic Mach number is ."
+. We make skewness and kurtosis moment maps of the |VP| images similar to the method of Burkhart et al. (, We make skewness and kurtosis moment maps of the $|\nabla \textbf{P}|$ images similar to the method of Burkhart et al. (
+2010).,2010).
+" This results in a smoothed moment image, which can be compared to the actual image of the LOS local sonic Mach number (LOS average sonic Mach number for each pixel) with the same smoothing kernel as the |VP| moment map."," This results in a smoothed moment image, which can be compared to the actual image of the LOS local sonic Mach number (LOS average sonic Mach number for each pixel) with the same smoothing kernel as the $|\nabla \textbf{P}|$ moment map."
+" In this section, we investigate the relationship between the LOS sonic Mach number and the values of skewness and kurtosis on a pixel-by-pixel basis."," In this section, we investigate the relationship between the LOS sonic Mach number and the values of skewness and kurtosis on a pixel-by-pixel basis."
+" The key questions one must consider when making a moment map are what the kernel shape should be, what resolution is appropriate for good statistics, and how to compare the values of the higher order moments to the sonic Mach number?"," The key questions one must consider when making a moment map are what the kernel shape should be, what resolution is appropriate for good statistics, and how to compare the values of the higher order moments to the sonic Mach number?"
+ The last point is particularly important., The last point is particularly important.
+ While Burkhart et al. (, While Burkhart et al. (
+"2010) took a linear fit between the sonic Mach number and the moments in the case of column density, we will test how effective the moment maps can distinguish between supersonic and subsonic regimes on a pixel-by-pixel bases.","2010) took a linear fit between the sonic Mach number and the moments in the case of column density, we will test how effective the moment maps can distinguish between supersonic and subsonic regimes on a pixel-by-pixel bases."
+" Thus we stress that we are not trying to calculate the exact sonic Mach number with this method, which is a topic for another work."," Thus we stress that we are not trying to calculate the exact sonic Mach number with this method, which is a topic for another work."
+ We only seek to determine how well a moment map is able to pick out regimes of supersonic and subsonic., We only seek to determine how well a moment map is able to pick out regimes of supersonic and subsonic.
+" For our initial test, we choose to use a boxcar kernel for both the |VP| maps and for averaging the sonic Mach number, since a boxcar can handle edges in a square image most effectively."," For our initial test, we choose to use a boxcar kernel for both the $|\nabla \textbf{P}|$ maps and for averaging the sonic Mach number, since a boxcar can handle edges in a square image most effectively."
+" How well the moment map is able to determine if the gas is supersonic or subsonic depends on three parameters: the threshold level for skewness (yr, above which the gasis considered supersonic and below it is"," How well the moment map is able to determine if the gas is supersonic or subsonic depends on three parameters: the threshold level for skewness $\gamma_T$ , above which the gasis considered supersonic and below it is"
+produced very precise metallicity measurements.,produced very precise metallicity measurements.
+ We use these more recent values in our final calculations., We use these more recent values in our final calculations.
+ We put all of our cluster ages on the same relative scale., We put all of our cluster ages on the same relative scale.
+" Percivaletal.(2003) present the main sequence fitting technique applied in S09, and derive a metallicity dependency of A(B-V) for their procedure."," \citet{Percival03} present the main sequence fitting technique applied in S09, and derive a metallicity dependency of $\Delta (B-V)= 0.154\Delta[Fe/H]$ for their procedure."
+" Their algorithm produces distances relative to an assumed distance to the Hyades, principally by matching the shape of the Hyades MS to the cluster CMDs."," Their algorithm produces distances relative to an assumed distance to the Hyades, principally by matching the shape of the Hyades MS to the cluster CMDs."
+" Assuming a linear slope with a magnitude of —5 for the relevant portionof the Hyades MS (Percivaletal.2003).. .N(B-V)=MM, /5: therefore. AM,=0.77sA[Fe/H]."," Assuming a linear slope with a magnitude of $\sim5$ for the relevant portionof the Hyades MS \citep{Percival03}, $\Delta (B-V)= \Delta M_v/5$ ; therefore, $\Delta M_v = 0.77*\Delta[Fe/H]$."
+" performed an investigation of MS fitting techniques using nearby field stars with Hipparcos parallaxes and precise metallicities and find AM,=0.98«A[Fe/H].", \citet{Twarog09} performed an investigation of MS fitting techniques using nearby field stars with Hipparcos parallaxes and precise metallicities and find $\Delta M_v = 0.98*\Delta[Fe/H]$.
+ The metallicity dependency determined by Percivaletal.(2003) and Twarogetal.(2009) are similar; implementing the Twarogetal.(2009) relation would not significantly impact our results., The metallicity dependency determined by \citet{Percival03} and \citet{Twarog09} are similar; implementing the \citet{Twarog09} relation would not significantly impact our results.
+ We modify cluster ages corresponding to changes in distance according to the errors presented in S09., We modify cluster ages corresponding to changes in distance according to the errors presented in S09.
+" We assume that, for à given cluster. the fractional error in distance is equal to the fractional error in age."," We assume that, for a given cluster, the fractional error in distance is equal to the fractional error in age."
+" If G2—M) is the S09 cluster distance subtracted from the new distance then Gn—M)/0,,σι,=Mu. where σι.Μι and o, are from S09 and the new cluster age 18 fij,=fiu,09—Mins."," If $\Delta(m-M)$ is the S09 cluster distance subtracted from the new distance then $\Delta (m-M)/\sigma_{(m-M)}* \sigma_{t_{clus}}=\Delta t_{clus}$, where $\sigma_{(m-M)}$ and $\sigma_{\tclus}$ are from S09 and the new cluster age is $t_{clus}=t_{clus, S09} - \Delta t_{clus}$."
+ We adopt ages relative to those of S09 for consistency., We adopt ages relative to those of S09 for consistency.
+ Below we summarize our revisions to these six clusters., Below we summarize our revisions to these six clusters.
+ Απcetal.(2008) present a new distance measurement to this cluster using empirically calibrated isochrones., \citet{An08} present a new distance measurement to this cluster using empirically calibrated isochrones.
+" Additionally, they obtain high signal to noise spectroscopy of several Praesepe stars, reporting |Fe/H]20.112-0.03."," Additionally, they obtain high signal to noise spectroscopy of several Praesepe stars, reporting $=0.11\pm0.03$ ."
+ This is between the [Fe/H|=0.04+0.06 from the literature compilation of Gratton(2000) used by S09 and the subsequent higher Fe/H]=0.27+0.10 found by Paceetal.(2008)., This is between the $=0.04\pm0.06$ from the literature compilation of \citet{Gratton00} used by S09 and the subsequent higher $=0.27\pm0.10$ found by \citet{Pace08}.
+. Anetal.(2008) find (77-M)y=6.33£0.04 assuming [Fe/H]20.142-0.02 (a weighted mean of their result and literature values excluding non-members) and 0.002., \citet{An08} find $(m-M)_0=6.33\pm0.04$ assuming $=0.14\pm0.02$ (a weighted mean of their result and literature values excluding non-members) and $E(B-V)=0.006\pm0.002$ .
+ We adopt the latter Anetal.(2007) value for our metallicity., We adopt the latter \citet{An07} value for our metallicity.
+" Comparing the metallicity used in this distancedetermination with that of S09, we find A|Fe/H]=0.14—0.040.10."," Comparing the metallicity used in this distancedetermination with that of S09, we find $\Delta\feh=0.14-0.04=0.10$."
+ Accounting for this change in metallicity according to the metallicity dependence in Percivaletal.(2003).. the de-reddened S09 distance is Gi—M)y=6.32£0.04.," Accounting for this change in metallicity according to the metallicity dependence in \citet{Percival03}, the de-reddened S09 distance is $(m-M)_0=6.32 \pm 0.04$."
+" The Anetal.(2008). distance is 0.01mag larger than expected, implying a relatively younger cluster age (744,9) of 637£50Myr usingisochrones incorporating convective overshoot (OS) and 440+40Myr assuming no overshoot (nOS)."," The \citet{An08} distance is $0.01\ mag$ larger than expected, implying a relatively younger cluster age ) of $637\pm50\ Myr$ usingisochrones incorporating convective overshoot (OS) and $440\pm40\ Myr$ assuming no overshoot (nOS)."
+ The errors in the cluster age remain unchanged., The errors in the cluster age remain unchanged.
+" Anetal.(2008). simultancously best-fit their photometry of NGC 2516 with /H]=-0.0440.05, E(B—V)=0.117x0.002, and (n—M),=8.03 £0.04."," \citet{An08} simultaneously best-fit their photometry of NGC 2516 with $=-0.04\pm0.05$, $E(B-V)=0.117\pm0.002$, and $(m-M)_0=8.03\pm0.04$ ."
+ This represents a substantial improvement in σιη over S09 (σμωη)= 0.11)., This represents a substantial improvement in $\sigma_{\feh}$ over S09 $\sigma_{\feh}=0.11$ ).
+" De-reddening and accounting for A[Fe/H]= 0.12, the S09 distance is GnM),= 8.04."," De-reddening and accounting for $\Delta\feh=0.12$ , the S09 distance is $_{0}=8.04$ ."
+ Hence. AG-M)y=—0.01mag yielding lifetimes of 795=1373:29Myr andzyos=91£26 Myr.," Hence, $\Delta\dm_0=-0.01\ mag$ yielding lifetimes of $\tau_{OS}=137\pm29\ Myr$ and$\tau_{NOS}=91\pm26\ Myr$ ."
+" The error in cluster age has beenreduced from their S09 values by the factor 6,arvios/Tun-an.soo= 0.57."," The error in cluster age has beenreduced from their S09 values by the factor $\sigma_{\dm,An08}/\sigma_{\dm, S09}=0.57$ ."
+the host ealaxy or [rom intervening systems.,the host galaxy or from intervening systems.
+ A systematic discussion of these issues is presented in Totanietal.(2006)., A systematic discussion of these issues is presented in \citet{totani06}.
+. To rule out an ionised IGM. with a GRB afterglow. a multi-parameter fit to its spectrum that assumes an ionised IM and that accounts for the aforementioned uncertainties must provide a poor fit.," To rule out an ionised IGM with a GRB afterglow, a multi-parameter fit to its spectrum that assumes an ionised IGM and that accounts for the aforementioned uncertainties must provide a poor fit."
+ A tov example of such a multi-parameter fit that accounts or some of these uncertainties isshown in Figure 6.., A toy example of such a multi-parameter fit that accounts for some of these uncertainties isshown in Figure \ref{fig:fits}.
+ In his figure. the simple case of a bubble with size 7 in a iomogeneous ICM with neutral fraction wy is shown.," In this figure, the simple case of a bubble with size $R_b$ in a homogeneous IGM with neutral fraction $\bar{x}_H$ is shown."
+" The solid. curves are absorption models that have the intrinsic »wanmeters S3=r 1.25. ty,τον=633 sg=05 and Ng,=107"" em (thick curves). Ng;=107"" em (medium-width curves). or ανμι=107 ? (thin We then fit to these model curves."," The solid curves are absorption models that have the intrinsic parameters $\beta = 1.25$ , $z_g= z_{\rm DLA} = 6.3$ $\bar{x}_H = 0.5$ and $N_{\rm HI} =
+ 10^{19}$ $^{-2}$ (thick curves), $N_{\rm HI} = 10^{20}$ $^{-2}$ (medium-width curves), or $N_{\rm HI} = 10^{21}$ $^{-2}$ (thin We then fit to these model curves."
+ Our fitting weights all frequency bins equally., Our fitting weights all frequency bins equally.
+ “Phe presence of night-skv. OLI ines and the wavelength dependence of the CCD sensitivity should make the weighting slishtlv. non-uniform in a real observation., The presence of night-sky OH lines and the wavelength dependence of the CCD sensitivity should make the weighting slightly non-uniform in a real observation.
+ The dashed. curves in Figure 6 are fits over AA= to a model that only include DLA absorption.," The dashed curves in Figure \ref{fig:fits} are fits over $\Delta
+ \lambda = 200\; (1+z_g)$ to a model that only include DLA absorption."
+ This model fits for the parameters Nyy. 2. and ot.," This model fits for the parameters $N_{\rm HI}$, $\beta$ , and $A$ ."
+" The absorption owing to a substantially neutral LOAD is more easily [it with a DLA model as 22), increases or as Nut increases."," The absorption owing to a substantially neutral IGM is more easily fit with a DLA model as $R_b$ increases or as $N_{\rm
+ HI}$ increases."
+ Lowe had fit over a larger AA. the dashed ancl solid. curves would diller by less at wavelengths where the absorption is important.," If we had fit over a larger $\Delta \lambda$, the dashed and solid curves would differ by less at wavelengths where the absorption is important."
+ The fits in Figure ο assume that τιν is known from metal lines. as was the case for CGItD050904.," The fits in Figure \ref{fig:fits} assume that $z_{\rm DLA}$ is known from metal lines, as was the case for GRB050904."
+ For GRBs in more metal-poor galaxies or for GRBs with smaller DLAs. this may not always be possible.," For GRBs in more metal-poor galaxies or for GRBs with smaller DLAs, this may not always be possible."
+ When fits include σον as a free parameter. a neutral IGM is much harder to distinguish from an ionised one (Barkana&Loeb2004)..," When fits include $z_{\rm DLA}$ as a free parameter, a neutral IGM is much harder to distinguish from an ionised one \citep{barkana04b}."
+" To quantify the capability of distinguishing cdilferent absorption models. we define themeasure where £4, and £ are the afterglow fluxes for models 1and 2."," To quantify the capability of distinguishing different absorption models, we define themeasure where $F_1$ and $F_2$ are the afterglow fluxes for models $1$and $2$."
+ This quantity represents the average difference in A7 between two mocels for a spectrum that has noise 0(À;) in channel A; and that has No wavelength. channels. assuming the data is drawn from model 1 or 2.," This quantity represents the average difference in $\chi^2$ between two models for a spectrum that has noise $\sigma(\lambda_i)$ in channel $\lambda_i$ and that has $N$ wavelength channels, assuming the data is drawn from model $1$ or $2$ ."
+ When comparing a model with a neutral IM. to an lonised one. two models can be distinguished at X confidence level (C.L.) i£ No of the total likelihood. is contained between the maximum likelihood in the neutral IGAL model and. exp.GA\7)/2] of that likelihood value (assuming Ciaussianitv).," When comparing a model with a neutral IGM to an ionised one, two models can be distinguished at $X$ confidence level (C.L.) if $X$ of the total likelihood is contained between the maximum likelihood in the neutral IGM model and $\exp[-\langle \Delta \chi^2 \rangle/2]$ of that likelihood value (assuming Gaussianity)."
+ For a 4 parameter fit. two models that οπου hy {A\73=8S signifies tf—at model 1 is. in the mean. preferred at 914 CL. over model 2 (N47)=5. τικ C.L. and (N47)=12. 98% C.L.).," For a $4$ parameter fit, two models that differ by $\langle \Delta \chi^2 \rangle = 8$ signifies that model 1 is, in the mean, preferred at $91\%$ C.L. over model 2 $\langle \Delta \chi^2 \rangle = 5$, $70\%$ C.L. and $\langle \Delta \chi^2 \rangle = 12$, $98\%$ C.L.)."
+ Th' àenificance levels at fixed AY? for fits with three and five parameters are similar., The significance levels at fixed $\Delta \chi^2$ for fits with three and five parameters are similar.
+ In what follows. we quote GA\7} in terms of AA;=3A and στ= OL. where AA; is the width of a [frequency channel and o is the standard deviation in the [lux in each spectral channel in units of A.," In what follows, we quote $\langle \Delta
+ \chi^2 \rangle$ in terms of $\Delta \lambda_i = 3 \; \AA$ and $\sigma_{\cal F} = 0.1$ , where $\Delta \lambda_i$ is the width of a frequency channel and $\sigma_{\cal F}$ is the standard deviation in the flux in each spectral channel in units of $A$."
+ These are approximately the values of AA; and ez for the Subaru FOCAS spectroscopic observation of CU3050904., These are approximately the values of $\Delta \lambda_i$ and $\sigma_{\cal F}$ for the Subaru FOCAS spectroscopic observation of GRB050904.
+" First. we consider the case in which the absorption spectrum is parameterised by Vgm107"" ?. eg=0.5. and J,=10 Alpe (similar to the fits in Pie. 6)"," First, we consider the case in which the absorption spectrum is parameterised by $N_{\rm HI} \approx 10^{19}$ $^{-2}$, $\bar{x}_H =
+ 0.5$, and $R_b = 10$ Mpc (similar to the fits in Fig. \ref{fig:fits})"
+) We fit to AA—50(1|24)2b redward of source-frame Lvo.," $^8$ We fit to $\Delta \lambda = 50 \;(1+z_g) \;
+ \AA$ redward of source-frame $\alpha$."
+ This fit results in Ay?—5.(BfAA;)(O.lfor)? between the model for a neutral LOAL and the model for an ionised ΑΛ.," This fit results in $\langle \Delta \chi^2 \rangle = 5 \times (3 \AA / \Delta \lambda_i
+ )\times (0.1/\sigma_{\cal F})^2$ between the model for a neutral IGM and the model for an ionised IGM."
+ Next. we fit to a model with thesame specifications as the previous fit but with Napz1077 ? Num1074 em 7].," Next, we fit to a model with thesame specifications as the previous fit but with $N_{\rm HI} \approx
+ 10^{20}$ $^{-2}$ $N_{\rm HI} \approx 10^{21}$ $^{-2}$ ]."
+ This fit results in GN47)2302]«(BAYAN(OLfor}.," This fit results in $\langle \Delta \chi^2 \rangle = 3 \; [0.2] \times (3
+ \AA/\Delta \lambda_i) \times (0.1/\sigma_{\cal F})^2$."
+" Therefore. AA;Z3ch and or&0.1are required.© for anyhopeof detecting a neutral IGM for the Nyym107"" ty& 0.5. and f,=10 Mpe case.and even better sensitivity is required for the other two examples."," Therefore, $\Delta \lambda_i \lesssim 3 \; \AA$ and $\sigma_{\cal F} \lesssim
+ 0.1$are required for anyhopeof detecting a neutral IGM for the $N_{\rm HI} \approx 10^{19}$ $^{-2}$ , $x_H \approx 0.5$ , and $R_b =
+ 10$ Mpc case,and even better sensitivity is required for the other two examples."
+ Let us investigate how the quoted N47)depend on the assumptions we have made., Let us investigate how the quoted $\langle \Delta \chi^2 \rangle$depend on the assumptions we have made.
+" For this discussion. we again assume an absorption model with ry= 0.5. 2,=10 Alpe. and :Nyyzz107""0 em 7. and we assume a measurement with. AA;D. and ozz 0.1."," For this discussion, we again assume an absorption model with $\bar{x}_H = 0.5$ , $R_b = 10$ Mpc, and $N_{\rm
+HI} \approx 10^{20}$ $^{-2}$ , and we assume a measurement with $\Delta \lambda_i \approx 3 \; \AA$ and $\sigma_{\cal F} \approx 0.1$ ."
+As noted above. SZPYIXIO reported that the purported UV object in the 1995 frames coincides with a bright star. Star S. seen in images taken al longer wavelengths.,"As noted above, SZPYK10 reported that the purported UV object in the 1995 frames coincides with a bright star, Star S, seen in images taken at longer wavelengths."
+ Fig., Fig.
+ 3 in SZPYI0 gives a finding ehart for this star. and it is also visible (but not marked) in Fig.," 3 in SZPYK10 gives a finding chart for this star, and it is also visible (but not marked) in Fig."
+ 1 ol PaperL., 1 of Paper.
+ Although I have just presented evidence that the 1995 object is a spurious case of CRs striking twice. I carried out a eareful registration of the 1995 frames with an ACS image obtained in an F435W filler in 2003. using stars detected in both images.," Although I have just presented evidence that the 1995 object is a spurious case of CRs striking twice, I carried out a careful registration of the 1995 frames with an ACS image obtained in an F435W filter in 2003, using stars detected in both images."
+ The result shows that the spurious UV object is offset [rom the real optical Star S byOQ!17. approximately 1l times the lo error of the registration.," The result shows that the spurious UV object is offset from the real optical Star S by$0\farcs17$, approximately 11 times the $\sigma$ error of the registration."
+ Thus. even if the UV source had been real. it did nol coincide with the star seen al optical wavelengths.," Thus, even if the UV source had been real, it did not coincide with the star seen at optical wavelengths."
+ llowever. SZPYKIO have futher reported that the optical star is blue and extremely hot. and that it is therelore a strong candidate for the remnant of M31 RV as well as providing support for the classical-nova hypothesis.," However, SZPYK10 have further reported that the optical star is blue and extremely hot, and that it is therefore a strong candidate for the remnant of M31 RV as well as providing support for the classical-nova hypothesis."
+ This is again in striking contradiction to the findings we reported in PaperL., This is again in striking contradiction to the findings we reported in Paper.
+ In Figure 3 of the present paper. I have reproduced our color-magnitude diagram (CMD) from Paper I. which plots stars that lie within the 30 positional error box lor M31RV.," In Figure 3 of the present paper, I have reproduced our color-magnitude diagram (CMD) from Paper I, which plots stars that lie within the $\sigma$ positional error box for M31."
+. This CMD was derived by. M. Siegel from PSF photometry of an excellent set of dithered: WFEPC2/PC frames obtained in 1999. as described in detail in PaperL.," This CMD was derived by M. Siegel from PSF photometry of an excellent set of dithered WFPC2/PC frames obtained in 1999, as described in detail in Paper."
+ The magnitudes in the WEPC2 F555W and FSi4W fillers were (ransformecl {ο the traditional ground-based Johnson-Ixron-C'óusins V. and / svstem., The magnitudes in the WFPC2 F555W and F814W filters were transformed to the traditional ground-based Johnson-Kron-Cousins $V$ and $I$ system.
+ Star ο is the brightest object in the error box. and is now labelled in Figure 3.," Star S is the brightest object in the error box, and is now labelled in Figure 3."
+ We had measured its magnitudes and color as V.—22.95. J= 21.33. and V—4=1.62. with errors of about 40.01 0.02 mag.," We had measured its magnitudes and color as $V=22.95$, $I=21.33$ , and $V\!-\!I=1.62$, with errors of about $\pm$ 0.01–0.02 mag."
+ As we already remarked in Paper L this CMD shows only a population of cool red giants belonging to the M31 bulge.," As we already remarked in Paper I, this CMD shows only a population of cool red giants belonging to the M31 bulge."
+ To verily that there are only old red giants in the error box. I have also added three isochrones to Figure 3. taken from the DaSTI (Pietrinferni οἱ 22004. 2006).," To verify that there are only old red giants in the error box, I have also added three isochrones to Figure 3, taken from the BaSTI (Pietrinferni et 2004, 2006)."
+ I chose parameters (vpically found in recent studies of the M31 bulge population (e.g.. Saglia et 22010 and references therein): an age of 10 Gyr and metallicities of ΔΕΗ)=+0.25». —0.25. and —0.66. with sealed solar element ratios.," I chose parameters typically found in recent studies of the M31 bulge population (e.g., Saglia et 2010 and references therein): an age of 10 Gyr and metallicities of $\rm [M/H] = +0.25$, $-0.25$, and $-0.66$, with scaled solar element ratios."
+" E then adjusted the isochrones for a reddening of LE(DB—-V)=0.062 (Schlegel. Finkbeiner. Davis 1998). corresponding to £(V—I)=0.08. and for an M31 distance modulus of (n—M),=24.4 (van den Dergh 2000)."," I then adjusted the isochrones for a reddening of $E(B\!-\!V)=0.062$ (Schlegel, Finkbeiner, Davis 1998), corresponding to $E(V\!-\!I)=0.08$, and for an M31 distance modulus of $(m-M)_0=24.4$ (van den Bergh 2000)."
+ The three adjusted isochrones nicely bracket the positions of all of these normal red giants in theCAID., The three adjusted isochrones nicely bracket the positions of all of these normal red giants in the.
+". S(ar ο, with a Jolnson-INron-C'ousins color index of (V—7)y= 1.54. has an implied effeclive temperature near Tuy=4000 Ix(e.g.."," Star S, with a Johnson-Kron-Cousins color index of $(V\!-\!I)_0 = 1.54$ , has an implied effective temperature near $\Teff=4000$ K(e.g.,"
+ Dessell 1979)., Bessell 1979).
+ How then did SZPYIX10 find, How then did SZPYK10 find
+Procvon A (herealter referred to Proevon). one of the brightest stars in the sky. is a well studied object.,"Procyon A (hereafter referred to Procyon), one of the brightest stars in the sky, is a well studied object."
+ Known lor its white dwarf companion. its orbit has been well determined bv speckle technique. and its mass. long quite uncertain. has now been evaluated (o be near 1.5 M. (Girard et al.," Known for its white dwarf companion, its orbit has been well determined by speckle technique, and its mass, long quite uncertain, has now been evaluated to be near 1.5 $M_{\odot}$ (Girard et al."
+ 2000)., 2000).
+ On the basis of evolutionary interior models. Procvon's internal structure is believed to consist of a convective core. wilh a radiative envelope ancl a thin surface convection zone.," On the basis of evolutionary interior models, Procyon's internal structure is believed to consist of a convective core, with a radiative envelope and a thin surface convection zone."
+" The mass in ils outer convection zone is 10.""M. at present (Guenther Demarque 1993).", The mass in its outer convection zone is $10^{-5} M_{\odot}$ at present (Guenther Demarque 1993).
+ Procvon thus offers an unusual opportunity to study. by seismic means both the structure of the outer lavers and (he elficieney of convective core overshoot in intermediate mass stars (Chabover οἱ al., Procyon thus offers an unusual opportunity to study by seismic means both the structure of the outer layers and the efficiency of convective core overshoot in intermediate mass stars (Chaboyer et al.
+ 1999: Straka et al., 1999; Straka et al.
+ 2005)., 2005).
+ For all these reasons. and in view of theoretical predictions that the p-mode oscillation amplitudes would be relatively large (Chiristensen-Dalseaard Frandsen 1050: IHoudek et al.," For all these reasons, and in view of theoretical predictions that the $p$ -mode oscillation amplitudes would be relatively large (Christensen-Dalsgaard Frandsen 1983; Houdek et al."
+ 1999). Procvon has been a prime candidate for seismic observations from (he eround and in space.," 1999), Procyon has been a prime candidate for seismic observations from the ground and in space."
+ Interestingly. now that reliable observations are possible. (here is an apparent discrepancy between the detection of p-mode oscillations on Procyon. based on radial velocity measurements made from the ground (Martié et al.," Interestingly, now that reliable observations are possible, there is an apparent discrepancy between the detection of $p$ -mode oscillations on Procyon, based on radial velocity measurements made from the ground (Martić et al."
+ 2004 and Eggenberger et al., 2004 and Eggenberger et al.
+ 2004). and," 2004), and"
+Üem for any direction jr.,0em for any direction $\mu_{J}$.
+ These new values of the coefficients can overrun the memory places of (he pervios ones. corresponding to the intial condition given by eq. (," These new values of the coefficients can overrun the memory places of the pervios ones, corresponding to the intial condition given by eq. ("
+25).,25).
+ 3em We repeat (he same exercice. namely (to emplov eq. (," 3em We repeat the same exercice, namely to employ eq. ("
+15) inside both the functional form [or S(ru). given by eq. (,"15) inside both the functional form for $S(\tau_{0})$, given by eq. ("
+"29). and that for 5""(7,). given by eq. (","29), and that for $S^{\prime\prime}(\tau_{0})$, given by eq. ("
+26). in order to recover the previous form for both of them. but now as a function of the upeoing intensities al τι insted of 75. hence with different coelfficients.,"26), in order to recover the previous form for both of them, but now as a function of the upgoing intensities at $\tau_{1}$ insted of $\tau_{0}$, hence with different coeffficients."
+" That is Jem ancl Now. just by solving for τη) and 5""(7,) we obtain the coefficients of the relations (27) and (28). earlier announced at the beginning of Section 3.2."," That is 0em and Now, just by solving for $S(\tau_{0})$ and $S^{\prime\prime}(\tau_{0})$ we obtain the coefficients of the relations (27) and (28), earlier announced at the beginning of Section 3.2."
+ These coefficients must be stored for further use., These coefficients must be stored for further use.
+ Ii such a wav we have achieved part of oul goal., In such a way we have achieved part of out goal.
+ 3em At this point let us describe the propagation of the downgoing intensities from L(7).77). eiven bv eq. (," 3em At this point let us describe the propagation of the downgoing intensities from $I^{-}(\tau_{0},\mu_{J})$, given by eq. ("
+30). lo £(74.0j) according to eq. (,"30), to $I^{-}(\tau_{1},\mu_{J})$ according to eq. ("
+16).,16).
+ By re-arrangement of the coefficients we gel (he new values corresponding to the relation Oem for all the directions joy., By re-arrangement of the coefficients we get the new values corresponding to the relation 0em for all the directions $\mu_{J}$.
+ Again these new values of the coellicients can overrun the previous ones. corresponding to eq. (," Again these new values of the coefficients can overrun the previous ones, corresponding to eq. ("
+30).,30).
+To resolve these problems. NJWIO5 proposed a new model for the origin of Q-tvpe NEAs by postulating that the optically-active laver on (heir surface has been recently reset bv the For example. applied (dal stresses applied may cause elements of a fractured body to move wilh respect to each other. ballistically displace surface material. or even liberate the surface lavers [from the asteroid.,"To resolve these problems, NJWI05 proposed a new model for the origin of Q-type NEAs by postulating that the optically-active layer on their surface has been recently reset by the For example, applied tidal stresses applied may cause elements of a fractured body to move with respect to each other, ballistically displace surface material, or even liberate the surface layers from the asteroid."
+ Alternatively. if the tidal torque spins up an asteroid. the weathered regolith lavers can be removed by carrving away the excess angular momentum.," Alternatively, if the tidal torque spins up an asteroid, the weathered regolith layers can be removed by carrying away the excess angular momentum."
+ To show the plausibility of this idea. NJWIO5 estimated (hat a twpical NEA suffers on average about one encounter to within 2 Roche radii (I4) from the Earth every 225 My.," To show the plausibility of this idea, NJWI05 estimated that a typical NEA suffers on average about one encounter to within 2 Roche radii $R_{\rm Roche}$ ) from the Earth every $\approx$ 5 My."
+ This time interval between encounters is comparable with the average orbital lifetime of NEAs (5 Mv according to Bottke et al., This time interval between encounters is comparable with the average orbital lifetime of NEAs $\approx$ 5 My according to Bottke et al.
+ 2002) ancl is also comparable with the range of the OW timescale discussed above., 2002) and is also comparable with the range of the SW timescale discussed above.
+" Consecuentlv. if Gdal encounters at 27225,4444 Can reset surfaces. Q-ivpe NEAs should be numerous."," Consequently, if tidal encounters at $2 R_{\rm Roche}$ can reset surfaces, Q-type NEAs should be numerous."
+ For comparison. encounters up to live planetary radii (this limit depends on the NEAs shape and spin vector) can produce strong shape distortions of a rubble-pile NEA and material stripping up to of the NEAs pre-encounter mass (Richardson et al.," For comparison, encounters up to five planetary radii (this limit depends on the NEA's shape and spin vector) can produce strong shape distortions of a rubble-pile NEA and material stripping up to of the NEA's pre-encounter mass (Richardson et al."
+ 1998)., 1998).
+ Strong evidence supporting (he NJWIOS model comes [rom the observed orbital distribution of Q-tvpe NEAs., Strong evidence supporting the NJWI05 model comes from the observed orbital distribution of Q-type NEAs.
+ Several trends were pointed out (Nesvorny οἱ al., Several trends were pointed out (Nesvorný et al.
+ 2004. NJWIOS. Marchi et al.," 2004, NJWI05, Marchi et al."
+ 2006): (a) the proportion of Q-tvpe NEAs increases with decreasing perihelion distance gq. (b) the orbital distribution of known Q-tvpe NEAs has a sharp edge αἱ q1 AU: and (c) concentrations of Q-tvpe NEAs occur for values of q that correspond to large collision probability with the terrestrial planets (g¢=1.0 and 0.72 AU for Earth and Venus: see. e.g.. Alorbidelli ancl Glaciman 1998).," 2006): (a) the proportion of Q-type NEAs increases with decreasing perihelion distance $q$; (b) the orbital distribution of known Q-type NEAs has a sharp edge at $q\sim1$ AU; and (c) concentrations of Q-type NEAs occur for values of $q$ that correspond to large collision probability with the terrestrial planets $q=1.0$ and 0.72 AU for Earth and Venus; see, e.g., Morbidelli and Gladman 1998)."
+ This correlation of the orbital distribution of Q-tvpe NEAs with the collision probabilitw is expected in the NJWIO5 model because orbits with laree collision probability also have frequent close encounters wilh the terrestrial planets (e.g... Bottke Melosh 1996).," This correlation of the orbital distribution of Q-type NEAs with the collision probability is expected in the NJWI05 model because orbits with large collision probability also have frequent close encounters with the terrestrial planets (e.g., Bottke Melosh 1996)."
+ Thus. the surfaces of these objects are expected to be ‘fresher on average showing fewer signs of the SW effects.," Thus, the surfaces of these objects are expected to be `fresher' on average showing fewer signs of the SW effects."
+ Recently. Binzel et al. (," Recently, Binzel et al. ("
+2010: herealier B10) presented results supporting the NJWI105 model.,2010; hereafter B10) presented results supporting the NJWI05 model.
+ They used a taxonomic classification based on the near-inlrared (near-I8.) spectroscopy which should more closely characterize surface mineralogy than previous taxononmies based on the visible wavelengths., They used a taxonomic classification based on the near-infrared (near-IR) spectroscopy which should more closely characterize surface mineralogy than previous taxonomies based on the visible wavelengths.
+ We have verified. however. (hat there exists a very good correlation between the visible aud near-IR. classifieation of Q tvpes.," We have verified, however, that there exists a very good correlation between the visible and near-IR classification of Q types."
+ For example. from 6 NEAs that were Classified as Qs from the visible spectra before IR data became available. 5 were classified as Qs based on the near-IR. spectra and only 1: must have been re-classified as Sq.," For example, from 6 NEAs that were classified as Qs from the visible spectra before IR data became available, 5 were classified as Qs based on the near-IR spectra and only 1 must have been re-classified as Sq."
+ This shows that the use of (he near-IR data does not really change the problem., This shows that the use of the near-IR data does not really change the problem.
+ D10 numerically integrated the orbits of 95 selected asteroids (Q-. Sep. and S-tvpe NEAs and more distant. Mars-crossers) for 0.5 My. recorded the history of their close encounters to the Earth and statistically analvzed the encounters in an attempt (o correlate the statistics with spectroscopic (wpe.," B10 numerically integrated the orbits of 95 selected asteroids (Q-, Sq-, and S-type NEAs and more distant Mars-crossers) for 0.5 My, recorded the history of their close encounters to the Earth and statistically analyzed the encounters in an attempt to correlate the statistics with spectroscopic type."
+ Two main results were obtained in DIO: (1) 20 Q- and 55 S/Seq-tvpe NEAs can collide with the Earth in 0.5 My. (although the actual probability of such a collision is low)., Two main results were obtained in B10: (1) 20 Q- and 55 S/Sq-type NEAs can collide with the Earth in 0.5 My (although the actual probability of such a collision is low).
+ Conversely. none of the remaining 13 NEAs and T Mars-crossing asteroids (AICAs) included in the study. all 5 and Sq. have any chance of impact in 0.5 Mx.," Conversely, none of the remaining 13 NEAs and 7 Mars-crossing asteroids (MCAs) included in the study, all S and Sq, have any chance of impact in 0.5 My."
+ The time interval of 0.5 My used in the DIO study was motivated by the recent estimate of the SW Umescale in Vernazza et al. (, The time interval of 0.5 My used in the B10 study was motivated by the recent estimate of the SW timescale in Vernazza et al. (
+2009) where it was proposed. based on a comparative study. of asteroid families aud. OC meteorites. that SW acts on «1 Mx to fully aller an asteroid spectrum from Q to S. Adopting this timescale. BLO computed,"2009) where it was proposed, based on a comparative study of asteroid families and OC meteorites, that SW acts on $<$ 1 My to fully alter an asteroid spectrum from Q to S. Adopting this timescale, B10 computed"
+values are quite comparable to the most recent study of IxXarachentsevetal.(2006) who determined the orbital mass of the Centaurus group to be 7.5x107 M... the virial mass 8.9x107> AL... and the mass determined. within. a zero-velocity. surface. including. dark energy G441.8x1077 NL. Karachents,"values are quite comparable to the most recent study of \cite{kar06}
+ who determined the orbital mass of the Centaurus group to be $7.5
+\times 10^{12}$ $_{\sun}$, the virial mass $8.9 \times 10^{12}$ $_{\sun}$, and the mass determined within a zero-velocity surface including dark energy $6.4\pm1.8 \times 10^{12}$ $_{\sun}$."
+eyοἱal.(2006). also determined a M/Lj=137 M. 1..., \cite{kar06} also determined a $_B = 137$ $_{\sun}$ $_{\sun}$.
+ The M/Lj; ratio for the Centaurus group is larger (han many other groups of galaxies in the nearby universe Γον—58—88 M. /L. (Ixarachentsev.2005))). but mass-to-light ratios are (vpically larger for groups dominated by eiat ellipticals. compared to spiral dominated eroups. such as the Milkv Way (M/Lj=28—29 M. /L.). C342 (Al/Lp=18—30 M. /L.). or MSI (M/Lgy=19—32 M. /L.) Uxarachentsey2005).," The $_B$ ratio for the Centaurus group is larger than many other groups of galaxies in the nearby universe $_B = 8-88$ $_{\sun}$ $_{\sun}$ \citep{kar05}) ), but mass-to-light ratios are typically larger for groups dominated by giant ellipticals, compared to spiral dominated groups, such as the Milky Way $_B = 28-29$ $_{\sun}$ $_{\sun}$ ), IC342 $_B = 18-30$ $_{\sun}$ $_{\sun}$ ), or M81 $_B = 19-32$ $_{\sun}$ $_{\sun}$ ) \citep{kar05}."
+. In brief. the satellite galaxies in the Centaurus group strongly indica(e a possible connection to the halo svstem of ils central group member.," In brief, the satellite galaxies in the Centaurus group strongly indicate a possible connection to the halo system of its central group member."
+ Theconfirmed members in both the Centaurus and M82 groups. wilh their brightest menibers separated by a projected radius of 1 Alpe. have many surrounding galaxies with distances and/or radial velocities (hat do nol have confirmed eroup membership.," The members in both the Centaurus and M83 groups, with their brightest members separated by a projected radius of 1 Mpc, have many surrounding galaxies with distances and/or radial velocities that do not have confirmed group membership."
+ It seems a possibility that these (wo eroups resemble the Milky Way and AISI as bound systems., It seems a possibility that these two groups resemble the Milky Way and M31 as bound systems.
+ Exploring this idea leads to the application of a simple Newtonian bounding argument. used recently by Broughetal.(2006)... al. (2004).. ancl others.," Exploring this idea leads to the application of a simple Newtonian bounding argument, used recently by \cite{brough06}, \cite{cortese04}, and others."
+ IE the two groups. Centaurus and M82. are part of a larger svsten. then the inequality holds.," If the two groups, Centaurus and M83, are part of a larger system, then the inequality holds."
+" In Equation 8.. Mοgroup+Ce=(9.605dE3.01)x1015 AL. is the total mass of the confirmed system. ve,=25 km ! is the relative velocity between the (wo groups. R,=907 kpeis the projected distance between the centroids ofthe (wo groups. and 7 is the projection angle between the plane of (he skv and the line that joins the centers of the two eroups."," In Equation \ref{eqn:bounding}, $M = M83_{group} + Cen_{group} =
+(9.65\pm3.01) \times 10^{12}$ $_{\sun}$ is the total mass of the confirmed system, $v_{r} = 25$ km $^{-1}$ is the relative velocity between the two groups, $R_p = 907$ kpc is the projected distance between the centroids ofthe two groups, and $\beta$ is the projection angle between the plane of the sky and the line that joins the centers of the two groups."
+" The inequality in Equation 8 holds for all projected angles between 5°<9<89"". using (he observed radial velocities of the Centaurus group ancl M83 centroids (o0 be those of NGC 5128. 541 km !. and NGC 5236. 516 kins !. respectively."," The inequality in Equation \ref{eqn:bounding} holds for all projected angles between $5^o \leq \beta \leq 89^o$, using the observed radial velocities of the Centaurus group and M83 centroids to be those of NGC 5128, $541$ km $^{-1}$, and NGC 5236, $516$ km $^{-1}$, respectively."
+ The probability (had Centaurus and M83 are bound can then be caleulated from (Girardietal. 2005).., The probability that Centaurus and M83 are bound can then be calculated from \citep{girardi05}. .
+ I find that the probability that these svstenis are bound is at the 91% level., I find that the probability that these systems are bound is at the $91\%$ level.
+in equaion (23)). we recover the perturbation equations of except that here i6 depends on the scale factor a and introduces a new source term. ὁ(0)=ΠΟ. in the density perturbation equation.,"in equation \ref{eq:densitypert}) ), we recover the perturbation equations of except that here $w$ depends on the scale factor $a$ and introduces a new source term, $\mathcal{E}(\dot{w})=
+3\mathcal{H}\dot{w}v/k$, in the density perturbation equation."
+ Note that this term depends on the derivative of the equation of state. d.," Note that this term depends on the derivative of the equation of state, $\dot{w}$."
+ Equation (24)) does not have a new terni., Equation \ref{eq:velocitypert}) ) does not have a new term.
+ As discussed by a single classical scalar field cannot evolve from a quintessence-like. woo1. to phantom-like. TE. behavior.," As discussed by a single classical scalar field cannot evolve from a quintessence-like, $w>-1$, to phantom-like, $w<-1$, behavior."
+ However. proposed a single scalar field model where a super-aceclerated phase (uw< 1) of the cosmic expansion can be achieved via quantum olfects.," However, proposed a single scalar field model where a super-accelerated phase $w<-1$ ) of the cosmic expansion can be achieved via quantum effects."
+ Later. showed that this model is stable.," Later, showed that this model is stable."
+ Alternatively. suggested the so-called quintom model: in this moclel. the effective equation of state of two combined scalar fields. one with (m1 and the other with ο.1. can cross the cosmological-constant boundary. w=1. as il evolves in time.," Alternatively, suggested the so-called quintom model: in this model, the effective equation of state of two combined scalar fields, one with $w>-1$ and the other with $w<-1$, can cross the cosmological-constant boundary, $w=-1$, as it evolves in time."
+ Other models that allow wr) to cross this boundary have also been proposed., Other models that allow $w(a)$ to cross this boundary have also been proposed.
+ However. for practical purposes. using an elective. evolving dark energy. equation of state produces a well-known divergence in equation (24)) when wa)=1.," However, for practical purposes, using an effective, evolving dark energy equation of state produces a well-known divergence in equation \ref{eq:velocitypert}) ) when $w(a)=-1$."
+ This divergence can be avoided. bv imposing ὁ0 and =0 within the logarithmic singularity region. w=1|[e]. where ο is infinitesimally. small.," This divergence can be avoided by imposing $\dot{\delta}=0$ and $\dot{v}=0$ within the logarithmic singularity region, $w=-1+|\epsilon|$, where $\epsilon$ is infinitesimally small."
+ Inaccuracies in this approximation have a negligible impact on the resulting CMD power spectra(???)., Inaccuracies in this approximation have a negligible impact on the resulting CMB power spectra.
+. In what follows. we present results for cases where dark energv clustering is either accounted for or ignored in the CAIB analysis (Section 5 4).," In what follows, we present results for cases where dark energy clustering is either accounted for or ignored in the CMB analysis (Section \ref{deevo}) )."
+ Dark energy. clustering docs not have a significant impact on the fi. analysis., Dark energy clustering does not have a significant impact on the $f_{\rm gas}$ analysis.
+ As described in Section 2.. to enable a direct comparison with the predicted dark energy constraints for other planned experiments. we parameterize our results in terms of the DETTE FoAL We present. results for a fiducial fi; |CMD data set. incorporating the statistical uncertainties. and systematic allowances described above. and with zero scatter about the fiducial curves.," As described in Section \ref{sec:de}, to enable a direct comparison with the predicted dark energy constraints for other planned experiments, we parameterize our results in terms of the DETF FoM. We present results for a fiducial $f_{\rm gas}$ +CMB data set, incorporating the statistical uncertainties and systematic allowances described above, and with zero scatter about the fiducial curves."
+ Phe absence of scatter in the simulated: samples ensures that the peaks of the posterior probability cüstributions occur at the expected values. in the," The absence of scatter in the simulated samples ensures that the peaks of the posterior probability distributions occur at the expected values, in the"
+just large enough for gravitational collapse to overcome he IGAL thermal pressure. but small cnough to have ιο direct progenitors above a cosmiological Jeans mass.,"just large enough for gravitational collapse to overcome the IGM thermal pressure, but small enough to have no direct progenitors above a cosmological Jeans mass."
+ After rejonization. such pristine star formation would ο fundamentally different than that taking place in uinibalos at very high redshifts. which depeuds seusitivelv on the presence of initial Πο O'Shea 2005).," After reionization, such pristine star formation would be fundamentally different than that taking place in minihalos at very high redshifts, which depends sensitively on the presence of initial $H_2$ O'Shea 2005)."
+ Iu small objects. molecular Lydrogen is casily shotodissociated bv 11.2-13.6 eV photons. to which the muiverse is otherwise transparent.," In small objects, molecular hydrogen is easily photodissociated by 11.2-13.6 eV photons, to which the universe is otherwise transparent."
+" Thus the emission frou he first stars quickly destroved all avenues for cooling by nolecular line emission (Dekel Rees 1987: Παππά, Rees.4. Loeb 1997: Ciardi. Ferrara. Abel 2000)."," Thus the emission from the first stars quickly destroyed all avenues for cooling by molecular line emission (Dekel Rees 1987; Haiman, Rees, Loeb 1997; Ciardi, Ferrara, Abel 2000)."
+ This quickA raised the münimuiun virial temperature necessary to cool effectively to approximately 104 EK. although the precise value of this transition is the subject of debate Glover Brand 2001 Yoshida 2003) aud is somewhat depeudent on the level of the lugh-redshift N-rav background (Tainan. Rees. Loch 1996: Oh 2001: Machack. Bryan. Abel 2003).," This quickly raised the minimum virial temperature necessary to cool effectively to approximately $10^4$ K, although the precise value of this transition is the subject of debate Glover Brand 2001 Yoshida 2003) and is somewhat dependent on the level of the high-redshift X-ray background (Haiman, Rees, Loeb 1996; Oh 2001; Machack, Bryan, Abel 2003)."
+ Later ametaltree star formation is mach more straightforward. aud is confined to halos with Των710! I. Iun. this case efficient atomic line cooling establishes a dense locally-stable disk. within which nou-equilibrimu free electrons catalyze Πω.," Later metal-free star formation is much more straightforward, and is confined to halos with $T_{\rm vir} \geq 10^4$ K. In this case efficient atomic line cooling establishes a dense locally-stable disk, within which non-equilibrium free electrons catalyze $H_2$."
+" Unlike iu less massive halos. IT, formation in these objects is larecly impervious to feedback from external UW felds. due to the high deusities achieved by atomic cooling (Ol Taman 2002)."," Unlike in less massive halos, $H_2$ formation in these objects is largely impervious to feedback from external UV fields, due to the high densities achieved by atomic cooling (Oh Haiman 2002)."
+" As discussed in SSEOU3. these loucly regious of late primordial star formation may have already been detected as a suberoup of high-redshift Lyiman-alpha cutters. although distinguishiug these objects frou, Population II stars is extremely difficult (Dawson 2001)."," As discussed in SSF03, these lonely regions of late primordial star formation may have already been detected as a subgroup of high-redshift Lyman-alpha emitters, although distinguishing these objects from Population II stars is extremely difficult (Dawson 2004)."
+ Iu this investigation we explore an alternative approach to search for such VAIS: the direct detection of the resulting SNe., In this investigation we explore an alternative approach to search for such VMS: the direct detection of the resulting SNe.
+" While many of the AZ,=—100ML. stars would have died forming black holes. progenitors within the 110300ΔΕ. mass range cud thei lives in tremendously powerful pair-production sipernovac (PPSNe))."," While many of the $M_\star \geq 100 \msun$ stars would have died forming black holes, progenitors within the $140-260 \msun$ mass range end their lives in tremendously powerful pair-production supernovae )."
+ In these SNe. ele pair creation softens the equation state at the cud of central carbon burning. leading to rapid collapse. followed by explosive burning up to the ion eroup Ober. El Eid. Fricke 1983: Bond. Árnett. Carr 1981: Ποσο Woosley 2002. hereafter ITWO02).," In these SNe, $e^+ e^-$ pair creation softens the equation state at the end of central carbon burning, leading to rapid collapse, followed by explosive burning up to the iron group Ober, El Eid, Fricke 1983; Bond, Arnett, Carr 1984; Heger Woosley 2002, hereafter HW02)."
+ The more massive the star. the higher the temperature at bounce and the heavier the elements that are produced by uuclear fusion.," The more massive the star, the higher the temperature at bounce and the heavier the elements that are produced by nuclear fusion."
+ Iu all cases. the star is completely disrupted and both the ejection energies (~3100«10i eres) and lasses (110260 ML.) are enormous. leading to observed properties qualitatively different frou those of Tyvpe-Ia and typical core collapse SNe.," In all cases, the star is completely disrupted and both the ejection energies $\sim
+3-100 \times 10^{51}$ ergs) and masses $140-260 \msun$ ) are enormous, leading to observed properties qualitatively different from those of Type-Ia and typical core collapse SNe."
+ Iu turn. these properties. such as bright UV. colors. allow for efficient intermecdiate-redshitt searches for verv lnassive star formation.," In turn, these properties, such as bright UV colors, allow for efficient intermediate-redshift searches for very massive star formation."
+ Receutly Wise Abel (2005) studied the feasibility of detecting at 2~20 with theTelescope. using extremely approximate liehteurves with huninosities that were taken to be proportional to the total PONT ass.," Recently, Wise Abel (2005) studied the feasibility of detecting at $z \sim 20$ with the, using extremely approximate lightcurves with luminosities that were taken to be proportional to the total $^{56}$ Ni mass."
+ A similar study was carried out by Weinman Lilly (2005). who also estimated παπανο couuts at 2=5 using a sinele light curve that was taken from Deeer (2002). which unfortunately contained a normalization Limited studies of lightcurves fron: have also been carried out by Woosley Weaver (1982). who cousidered progenitors with onulv two masses aud found features simular to those described below. aud by Ierzig (1990) who modeled the disruption of a GL ML... Wolt-Ravet star. eeneratiug a relatively faint SN.," A similar study was carried out by Weinmann Lilly (2005), who also estimated number counts at $z=5$ using a single light curve that was taken from Heger (2002), which unfortunately contained a normalization Limited studies of lightcurves from have also been carried out by Woosley Weaver (1982), who considered progenitors with only two masses and found features similar to those described below, and by Herzig (1990) who modeled the disruption of a 61 $\msun$ Wolf-Rayet star, generating a relatively faint SN."
+ Hore we use the IKEPLER code to directly construct a suite of PPSN models that addresses the range of theoretical possibilities and accounts for the detailed ράσα) processes imvolved up until the poiut at which the SN becomes optically thin., Here we use the KEPLER code to directly construct a suite of PPSN models that addresses the range of theoretical possibilities and accounts for the detailed physical processes involved up until the point at which the SN becomes optically thin.
+" These curves allow us to identify the properties that 10st distinguish from Type Ia aud core collapse SNe. aud discuss thei detectability at 2<6 in curent and ""ture survevs."," These curves allow us to identify the properties that most distinguish from Type Ia and core collapse SNe, and discuss their detectability at $z \lesssim
+6$ in current and future surveys."
+ The outline of this paper is as follows., The outline of this paper is as follows.
+ Iu 82 we use he imuiplicit lvdrodvuamical code KEPLER to coustruct a suite of detailed PPSN progenitor models., In 2 we use the implicit hydrodynamical code KEPLER to construct a suite of detailed PPSN progenitor models.
+ Iu 825 we eenerate simulated spectra and Leltcurves from these uodels. lighliehting the key features that distinguish from other SNe.," In 3 we generate simulated spectra and lightcurves from these models, highlighting the key features that distinguish from other SNe."
+ In 81 owe apply these lishteurves O letermine the coustraiuts on obtainable bx existing and planned SNe searches., In 4 we apply these lightcurves to determine the constraints on obtainable by existing and planned SNe searches.
+ In 85 we discuss the xoperties of the most likely cuvirouments ofPPSNe.. and we sunnaaurize our results iu 86.," In 5 we discuss the properties of the most likely environments of, and we summarize our results in 6."
+ The progenitors of are confined to a range of initial masses that are large enough to collapse via the paii-production instability. but simall enough for this collapse to be reversed bv nuclear burning.," The progenitors of are confined to a range of initial masses that are large enough to collapse via the pair-production instability, but small enough for this collapse to be reversed by nuclear burning."
+ Iu the non-rotatiug case. these linuts are clearly defined aud well uuderstood muuercallv (Ober. El Eid. Fricke 1983: ΠΑΝΟΣ: Teeer 2003).," In the non-rotating case, these limits are clearly defined and well understood numerically (Ober, El Eid, Fricke 1983; HW02; Heger 2003)."
+ Iu stars above ~100ML... after central heliuua burning. the core reaches a sufficiently biel temperature aud eutropy to produce ¢!e pairs convertiug some of the internal enerev of the eas into rest mmass. Which softens the equation of state to an aciabatic iudex below the critical value of 1/3 (Barkat. Rakavy. Sack 1967: Dond. Arnett. Carr 1981).," In stars above $\sim100 \msun$, after central helium burning, the core reaches a sufficiently high temperature and entropy to produce $e^+e^-$ pairs, converting some of the internal energy of the gas into rest mass, which softens the equation of state to an adiabatic index below the critical value of 4/3 (Barkat, Rakavy, Sack 1967; Bond, Arnett, Carr 1984)."
+ The result is a runaway collapse that leads to explosive burning in the carbon-oxveen core., The result is a runaway collapse that leads to explosive burning in the carbon-oxygen core.
+ For stars with initial masses more than about L10. AL... the enerey released is sufficient to completely disrupt the star.," For stars with initial masses more than about 140 $\msun,$ the energy released is sufficient to completely disrupt the star."
+ A shock moves outward from the edee of the core. initiating the supernova outburst when it reaches the stellar surface.," A shock moves outward from the edge of the core, initiating the supernova outburst when it reaches the stellar surface."
+ Just above the 140. AL. lint. weal silicon burning occurs and onlv trace amounts of radioactive PONT are produced.," Just above the 140 $\msun$ limit, weak silicon burning occurs and only trace amounts of radioactive $^{56}$ Ni are produced."
+ It is this 70Ni that powers the late- supernova light curves., It is this $^{56}$ Ni that powers the late-time supernova light curves.
+ The amount of '?Ni produced Increases in larger progenitors. aud iu 260 ML. progenitors up to 50 AL. may be svuthesizecd.  100 times more than iu a typical Type Ia supernova.," The amount of $^{56}$ Ni produced increases in larger progenitors, and in 260 $\msun$ progenitors up to 50 $\msun$ may be synthesized, $\sim$ 100 times more than in a typical Type Ia supernova."
+ For stars with masses above 260 ALL. however. the ouset of photodisiuteeration in the center imposes an additional mstabilitv that collapses most of the star iuto a black hole (Boud. Arnett Carr 1981: Πο: Fever. Woosley. IHeeger 2001).," For stars with masses above 260 $\msun,$ however, the onset of photodisintegration in the center imposes an additional instability that collapses most of the star into a black hole (Bond, Arnett Carr 1984; HW02; Fryer, Woosley, Heger 2001)."
+" Pai-production superunovae are anmone most powerful thermonuclear explosious in the universe. with a total euereies ranging from ὃς1075 eres for a 110. AL. star (61 helimm M. core] to alinost 100:«1071 eres for a 260 M, "," Pair-production supernovae are among most powerful thermonuclear explosions in the universe, with a total energies ranging from $3 \times
+10^{51}$ ergs for a 140 $\msun$ star (64 helium $\msun$ core) to almost $100 \times 10^{51}$ ergs for a 260 $\msun$ "
+A detailed description of the reduction steps with the set of selection criteria. that must be satisfied [ον (he identification of acceptable bipolar magnetic regions is given in (he Appendix.,A detailed description of the reduction steps with the set of selection criteria that must be satisfied for the identification of acceptable bipolar magnetic regions is given in the Appendix.
+ After the IDL program has looped through all the 73.838 magnetogram files. ib has written an IDL save file with the extracted bipolar region information for each of the analvsed magnetograms.," After the IDL program has looped through all the 73,838 magnetogram files, it has written an IDL save file with the extracted bipolar region information for each of the analysed magnetograms."
+ These save files are then merged into one single master file that contains the extracted information for a total of of 245.733 identified bipolar regions. together with the relevant housekeeping data. like time of the magnetogram. heliographic longitude of the central meridian. heliographie latitude of disk center. Carrington rotation number. etc.," These save files are then merged into one single master file that contains the extracted information for a total of of 245,733 identified bipolar regions, together with the relevant housekeeping data, like time of the magnetogram, heliographic longitude of the central meridian, heliographic latitude of disk center, Carrington rotation number, etc."
+ From (his data set we remove the bipolar regions. for which the derived centers οἱ eravitv do not lie within r/r.«0.3 for polarities. in order (ο avoid errors that. can occur for regions that are located close to the solar limb.," From this data set we remove the bipolar regions, for which the derived centers of gravity do not lie within $r/r_\odot <0.8$ for polarities, in order to avoid errors that can occur for regions that are located close to the solar limb."
+ r/r.=0.8 corresponds to the cosine of the heliocentric angle pp=0.6., $r/r_\odot =0.8$ corresponds to the cosine of the heliocentric angle $\mu=0.6$.
+ This limb zone is conservatively chosen to be (his wide to minimize sources of error due to projection effects aud noise., This limb zone is conservatively chosen to be this wide to minimize sources of error due to projection effects and noise.
+ 85.654 regions. or with a rest of 160.079 bipolar regions. which satisfy all our criteria and enter into our statistical analvsis.," 85,654 regions, or with a rest of 160,079 bipolar regions, which satisfy all our criteria and enter into our statistical analysis."
+ For each of these bipolar magnetic regions the heliographic coordinates for the centers of gravity and the total magnetic fluxes of each of the two magnetic polarilies are determined and saved. accounting for the geometric foreshortening effect [or pixels away from disk center.," For each of these bipolar magnetic regions the heliographic coordinates for the centers of gravity and the total magnetic fluxes of each of the two magnetic polarities are determined and saved, accounting for the geometric foreshortening effect for pixels away from disk center."
+ Using spherical irigonomeltry. we connect the centers of gravity of the two polarities with a great circle and calculate the tilt angle of this great circle with respect to the parallel circle Chat defines the E-W direction. as well as (he polarity separation (in Man along the great. circle).," Using spherical trigonometry, we connect the centers of gravity of the two polarities with a great circle and calculate the tilt angle of this great circle with respect to the parallel circle that defines the E-W direction, as well as the polarity separation (in Mm along the great circle)."
+ For the N hemisphere the tilt is defined (o increase in the clockwise direction. for the 5 hemisphere in the counter-clockwise direction.," For the N hemisphere the tilt is defined to increase in the clockwise direction, for the S hemisphere in the counter-clockwise direction."
+ The zero point of the tilt angle represents, The zero point of the tilt angle represents
+The discovery of high galactic latitude cool (F. G. K) and hot (OB) post- AGB supergiants. e.g. HD 161796 (Parthasarathy Pottasch 1986) and LSII 434726 (Parthasarathy 1993) indicated that K. G. Ε. A. OB post- AGB supergiants form an evolutionary sequence in the transition region from the tip of the AGB into early stages of planetary nebulae (PNe).,"The discovery of high galactic latitude cool (F, G, K) and hot (OB) $-$ AGB supergiants, e.g. HD 161796 (Parthasarathy Pottasch 1986) and LSII $^{\circ}$ 26 (Parthasarathy 1993) indicated that K, G, F, A, OB $-$ AGB supergiants form an evolutionary sequence in the transition region from the tip of the AGB into early stages of planetary nebulae (PNe)."
+ Since then. several cool and hot post-AGB candidates have been identified (Pottasch Parthasarathy 1988. Parthasarathy Pottasch 1989. Garcíaa- Lario et al.," Since then, several cool and hot $-$ AGB candidates have been identified (Pottasch Parthasarathy 1988, Parthasarathy Pottasch 1989, $-$ Lario et al."
+ 19972; Parthasarathy et al., 1997a; Parthasarathy et al.
+ 20003)., 2000a).
+ Gauba Parthasarathy (2003. 2004) analysed the UV spectra and circumstellar dust envelopes of several hot post-AGB stars from the above lists. including LS III +52°24.," Gauba Parthasarathy (2003, 2004) analysed the UV spectra and circumstellar dust envelopes of several hot $-$ AGB stars from the above lists, including LS III $^{\circ}$ 24."
+ High resolution optical spectra of only a few hot post- AGB stars have been studied (Sarkar et al., High resolution optical spectra of only a few hot $-$ AGB stars have been studied (Sarkar et al.
+ 2005 and references therein)., 2005 and references therein).
+ To ascertain their evolutionary nature. atmospheric parameters and chemical compositions. there ts a clear need to study a bigger sample of these objects.," To ascertain their evolutionary nature, atmospheric parameters and chemical compositions, there is a clear need to study a bigger sample of these objects."
+ The optically bright B-type star. LS II] +52°24 identified with the IR source IRAS 22023+5249 (hereafter. 122023) has far-IR colors similar to PNe (Table 1).," The optically bright $-$ type star, LS III $^{\circ}$ 24 identified with the IR source IRAS 22023+5249 (hereafter, I22023) has $-$ IR colors similar to PNe (Table 1)."
+ It is listed in Wackerling’s (1970) Catalog of early-type emission-line stars., It is listed in Wackerling's (1970) Catalog of $-$ type $-$ line stars.
+ In this paper. we present the first high resolution optical spectrum of 122023.," In this paper, we present the first high resolution optical spectrum of I22023."
+ Recent. ground-based. high resolution. near-IR images. have shown H» emission arising close to the central star. possibly in a bipolar," Recent, $-$ based, high resolution, $-$ IR images, have shown $_{2}$ emission arising close to the central star, possibly in a bipolar"
+angle 7 increases. because a=150°and30° correspond to the collision angle at superior conjunction (SUPC) and inferior conjunction (INFC). respectively. when ¢=60° (see Figure 1)),"angle $i$ increases, because $\alpha = 150^{\circ} \text{ and } 30^{\circ}$ correspond to the collision angle at superior conjunction (SUPC) and inferior conjunction (INFC), respectively, when $i = 60^{\circ}$ (see Figure \ref{modpic}) )."
+ In the same manner. a=120°and607 correspond to (hat al SUPC and INFC when i= 307. and a=105°and75° correspond to that of 7=15° (see Figure 1)).," In the same manner, $\alpha = 120^{\circ} \text{ and } 60^{\circ}$ correspond to that at SUPC and INFC when $i = 30^{\circ}$ , and $\alpha = 105^{\circ} \text{ and } 75^{\circ}$ correspond to that of $i = 15^{\circ}$ (see Figure \ref{modpic}) )."
+ Moreover. the ratio of GeV flux to TeV flux is large when p=2.5. while it is small when p=2.0 It is inevitable for a hieh enerev photon to be absorbed in the stellar radiation fielcl.," Moreover, the ratio of GeV flux to TeV flux is large when $p = 2.5$, while it is small when $p=2.0$ It is inevitable for a high energy photon to be absorbed in the stellar radiation field."
+ Figure 3. shows the optical depth lor 55 absorption., Figure \ref{tau} shows the optical depth for $\gamma \gamma$ absorption.
+ According to this figure. the optical depth increases as (he angle a increases. and the photon energv al the maximum 7 shifts from ~ I TeV to ~ 100 GeV. Furthermore. when 7<30° (60°<a«120°). 7 exceeds unity in anv orbital phase.," According to this figure, the optical depth increases as the angle $\alpha$ increases, and the photon energy at the maximum $\tau$ shifts from $\sim$ 1 TeV to $\sim$ 100 GeV. Furthermore, when $i \le 30^{\circ}$ $60^{\circ} < \alpha < 120^{\circ}$ ), $\tau$ exceeds unity in any orbital phase."
+ Thus. +> absorption is very efficient in LS5039. so Chat electromagnetic cascade develops in the case where the svnchrotron energy loss of electrons can be ignored.," Thus, $\gamma \gamma$ absorption is very efficient in LS5039, so that electromagnetic cascade develops in the case where the synchrotron energy loss of electrons can be ignored."
+ The emission [rom the svstem L55039 can be considered as a stationary one if a divided phase has a range of a tenth of one orbital period., The emission from the system LS5039 can be considered as a stationary one if a divided phase has a range of a tenth of one orbital period.
+" This is meaninelul when electrons whose enerey range is 2.x10*«5,< cool completely on much shorter time scale than that jxhase range.", This is meaningful when electrons whose energy range is $3\times 10^3 < \gamma_e < 10^8$ cool completely on much shorter time scale than that phase range.
+" The ratio of IC cooling time for the electrons in the Thomson regime5 /icLVL toa tenth of an orbital period T,/10~3.4x10! s is On the other hand. high energy electrons cool in the Klein-Nishina regime. sothat where Nounο.ο/YTeV)""'s is the cooling time of electrons in the Ixlein-Nishina regime (Aharonianetal. 2006b).."," The ratio of IC cooling time for the electrons in the Thomson regime $t_{\text{IC,T}}$ toa tenth of an orbital period $T_{\text{orb}} / 10 \sim 
+3.4\times 10^4$ s is On the other hand, high energy electrons cool in the Klein-Nishina regime, sothat where $t_{\text{IC,KN}} \sim 22\ (a/a_{\text{peri}})^2(E_e/1\text{TeV})^{0.7} \text{s}$ is the cooling time of electrons in the Klein-Nishina regime \citep{aha06b}. ."
+" Thus. the electrons with 3x10*<>,<10 "," Thus, the electrons with $3\times 10^3 < \gamma_e < 10^8$ "
+of Schoenibs&Zoeschiuger(1990). that thefine of optical bursts (without simultancous X-ray coverage) was correlated with the orbital phase in the same seuse.,of \citet{Schoembs:1990a} that the of optical bursts (without simultaneous X-ray coverage) was correlated with the orbital phase in the same sense.
+ Obviously before extensive interpretation. more sophisticated analysis is warranted to verity these results.," Obviously before extensive interpretation, more sophisticated analysis is warranted to verify these results."
+ The quality of the CCFs are superb aud the major concern relates to the mieauingfuluess of the lag derived., The quality of the CCFs are superb and the major concern relates to the meaningfulness of the lag derived.
+ Fundamental to the CCF approach is the assumption that the optical/UV is simply a lageed version of the N-vavs. but this is clearly not the case im the burst hehtcurves. and is not expected theoretically.," Fundamental to the CCF approach is the assumption that the optical/UV is simply a lagged version of the X-rays, but this is clearly not the case in the burst lightcurves, and is not expected theoretically."
+ Several factors nav be important., Several factors may be important.
+ One is light travel tines., One is light travel times.
+ These will not only produce a mean lag. but will tend to sincar out the signal. as different lags are produced bv cifferent reprocessing sites (OBrienetal.2002)..," These will not only produce a mean lag, but will tend to smear out the signal, as different lags are produced by different reprocessing sites \citep{OBrien:2002a}."
+ A second factor having a simular effect is that there may ve a finite diffusion fine associated with reprocessing: if N-ravs deposit cnerev at a significant optical deptl it cannot be re-radiated instantly. but there will be urther lageius aud smnemuring as the energv cdiffuses outward.," A second factor having a similar effect is that there may be a finite diffusion time associated with reprocessing; if X-rays deposit energy at a significant optical depth it cannot be re-radiated instantly, but there will be further lagging and smearing as the energy diffuses outward."
+ Finally. the tails of the bursts me clearly different. indicating non-linear relationships between the iehteurves.," Finally, the tails of the bursts are clearly different, indicating non-linear relationships between the lightcurves."
+" This occurs because the bursts have a large zaüuplitude aud substantially change the temperature of voth the neutron star aud the reprocessing site,", This occurs because the bursts have a large amplitude and substantially change the temperature of both the neutron star and the reprocessing site.
+ As a site cools. the peak of the euütted spectrum moves to ower energies.," As a site cools, the peak of the emitted spectrum moves to lower energies."
+ In the case of the reprocessed light. it uoves out of the £u-UV. aud iuto the optical. havine he effect of accelerating the £u-UV decay rate (relative o the bolometric decay) and suppressing the optical decay.," In the case of the reprocessed light, it moves out of the far-UV, and into the optical, having the effect of accelerating the far-UV decay rate (relative to the bolometric decay) and suppressing the optical decay."
+ None of these factors are accounted for m cross-correlations. so a more sophisticated model is uecded.," None of these factors are accounted for in cross-correlations, so a more sophisticated model is needed."
+ For example. lags can be over or uudoer-estinmated if the timing characteristics of the two ligbiteurves differ (Ίνοσι2003).. reflecting the more general problem that in this case there is not a uuiquely defined lag.," For example, lags can be over or under-estimated if the timing characteristics of the two lightcurves differ \citep{Koen:2003a}, reflecting the more general problem that in this case there is not a uniquely defined lag."
+ There is a range. from which a CCF favors an average value.," There is a range, from which a CCF favors an average value."
+ We next attempt to deconvolve the lehtcurves to obtain the transfer function between Nav and optical wavelengths independent of the width of the burst., We next attempt to deconvolve the lightcurves to obtain the transfer function between X-ray and optical wavelengths independent of the width of the burst.
+ This will be more readily interpreted than the CCF which is much broader than the trausfer function., This will be more readily interpreted than the CCF which is much broader than the transfer function.
+ We iuitiallv use the maxima eutropy method (MEM: Ilorueetal.10011: Horne199 L3). which does not require au assunied model for the trausfer function. bevoud an assumed default used in calculating eutropy.," We initially use the maximum entropy method (MEM; \citealt{Horne:1991a}; \citealt{Horne:1994a}) ), which does not require an assumed model for the transfer function, beyond an assumed default used in calculating entropy."
+ Iu principle. the NEM method cau then resolve distiuct. 5ub-structures. if they are present. for example the disk aud companion star.," In principle, the MEM method can then resolve distinct sub-structures, if they are present, for example the disk and companion star."
+ For cach burst we use our ταν flux lighteurves as the driver signal aud attempt to fit the optical or UV. echo over the range 50 to |300 9x relative to the burst rise., For each burst we use our X-ray flux lightcurves as the driver signal and attempt to fit the optical or UV echo over the range $-50$ to $+300$ s relative to the burst rise.
+" The transfer fiction was calculated for lags of 50 ss to |200 «x. We use a narrow Caussian default to minimize euforced suoothiue of the derived transfer ΠιοΊο»,", The transfer function was calculated for lags of $-50$ s to $+200$ s. We use a narrow Gaussian default to minimize enforced smoothing of the derived transfer functions.
+ The results are shown in Fie. 5., The results are shown in Fig. \ref{MEMFig}.
+ Negligible response was seen before —10 ss. aud slow declines after |50-ss: these regions are not shown.," Negligible response was seen before $-10$ s, and slow declines after $+50$ s; these regions are not shown."
+" As expected. the derived transfer fuuctious are 1uuch narrower than the (Εν,"," As expected, the derived transfer functions are much narrower than the CCFs."
+ All bursts suggest a siugle-peaked response peaking at a lag of a few seconds. very close to the locations of the maxima of the corresponding CCFs.," All bursts suggest a single-peaked response peaking at a lag of a few seconds, very close to the locations of the maxima of the corresponding CCFs."
+ Besices the main peak. all the trausfer functions have a tail extending to larger lags.," Besides the main peak, all the transfer functions have a tail extending to larger lags."
+ This may be a real effect. as models of reprocessing of bursts by stellar atinospleres do predict a siuall amount of the energv emerges with large delavs (Conüuskyv.Loudon.&lei 1987: see Section 7.2. for more discussion).," This may be a real effect, as models of reprocessing of bursts by stellar atmospheres do predict a small amount of the energy emerges with large delays \citealt{Cominsky:1987a}; see Section \ref{DiffusionSection} for more discussion)."
+ We caunot sav with confidence that this effect is real. however. as there is another explanation for it.," We cannot say with confidence that this effect is real, however, as there is another explanation for it."
+ As noted earlier. the baudpass-limited respouse is non-linear. aud in particular the optical burst decay is much slower than that in X-rays because the reprocessed spectruni shifts to longer waveleugths Guto the optical baud) as it cools.," As noted earlier, the bandpass-limited response is non-linear, and in particular the optical burst decay is much slower than that in X-rays because the reprocessed spectrum shifts to longer wavelengths (into the optical band) as it cools."
+ A method which assumes a linear response will require an extended tail to reproduce this., A method which assumes a linear response will require an extended tail to reproduce this.
+ Iu support of this interpretation. the effect is unich weaker im the transfer fiction derived ποια UV data as expected.," In support of this interpretation, the effect is much weaker in the transfer function derived from UV data as expected."
+to be approximated by a power law) of the absorber sample.,to be approximated by a power law) of the absorber sample.
+ The geometric mean cornposite spectrum for an absorber sample was generated by shifting the spectra to the absorber rest frames and taking their geometrice mean., The geometric mean composite spectrum for an absorber sample was generated by shifting the spectra to the absorber rest frames and taking their geometric mean.
+" The non-absorber sample was compiled by selecting «one non-absorber QSO (not having any absorption lines in its spectrum) for each QSO in the absorber sample. having nearly the same z,.,, and M; (see YOO for details)."," The non-absorber sample was compiled by selecting one non-absorber QSO (not having any absorption lines in its spectrum) for each QSO in the absorber sample, having nearly the same $_{em}$ and $_i$ (see Y06 for details)."
+ To generate the geometric mean of the non-absorber sample. each non-absorber spectrum was first shifted to the rest frame of the absorber in the corresponding absorber QSO.," To generate the geometric mean of the non-absorber sample, each non-absorber spectrum was first shifted to the rest frame of the absorber in the corresponding absorber QSO."
+" As long as the spectroscopic properties of the QSOs in the absorber and non-absorber samples are statistically equivalent (which is ensured by their having similar z,.,,, and M,). except for the effects of the intervening absorption systems. dividing the absorber composite by the non-absorber composite will vield the relative extinction curve in the absorber frame."," As long as the spectroscopic properties of the QSOs in the absorber and non-absorber samples are statistically equivalent (which is ensured by their having similar $_{em}$ and $_i$ ), except for the effects of the intervening absorption systems, dividing the absorber composite by the non-absorber composite will yield the relative extinction curve in the absorber frame."
+ In doing this. the non-absorber spectrum is scaled so that the average flux densities in the reddest 100 of both spectra are equal.," In doing this, the non-absorber spectrum is scaled so that the average flux densities in the reddest 100 of both spectra are equal."
+ The extinction curve can be fitted with the SMC. LMC or Milky Way (MW) extinction curves to obtain the average E(BV) for the absorrber sample.," The extinction curve can be fitted with the SMC, LMC or Milky Way (MW) extinction curves to obtain the average $E(B-V)$ for the absorber sample."
+ For generating both types of composites. we used only parts of spectra having A>1250A in the rest frame of QSOs in order to avoid any contamination by the Lyman alpha forest in the spectra.," For generating both types of composites, we used only parts of spectra having $\lambda > 1250$ in the rest frame of QSOs in order to avoid any contamination by the Lyman alpha forest in the spectra."
+ Although FPIO use the method of composite spectra as advocated by YO. their procedure differed from YO6 in respect of the use of the matching non-absorber sample.," Although FP10 use the method of composite spectra as advocated by Y06, their procedure differed from Y06 in respect of the use of the matching non-absorber sample."
+ They used the sample of QSOs not having DLAs (non-DLA sample) as their non-absorber sample. arguing that any other dower H I column density) systems will be found with equal probability in the absorber as well as the non-DLA sample and the effect of these systems should cancel out.," They used the sample of QSOs not having DLAs (non-DLA sample) as their non-absorber sample, arguing that any other (lower H I column density) systems will be found with equal probability in the absorber as well as the non-DLA sample and the effect of these systems should cancel out."
+ We use both types of comparison samples: one comprising of QSOs having no absorbers along the line of sight (the non-absorber sample) and the other comprising of QSOs having no DLAs along the line of sight irrespective of the presence or absence of other absorption systems (the non-DLA sample}., We use both types of comparison samples: one comprising of QSOs having no absorbers along the line of sight (the non-absorber sample) and the other comprising of QSOs having no DLAs along the line of sight irrespective of the presence or absence of other absorption systems (the non-DLA sample).
+ The composite spectra of the non-absorber sample would be completely devoid of reddening due to intervening absorbers and a comparison with the composite of QSOs having DLAs along their lines of sight should reveal any weak reddening present in these QSOs., The composite spectra of the non-absorber sample would be completely devoid of reddening due to intervening absorbers and a comparison with the composite of QSOs having DLAs along their lines of sight should reveal any weak reddening present in these QSOs.
+ On the other, On the other
+We will neglect the small mass difference between neutrons and protons and write simply 20 for the nucleon mass. which we take equal to the atomic mass unit.,"We will neglect the small mass difference between neutrons and protons and write simply $m$ for the nucleon mass, which we take equal to the atomic mass unit."
+ For consistency. we thus take the electron mass 1)=0., For consistency we thus take the electron mass $m^e=0$.
+ We also neglect the muon mass except at the microscopic scale for calculating the internal energy. density., We also neglect the muon mass except at the microscopic scale for calculating the internal energy density.
+" In the limit of small currents. the Lagrangian density A can be decomposed into a dynamical part Aiea. which depends on the particle currents nf. and a statie part in. which depends only on the particle densities given by The dynamical Lagrangian density (neglecting the contribution of the leptons) can be written as a quadratic form of the nucleon currents The coellicients of the symmetric mobility matrix AK"". (q.q'=n.p for neutrons. protons respectively) are functions of the nucleon densities and will be evaluated in Section δι."," In the limit of small currents, the Lagrangian density $\Lambda$ can be decomposed into a dynamical part $\Lambda_{\rm dyn}$, which depends on the particle currents $n_{_{\rm X}}^{\, \mu}$ , and a static part $\Lambda_{\rm ins}$ which depends only on the particle densities given by The dynamical Lagrangian density (neglecting the contribution of the leptons) can be written as a quadratic form of the nucleon currents The coefficients of the symmetric mobility matrix ${\cal K}^{q q^\prime}$ $q,q^\prime=n,p$ for neutrons, protons respectively) are functions of the nucleon densities and will be evaluated in Section \ref{sect.micro}."
+" Due to the Galilean invariance. the matrix elements are related to each other by ‘Taking the partial derivative of the above equation with respect to the nucleon density leads to the following identity Since the left-hancl side of Eq.(14)) has to be negative for any neutron ancl proton densities. we obtain the following inequalities The non-diagonal coellicient K""=KP"" accounts for the non-dissipative entrainment effects and arises from the strong interactions between nucleons."," Due to the Galilean invariance, the matrix elements are related to each other by Taking the partial derivative of the above equation with respect to the nucleon density leads to the following identity Since the left-hand side of \ref{eq.gal.inv2}) ) has to be negative for any neutron and proton densities, we obtain the following inequalities The non-diagonal coefficient ${\cal K}^{np}={\cal K}^{pn}$ accounts for the non-dissipative entrainment effects and arises from the strong interactions between nucleons."
+" ΙΓ the nucleons could. be treated as ideal Fermi. gases. this coellicient would simply vanish K'""=0."," If the nucleons could be treated as ideal Fermi gases, this coefficient would simply vanish ${\cal K}^{np}=0$."
+" The static contribution Aji. is related to the static internal (non gravitational) energy density Ci, by where p is the total barvon mass density and Cu,=pO is the gravitational potential energy. density. © being the gravitational scalar potential."," The static contribution $\Lambda_{\rm ins}$ is related to the static internal (non gravitational) energy density $U_{\rm ins}$ by where $\rho$ is the total baryon mass density and $U_{\rm pot}=\rho \phi$ is the gravitational potential energy density, $\phi$ being the gravitational scalar potential."
+ Xs cliscussed previously nucleons are not relativistic neither at the microscopic scale nor at the macroscopic scale., As discussed previously nucleons are not relativistic neither at the microscopic scale nor at the macroscopic scale.
+ Nevertheless. in order to prepare the generalization to relativistic I[uids in the next section. the nucleon rest mass energy density is included in C.," Nevertheless, in order to prepare the generalization to relativistic fluids in the next section, the nucleon rest mass energy density is included in $U_{\rm ins}$."
+ lt can be shown that in the Newtonian framework this extra term has no effect on the variational principle (thereby. on the dynamical equations of the IEuids) provided the total mass current is conserved (?).., It can be shown that in the Newtonian framework this extra term has no effect on the variational principle (thereby on the dynamical equations of the fluids) provided the total mass current is conserved \citep{carter-06}.
+ Let us point out that the total internal (non 3gravitational) energy.S. density Ci is not simply ABgiven by Ci. but contains an additional entrainment term cdoefined bv Cont=(aymnCin. where Cayn=AWNu and is the total kinetic energy density.," Let us point out that the total internal (non gravitational) energy density $U_{\rm int}$ is not simply given by $U_{\rm ins}$ , but contains an additional entrainment term defined by $U_{\rm ent}= U_{\rm dyn} - U_{\rm kin}$, where $U_{\rm dyn}=\Lambda_{\rm dyn}$ and is the total kinetic energy density."
+ We thus have Ci=Cus|Uu 1n, We thus have $U_{\rm int}=U_{\rm ins}+U_{\rm ent}$.
+" summary. the non-relativistic Lagrangian density of the Duids is therefore given by where Aga, and Ais are given by Eqs. (12))"," In summary, the non-relativistic Lagrangian density of the fluids is therefore given by where $\Lambda_{\rm dyn}$ and $\Lambda_{\rm ins}$ are given by Eqs. \ref{eq.lambda.dyn}) )"
+ and (16)) respectively., and \ref{eq.lambda.ins}) ) respectively.
+ The dvnamical equations of the gravitational field can be obtained from the variational principle bv simplyadding the contribution C beingo the ogravitational constant., The dynamical equations of the gravitational field can be obtained from the variational principle by simplyadding the contribution $G$ being the gravitational constant.
+" Consideringo variations of the LagrangianD density Av,=uA|Ayurlwith respect to © leads to the usual Poisson's equation", Considering variations of the Lagrangian density $\Lambda_{\rm tot}=\Lambda+\Lambda_{\rm grf}$with respect to $\phi$ leads to the usual Poisson's equation
+could simply be due to them having had less time to deplete their newly-acquirecl eas reservoirs.,could simply be due to them having had less time to deplete their newly-acquired gas reservoirs.
+ We alleviate this potential problem by requiring the dust lane control and sample galaxies to span the same range in starburst age. LOO<{ο«400 Myr.," We alleviate this potential problem by requiring the dust lane control and sample galaxies to span the same range in starburst age, $100 < t_2 < 400$ Myr."
+ As Fig. S((, As Fig. \ref{fig:SFtimescales}( (
+a) shows. this preferentially picks out the AGN population (as defined in the DIE. plot). and. should. be representative of recent mereing events in both the dust lane and overall earlv-tvpe populations.,"a) shows, this preferentially picks out the AGN population (as defined in the BPT plot), and should be representative of recent merging events in both the dust lane and overall early-type populations."
+ We note that the stellar mass distributions of the two subsamples are very. similar., We note that the stellar mass distributions of the two subsamples are very similar.
+ Fig. O(, Fig. \ref{fig:gasAccretion}( (
+a) shows that. even alter correcting Lor the dilference in starburst age distributions. dust lane galaxies exhibit star formation rates that are higher by a factor of 3 than the control sample.,"a) shows that, even after correcting for the difference in starburst age distributions, dust lane galaxies exhibit star formation rates that are higher by a factor of 3 than the control sample."
+ HE dust. corrections are important. this number is a lower limit. and the actual star formation rates in dust lane earlv-tvpe galaxies may be even higher.," If dust corrections are important, this number is a lower limit, and the actual star formation rates in dust lane early-type galaxies may be even higher."
+ Fig. 9((, Fig. \ref{fig:gasAccretion}( (
+b) shows the distribution of the black hole accretion rate. in Edcington units.,"b) shows the distribution of the black hole accretion rate, in Eddington units."
+ Bolometric luminosity is derived. using the O14]] line luminosity. and. the usual correction factor of 3500 is applied.," Bolometric luminosity is derived using the ] line luminosity, and the usual correction factor of 3500 is applied."
+ The black hole mass is estimated. via its velocity dispersion2002)., The black hole mass is estimated via its velocity dispersion.
+. Dust lane carly types show aceretion rates that are again a [actor of three higher than the control sample. consistent. with the discrepancy in the star formation rates.," Dust lane early types show accretion rates that are again a factor of three higher than the control sample, consistent with the discrepancy in the star formation rates."
+ A final confirmation of this picture comes from emission-line AGN fractions in the two samples., A final confirmation of this picture comes from emission-line AGN fractions in the two samples.
+ Almost all (δὲ per cent) dust lane early types are classified as either σον[οί LINER. or transition objects. compared to only. 24 » cent for the control sample.," Almost all (84 per cent) dust lane early types are classified as either Seyfert, LINER or transition objects, compared to only 24 per cent for the control sample."
+ More than 60. per cent of the objects in the control sample have emission. lines that are too weak to place them on the BPP diagnostic. sugeesting that these objects undergo very low-level (if any) star formation and AGN activity. despite having relatively voung starbursts.," More than 60 per cent of the objects in the control sample have emission lines that are too weak to place them on the BPT diagnostic, suggesting that these objects undergo very low-level (if any) star formation and AGN activity, despite having relatively young starbursts."
+ This is in stark contrast to dust. lane hosts with the same starburst ages. and confirms that the dust. features are simply associated with galaxies that have acquired large amounts of gas. most likely in a recent merger.," This is in stark contrast to dust lane hosts with the same starburst ages, and confirms that the dust features are simply associated with galaxies that have acquired large amounts of gas, most likely in a recent merger."
+ We consider the origin and. evolution of dust. lane carly-tvpe galaxies., We consider the origin and evolution of dust lane early-type galaxies.
+ In a companion paper. we showed that these objects are preferentially located in poor environments. are morphologically clisturbed. and have higher star formation rates than the mean for earlv-tvpe galaxies.," In a companion paper, we showed that these objects are preferentially located in poor environments, are morphologically disturbed, and have higher star formation rates than the mean for early-type galaxies."
+ In this work we addressed the issue of the origin of these objects., In this work we addressed the issue of the origin of these objects.
+ By emploving optical (SDSS) and racio (ΕΤ and NVSS) surveys we examined [GN properties of dust [ane earlv-tvpe galaxies ancl a relevant control sample of, By employing optical (SDSS) and radio (FIRST and NVSS) surveys we examined AGN properties of dust lane early-type galaxies and a relevant control sample of early-type.
+ Our findings are as follows., Our findings are as follows.
+The time evolution in case A (Figs. 1..2)),"The time evolution in case A (Figs. \ref{fig:A1}, \ref{fig:A2}) )"
+ is analogous to the models of Lelloch Lazarell (10994)., is analogous to the models of Lefloch Lazareff \shortcite{lefloch94}.
+". One can distinguish between four cdillerent phases: Placing our models with H5,Spo294.71 and m.c40 into the Figure 14 in Bertoldi Alelxee (1990). (reprinted here as Fig. 5))."," One can distinguish between four different phases: Placing our models with $R_{\rm pc}^2/S_{49} \simeq 4.71$ and $m_\odot
+\simeq 40$ into the Figure 14 in Bertoldi McKee \shortcite{bertoldi90} (reprinted here as Fig. \ref{fig:bertoldi}) ),"
+ they lic far in the regime where gravitational instability is expected., they lie far in the regime where gravitational instability is expected.
+ Simulations with smooth initial conditions including self-eravity. verily this conclusion., Simulations with smooth initial conditions including self-gravity verify this conclusion.
+ The. overall evolution is indistinguishable to the one in case A. until the point of maximum compression is reached.," The overall evolution is indistinguishable to the one in case A, until the point of maximum compression is reached."
+ At that moment eravitational collapse sets in. which can be seen in the sudden rise in the maximum. density of the system (see Fie. 6)).," At that moment gravitational collapse sets in, which can be seen in the sudden rise in the maximum density of the system (see Fig. \ref{fig:maxdens}) )."
+ 1n these models. the evolution fundamentally cülfers fron the one in model D. During the compression phase. density enhancements are amplified in the shocked laver. similar to what is observed in the simulations by Elmeereen et al.," In these models, the evolution fundamentally differs from the one in model B. During the compression phase, density enhancements are amplified in the shocked layer, similar to what is observed in the simulations by Elmegreen et al."
+ (1995). (Figs. 3.. 4))," \shortcite{elmegreen95} (Figs. \ref{fig:C1}, \ref{fig:C2}) )"
+ and those of Williams et al. (, and those of Williams et al. (
+2001) (lig.,2001) (Fig.
+ 6D)., 6f).
+ They are reminiscent of the observed. small-scale structures seen in MIG (Llester et al., They are reminiscent of the observed small-scale structures seen in M16 (Hester et al.
+ 1996: White et al., 1996; White et al.
+ 1999)., 1999).
+ Contrary to case D. no collapse occurs at the moment of maximum compression.," Contrary to case B, no collapse occurs at the moment of maximum compression."
+ We rather observe an evolution analogous to the phases 3 and 4 of case A. After a delay of 300 kvrs. the laver starts fragmenting by gravitational instability.," We rather observe an evolution analogous to the phases 3 and 4 of case A. After a delay of 300 kyrs, the layer starts fragmenting by gravitational instability."
+ As already pointed out in Chapter 2.1.. the initial conditions were chosen such that the integration volume contains ~1 Jeans mass.," As already pointed out in Chapter \ref{chapt:inconds}, the initial conditions were chosen such that the integration volume contains $\sim 1$ Jeans mass."
+ Looking at the Jeans mass Aj and Jeans length Ly. it follows that Aj as well as Zip drop with increasing density.," Looking at the Jeans mass $M_{\rm J}$ and Jeans length $L_{\rm J}$, and the freefall timescale, it follows that $M_{\rm J}$ as well as $T_{\rm ff}$ drop with increasing density."
+ At the moment of highest. compression in. model A. regions with densities higher than a threshold. density fu5:101T'gem5 contain. 5.3⋅AJ..," At the moment of highest compression in model A, regions with densities higher than a threshold density $\rho_{\rm m} = 5 \cdot 10^{-17} {\rm g \, cm}^{-3}$ contain $5.3\; M_\odot$."
+ For∖ this. py. one finds. . Lim1-10?pe and Yip=10kvrs.," For this $\rho_{\rm m}$, one finds $M_{\rm J} = 2.7 \, {\rm
+M}_\odot$, $L_{\rm J} = 7 \cdot 10^{-3}\, {\rm pc}$ and $T_{\rm ff} = 10 \, {\rm kyrs}$."
+ Compared to these values. the compressed region in the tip of the cloud contains 2 Jeans masses and is large enough. and Zip is small compared to other timescales in the svstem.," Compared to these values, the compressed region in the tip of the cloud contains 2 Jeans masses and is large enough, and $T_{\rm ff}$ is small compared to other timescales in the system."
+ One can thus expect that the elobule is unstable to gravitational collapse at this moment., One can thus expect that the globule is unstable to gravitational collapse at this moment.
+ This explains why we observe collapse when including gravity in model D. One should however use the Jeans masses criterion with care. as the cases € and D illustrate.," This explains why we observe collapse when including gravity in model B. One should however use the Jeans masses criterion with care, as the cases C and D illustrate."
+ Lt only considers the balance between internal energy. and gravity and is thus a static approach. neglecting the ellects of internal motion.," It only considers the balance between internal energy and gravity and is thus a static approach, neglecting the effects of internal motion."
+ It is às such only applicable in very lew cases., It is as such only applicable in very few cases.
+ Since the initial conditions are chosen such that they contain 1. Jeans mass for all cases. the stabilization of cases € and D cannot be caused. by gas pressure alone.," Since the initial conditions are chosen such that they contain $\sim 1$ Jeans mass for all cases, the stabilization of cases C and D cannot be caused by gas pressure alone."
+ The only mechanism which ολους in the simulations and which can be responsible is undirected gas motion. i.e. disordered. kinetic Onorgv.," The only mechanism which differs in the simulations and which can be responsible is undirected gas motion, i.e. disordered kinetic energy."
+ 7, Fig.
+ shows the total amount of kinetic energy Lu relative to the center of mass compared το the eravitational energy. fon. in the system for case C. During the compression phase. more and more material is collected bv the post-shock [aver and. accelerated: to. velocities," \ref{fig:energies} shows the total amount of kinetic energy $E_{\rm kin}$ relative to the center of mass compared to the gravitational energy $E_{\rm grav}$ in the system for case C. During the compression phase, more and more material is collected by the post-shock layer and accelerated to velocities"
+"These rates (n mits cin? ty are: 52=9.5«1019, VE710 Mand 54=L71027.","These rates (in units $^3$ $^{-1}$ ) are: $\gamma^{\small{H}}_{21}=9.5\times10^{-10}$, $\gamma^{\small{H}}_{31}=5.7\times10^{-10}$ and $\gamma^{\small{H}}_{32}=4.7\times10^{-10}$."
+" Tu the absence of information ou hieher-.7 levels in the °D umltiplet. we have adopted de-excitation rates of 5.&1019 eni? 1, constant with temperature."," In the absence of information on $J$ levels in the $^6$ D multiplet, we have adopted de-excitation rates of $5 \times 10^{-10}$ $^3$ $^{-1}$, constant with temperature."
+" The critical density of the first excited state. Όρο, resulting frou IT? de-excitation Is He,=2.2ς10°ín.7 "," The critical density of the first excited state, $^6$ $_{7/2}$, resulting from $^{\circ}$ de-excitation is $n_{\rm cr} = 2.24 \times 10^6~{\rm cm}^{-3}$."
+We computed the radiative cooling rate from ποια the statistical equilibriunn populations of the excited fine-structure levels. multiplied by their radiative decay rates and line energies.," We computed the radiative cooling rate from from the statistical equilibrium populations of the excited fine-structure levels, multiplied by their radiative decay rates and line energies."
+ At T<SOO0 Is. excitation of lis no longer an effiieut cooling process.," At $T < 8000$ K, excitation of is no longer an efficient cooling process."
+ Similarly. at LF<60 E. cooling by pure rotational lines from Dhecomes ineffective.," Similarly, at $T < 60$ K, cooling by pure rotational lines from becomes ineffective."
+ At very low temperatures (Z<60 K. WD cooling can be effective. despite the simall primordial abundance. rin226 ppm. owing to its weak dipole transitions (??)..," At very low temperatures $T < 60$ K, HD cooling can be effective, despite the small primordial abundance, $_{\rm prim} \approx 26$ ppm, owing to its weak dipole transitions \citep{LS84,LSD02}."
+ At T«1000 K. in the absence of external photoionizing sources. the hydrogen eas is completly neutral. and low-cucrey fine-structure transitions ofΠΠ.I]..H].. and II].. excited by IP collisions. become the most miportaut coolauts. (," At $T < 1000$ K, in the absence of external photoionizing sources, the hydrogen gas is completly neutral, and low-energy fine-structure transitions of, and , excited by $^{\circ}$ collisions, become the most important coolants. ("
+We asstuue that C. $i. aud Fe are phlotoionized bv zunmnbieut UV radiation from newly formed hot stars.),"We assume that C, Si, and Fe are photoionized by ambient UV radiation from newly formed hot stars.)"
+ These fine-structure transitions can also be excited by electrons. once the electron fraction. wv). rises above about LO°.," These fine-structure transitions can also be excited by electrons, once the electron fraction, $x_e$ rises above about $10^{-3}$."
+" At redshifts. +20. the residual electron Yaction is,z10Danο aud the electrons contributed by ootoionization of trace ietals (Ch. Si! S. Fel) are jcelieible."," At redshifts, $z < 20$, the residual electron fraction is $x_e \approx 10^{-3.3}$, and the electrons contributed by photoionization of trace metals $^+$, $^+$, $^+$, $^+$ ) are negligible."
+ As the cloud cools and condeuses. ος declines owing to recombination ofthe trace IT!.," As the cloud cools and condenses, $x_e$ declines owing to recombination of the trace $^+$."
+" Iu Figures d. and 2. we plot the four metal cooling ""unctious together with £go at metallicities of 1027. and 10.1Z.. respectively."," In Figures \ref{fig:coolhighZ} and \ref{fig:coollowZ} we plot the four metal cooling functions together with ${\cal L}_{\rm H2}$ at metallicities of $10^{-2} Z_{\odot}$ and $10^{-4} Z_{\odot}$, respectively."
+ In both figures. we show the otal cooling function. as well as individual rates frouIL»...ΤΠ.I].. H].. aud TIJ.," In both figures, we show the total cooling function, as well as individual rates from, and ."
+. The low-deusity cooling rates are for gas with total number density me=107ocm7., The low-density cooling rates are for gas with total number density $n_t = 10^{-2}~{\rm cm}^{-3}$.
+" At high densities (iy,9105 3) and low imetallicities (Z10? Z.) cooling bx IE excitation of bbecoimes important at F<200/—LOO I& compared to the metal lines."," At high densities $n_H \geq 10^3$ $^{-3}$ ) and low metallicities $Z \leq 10^{-2}~Z_{\odot}$ ), cooling by $^{\circ}$ excitation of becomes important at $T < 200-400$ K compared to the metal lines."
+ The criterion for fragmentation. used to derive Zerit. conies from equating leating aud cooling functions of he form. Tyg=oT? and £i.T).," The criterion for fragmentation, used to derive $Z_{\rm crit}$, comes from equating heating and cooling functions of the form, $\Gamma_{\rm ad} = a T n^{3/2}$ and ${\cal L}(n,T)$."
+ Frou equation (5). we have à=2.7IkGp). where p=1.323.," From equation (5), we have $a = 2.74 k (G \mu)^{1/2}$, where $\mu = 1.323 m_H$."
+ Above he critical density. when the fiuc-tructure levels reach LTE (Boltzmann) populations. the LTE cooling functiou can be written £— nZf(T). where f(T) can be derived your the LTE excited populatious times the radiative ransitiou rates CÀoi) and line energies (£54).," Above the critical density, when the fine-structure levels reach LTE (Boltzmann) populations, the LTE cooling function can be written ${\cal L} = n Z f(T)$ , where $f(T)$ can be derived from the LTE excited populations times the radiative transition rates $A_{21}$ ) and line energies $E_{21}$ )."
+" Thus. the critical metallicity is The LTE cooling function. f(T) can casily be evaluated for a two-level atom of clemeut (4) of solar abundance ly; relative to hydrogen: where g4 and go are the statistical weights of the lower and upper states. respectively, of the two-level transition."," Thus, the critical metallicity is The LTE cooling function, $f(T)$, can easily be evaluated for a two-level atom of element (i) of solar abundance $A_i$ relative to hydrogen: where $g_1$ and $g_2$ are the statistical weights of the lower and upper states, respectively, of the two-level transition."
+ From the above.we arrive at a convenieut analytic expression for the critical metallicity for à eiven atom and fine-structure coolant. that depeuds only ou its reference abundance and theatomic parameters of the transition:," From the above,we arrive at a convenient analytic expression for the critical metallicity for a given atom and fine-structure coolant, that depends only on its reference abundance and theatomic parameters of the transition:"
+rave a close companion at a projected distance of about 10 spe.,have a close companion at a projected distance of about 10 kpc.
+" This result sugeests the preseuce of niergiug at 2« 1,5. which could increase the star formation. explainiug he blue colors of the sources."," This result suggests the presence of merging at $z<0.5$ , which could increase the star formation, explaining the blue colors of the sources."
+ Discutaneling the effects of luuinosity and umber density evolution is fundamental ia understanding the rature of field galaxy formation aud evolution. allowing a direct comparison with models," Disentangling the effects of luminosity and number density evolution is fundamental in understanding the nature of field galaxy formation and evolution, allowing a direct comparison with models."
+ For iustauce. soeni-analytical models of galaxy formation i hicrarchical clustering scenarios. assunmdue that ealaxics assenible through successive merecr events from smaller sub-uuits. make concrete predictions about the evolution of merger rate. morphological mix aud relative redshift distribution (c.g. Daugh et al.," For instance, semi-analytical models of galaxy formation in hierarchical clustering scenarios, assuming that galaxies assemble through successive merger events from smaller sub-units, make concrete predictions about the evolution of merger rate, morphological mix and relative redshift distribution (e.g. Baugh et al."
+ 1996: Ikauftinaun. 1996: I&auffuiuuimn Charlot 1998)., 1996; Kauffmann 1996; Kauffmann Charlot 1998).
+ In this scenario he nunber density of ealaxies should increase wi1 redshift. their sizes should decrease and the fraction of high redshift ealaxies (7> 2) in an IB-selected sample should be very low.," In this scenario the number density of galaxies should increase with redshift, their sizes should decrease and the fraction of high redshift galaxies $z>2$ ) in an IR-selected sample should be very low."
+ Moreover the morphological nix should change iu favour of late type aud mregular galaxies which thus would dominate both galaxy counts at faint magnitudes and the high redshüft tail of the redshift distribution. where a lack of elliptical galaxies should be observed.," Moreover the morphological mix should change in favour of late type and irregular galaxies which thus would dominate both galaxy counts at faint magnitudes and the high redshift tail of the redshift distribution, where a lack of elliptical galaxies should be observed."
+ Most of these predictions have heen the olajec of detailed: analvsis of data resulting im conflictiue results;, Most of these predictions have been the object of detailed analysis of data resulting in conflicting results.
+ For instance. sone works clau a deficit of ellipticals at 2>1 (Zepf 1997: Franceschini et al.," For instance, some works claim a deficit of ellipticals at $z>1$ (Zepf 1997; Franceschini et al."
+ 1998: Barecr et al., 1998; Barger et al.
+ 1999) aud a very low fraction ( 5%) of :2 galaxies ina K«21 selected suuple (Foutaua et al., 1999) and a very low fraction $\sim5\%$ ) of $z>2$ galaxies in a $<21$ selected sample (Fontana et al.
+ 1999)., 1999).
+ Ou the coutrary other authors find a constaut comoving density of cllipticals down to :o~2 (es. Beuittez et al., On the contrary other authors find a constant comoving density of ellipticals down to $z\sim2$ (e.g. tez et al.
+ 1999: Broadhurst Bowens 2000) and a strong clustering «igual of EROs in R-selected samples (Daddi ct al., 1999; Broadhurst Bowens 2000) and a strong clustering signal of EROs in K-selected samples (Daddi et al.
+ 2000)., 2000).
+ IIST. Meciuu Deep Survey (AIDS). ou he other laud. mdicates that a laree fraction of galaxies at amocerate redshift (2~ 0.5) are actually intrinsically faint aud exhibit peculiar morphologics sugeestive of 1uicreimeg (€Caitfiths et al.," HST Medium Deep Survey (MDS), on the other hand, indicates that a large fraction of galaxies at moderate redshift $z\sim0.5$ ) are actually intrinsically faint and exhibit peculiar morphologies suggestive of merging (Griffiths et al."
+ 1991: Driver et al., 1994; Driver et al.
+ 1995: Clazebrook et al., 1995; Glazebrook et al.
+ 1995: Neuschacter et al., 1995; Neuschaefer et al.
+ 1997)., 1997).
+ The results derived frou the morphological analvsis ou the ITIDE-N data (van den Dereh et al., The results derived from the morphological analysis on the HDF-N data (van den Bergh et al.
+ 1996: Abraham et al., 1996; Abraham et al.
+ 1996: Driver et al., 1996; Driver et al.
+ 1998) sccm Oo confuiu the Heh fraction of apparent nregular/peculiar galaxies with respect to the local universe which may be the cause of the faint blue galaxies excess (Ellis 1997)., 1998) seem to confirm the high fraction of apparent irregular/peculiar galaxies with respect to the local universe which may be the cause of the faint blue galaxies excess (Ellis 1997).
+ Williams et al. (, Williams et al. (
+1996) stress also that IIDE-N ealaxics show a strong bluiug color trend: while red galaxies are compatible with an elliptical galaxy. colors. most of blue sources are bluer than an Bregular galaxy. aud these colors night be accounted for ouly cousideriug Iuuinosity evolution.,"1996) stress also that HDF-N galaxies show a strong bluing color trend: while red galaxies are compatible with an elliptical galaxy colors, most of blue sources are bluer than an irregular galaxy, and these colors might be accounted for only considering luminosity evolution."
+ Thus it seems that the overall evolutionary scenario is still missed and the uuderstaudiug of the contribution of the different populations of galaxies to this scenario rather confused., Thus it seems that the overall evolutionary scenario is still missed and the understanding of the contribution of the different populations of galaxies to this scenario rather confused.
+ Iu this work we preseut an analysis of nuuber counts. colors aud morphological distributions of IIDE-S ealaxics in an attempt to reveal aud better coustraiut the evolving populations ofgalaxies.," In this work we present an analysis of number counts, colors and morphological distributions of HDF-S galaxies in an attempt to reveal and better constraint the evolving populations of galaxies."
+ The plan of the paper is as follows: iu Section 2 we sunmnaanze the steps leading to the creation of the catalogue. the computation of optical galaxy counts is explained in Section 3. iu Section [| the technique used to determine ealaxv colors.," The plan of the paper is as follows: in Section 2 we summarize the steps leading to the creation of the catalogue, the computation of optical galaxy counts is explained in Section 3, in Section 4 the technique used to determine galaxy colors."
+ In Section 5 we discuss our results. both counts aud colors. while in Section 6 we present aud discuss morphological imiuber counts and im Section 7. by adding the iuformation given by photometric redshifts and nuuber counts models. we cdiscutanele the contribution of different galaxw populations aud discuss the interpretation of redshift distributious.," In Section 5 we discuss our results, both counts and colors, while in Section 6 we present and discuss morphological number counts and in Section 7, by adding the information given by photometric redshifts and number counts models, we disentangle the contribution of different galaxy populations and discuss the interpretation of redshift distributions."
+ Finally. iu Section N owe summarize our results and ONpLress οἱ) conclusions.," Finally, in Section 8 we summarize our results and express our conclusions."
+ The samples here used are not extracted from the public TIDF-S catalogs available on the Web., The samples here used are not extracted from the public HDF-S catalogs available on the Web.
+ On the coutrary they have been independently created beeline from the optical UST images in order to have a coutrol of the etection procedure and of magnitude estimate., On the contrary they have been independently created beginning from the optical HST images in order to have a control of the detection procedure and of magnitude estimate.
+ A detailed description of the procedure adopted aud the selection criteria used to obtain the catalogs in the different bands as well as magnitude estimates. cau be found in Volouteri et al. (," A detailed description of the procedure adopted and the selection criteria used to obtain the catalogs in the different bands, as well as magnitude estimates, can be found in Volonteri et al. ("
+2000).,2000).
+ Tere we briefly describe the steps related to the derivation of number counts., Here we briefly describe the steps related to the derivation of number counts.
+ The object detection was pertormed across the area observed in optical bands (WFPC2 field) using SExtractor (Bertin and Arnouts 1996)., The object detection was performed across the area observed in optical bands (WFPC2 field) using SExtractor (Bertin and Arnouts 1996).
+ Tages were first 1100thied with a Caussian function having FWOAL equal to the one measured on the images (27 0.2 arcsec)., Images were first smoothed with a Gaussian function having FWHM equal to the one measured on the images $\approx$ 0.2 arcsec).
+ A detection threshold of 1 sigma per pixel aud a minima detection area equal to the secing disk (m0.02 aresec?=19 pixel) were adopted to peal up objects., A detection threshold of 1 sigma per pixel and a minimum detection area equal to the seeing disk $\approx 0.02$ $^2 = 13$ pixel) were adopted to peak up objects.
+" This +ireshold. corresponds to a limi signal-to-noise ratio of SYMpr=131 and ορε=0.67 for the fuutest sources detectable on the WF area aud on the PC area respectively,", This threshold corresponds to a minimum signal-to-noise ratio of $S/N_{WF}=1.34$ and $S/N_{PC}=0.67$ for the faintest sources detectable on the WF area and on the PC area respectively.
+ Iun Table 1 we report for each filter the zero-point (AB inagnuitude. Oke 1971). the sky RAIS estimated by SExtractor aud the corresponding 5o limiting maguitude for a point source.," In Table 1 we report for each filter the zero-point (AB magnitude, Oke 1974), the sky RMS estimated by SExtractor and the corresponding $\sigma$ limiting magnitude for a point source."
+The slobal heating function which drives the flare evolution of Fig.,The global heating function which drives the flare evolution of Fig.
+ 6 is shown in Fig. 9.., \ref{fig:bs6} is shown in Fig. \ref{fig:heat}.
+" Given the close sinulavity of the heating function of the two flaring structures. mdependeut although probably adjaccut. we navy advance the hypothesis that the two heating components may often be roth esent iu inauy flares, aud nay represent a general characteristics of solar aud stellar Hares."," Given the close similarity of the heating function of the two flaring structures, independent although probably adjacent, we may advance the hypothesis that the two heating components may often be both present in many flares, and may represent a general characteristics of solar and stellar flares."
+ The idea that one reating mechanisni is probably iusufücieut to explain coroual flares is not new (c.g. Peres et al., The idea that one heating mechanism is probably insufficient to explain coronal flares is not new (e.g. Peres et al.
+ 1987. Masuda et al.," 1987, Masuda et al."
+" 190, but our analysis may xovide a detailed aud quantitative patter o be explored in other events."," 1994), but our analysis may provide a detailed and quantitative pattern to be explored in other events."
+ Note. iu particular. that the heating at the ootpoiuts may be driveu x ligh-enereyv electron beans oxecipitatiue along the loop from a reconnection site hieli in the loop. as often meutioned in the literature (o.g. Masuda et al.," Note, in particular, that the heating at the footpoints may be driven by high-energy electron beams precipitating along the loop from a reconnection site high in the loop, as often mentioned in the literature (e.g. Masuda et al."
+ 1991). aud as also traced by the optical ight curve (Paper IT).," 1994), and as also traced by the optical light curve (Paper II)."
+" On the basis of the inodel which best describes the N-rayv ata on this flare, we can male some considerations about ous enerev budect."," On the basis of the model which best describes the X-ray data on this flare, we can make some considerations about its energy budget."
+" Consideriug the loop aspect Ry/L—V1 for loop À aud Rp/L~0.22 for loop DB. we obtain that the αλα heating rates injected im the main oop and oe ithe arcade are 2οσον27«1077 eres; aud z1LL.DS1077 ergs.; respectively,"," Considering the loop aspect $R_A/L \approx 0.1$ for loop A and $R_B/L
+\approx 0.22$ for loop B, we obtain that the maximum heating rates injected in the main loop and in the arcade are $\approx 27 \times
+10^{28}$ erg/s and $\approx 14 \times 10^{28}$ erg/s, respectively."
+ For comparison. the maxima Imunuinositv of the flare las been estimated⋅ to be Lyyasiw223.9.«MEN1077 eres (Paper ID. 1e. ~15% of the maxinuuu heating rate.," For comparison, the maximum luminosity of the flare has been estimated to be $L_{X, 0.15-10} \approx 3.9 \times 10^{28}$ erg/s (Paper II), i.e. $\sim 15$ of the maximum heating rate."
+ The rest of the heating rate goes iuto enthalpy. conduction of thermal energy dowmwards to the chromosphere. aud radiation at lower energies.," The rest of the heating rate goes into enthalpy, conduction of thermal energy downwards to the chromosphere, and radiation at lower energies."
+ The rate of the unifoiui heating at the beeline of he decays ave z27.2«1075 cre/s aud zzL8<107 org/s. respectively.," The rate of the uniform heating at the beginning of the decays are $\approx 7.2 \times 10^{28}$ erg/s and $\approx 1.8 \times 10^{28}$ erg/s, respectively."
+ If we integrate iu time. we obtain that the otal enerev released over the decay plase (3.2«107? cre or loop A and 0.8&107? erg for loop D) is for loop A wice. and for loop B equal to. the energy released im the jcating pulse.," If we integrate in time, we obtain that the total energy released over the decay phase $3.2 \times 10^{32}$ erg for loop A and $0.8 \times 10^{32}$ erg for loop B) is for loop A twice, and for loop B equal to, the energy released in the heating pulse."
+ Sununiug over the whole fare we obtain a total fare hermal enerey input o6.5 «1077 ere. to be compared to a total energy. racliatec in the 015-10 keV of 1.5«107? erg (Paper D. ie. —25% of the total injected energy.," Summing over the whole flare we obtain a total flare thermal energy input of $6.5 \times 10^{32}$ erg, to be compared to a total energy radiated in the 0.15-10 keV of $1.5 \times 10^{32}$ erg (Paper I), i.e. $\sim 25$ of the total injected energy."
+" For comparison. the total euergev. (of the analyzed part) and the peak N-rav numositv of the flare observe on Proxima Centauri with the satellite were ο,ς10% ore and τή1.2«1053 cre/s. respectively. ic. about 1/3 and 1/l of he total enerevOo, and peak rate of the heating used in the relevaut lyvdrodvuamic modeling (Reale et al."," For comparison, the total energy (of the analyzed part) and the peak X-ray luminosity of the flare observed on Proxima Centauri with the satellite were $\approx 2 \times 10^{31}$ erg and $\approx 1.2 \times
+10^{28}$ erg/s, respectively, i.e. about 1/3 and 1/4 of the total energy and peak rate of the heating used in the relevant hydrodynamic modeling (Reale et al."
+ 1988)., 1988).
+ Tha flare appeared to be less energetic aud therefore. consistently. softer aud with less efficieu thermal conduction than the flare aualvzed here.," That flare appeared to be less energetic and therefore, consistently, softer and with less efficient thermal conduction than the flare analyzed here."
+ For comparison with a solar flare. the total thermal energv involved iu the Bastille day flare has heen estimated to be ~0.5&]077- ere (Asclawanden Alexander Re]20013. about 1/10 of the euerev iuput of this Prox Cen flare.," For comparison with a solar flare, the total thermal energy involved in the Bastille day flare has been estimated to be $\sim 0.5 \times
+10^{32}$ erg (Aschwanden Alexander 2001), about 1/10 of the energy input of this Prox Cen flare."
+ The bes fittingzl model provides us with au insighto on the evolution of the plasina involved iu the flare and confined in the faring structures., The best fitting model provides us with an insight on the evolution of the plasma involved in the flare and confined in the flaring structures.
+ Some features are shown in Fig., Some features are shown in Fig.
+ 10 (see also Table 1)., \ref{fig:profs} (see also Table 1).
+ The most dvuanüc aud luteuse ewution of the flare occurs iu loop A: the plasma teniperatiEὉ rises in few seconds from about 3 MIS to 30 AQ. and to an absolute maxim close to 50 MIS. after the first 100 s: if then settles at about LO AUS during the time the heat pulse is on.," The most dynamic and intense evolution of the flare occurs in loop A: the plasma temperature rises in few seconds from about 3 MK to 30 MK, and to an absolute maximum close to 50 MK after the first 100 s; it then settles at about 40 MK during the time the heat pulse is on."
+ After the first 10 s plasma begius to evaporate significantly upwards frou he chromosphere with velocities above 1000 Em/s aud making the coronal chsity increase by almost two orders of magnitude to out 1029 ?.," After the first 10 s plasma begins to evaporate significantly upwards from the chromosphere, with velocities above 1000 km/s and making the coronal density increase by almost two orders of magnitude to about $10^{10}$ $^{-3}$."
+ In spite of the high velocity. the relevant Doppler (blue) shifts would be dificult to be detecte evel in the case of a favourable loop oricutation to the line sieht. because restricted to very hot lines (>10 MEK) undetected with the RCS. aud a very short time interval (the initial 2-3 mun) of the flare. as typically occurs iu solar flares (c.g. Autonucci ct al.," In spite of the high velocity, the relevant Doppler (blue) shifts would be difficult to be detected even in the case of a favourable loop orientation to the line of sight, because restricted to very hot lines $> 10$ MK) undetected with the RGS, and a very short time interval (the initial 2-3 min) of the flare, as typically occurs in solar flares (e.g. Antonucci et al."
+ 1987)., 1987).
+ The bluc-shiftea liue component is herefore lost. because highly diluted iu the typical time intervals of photon integration LO win (Paper I).," The blue-shifted line component is therefore lost, because highly diluted in the typical time intervals of photon integration $\ga 10$ min (Paper I)."
+ After ~1 uiu he strong evaporation front reaches tli top of the loop., After $\sim 1$ min the strong evaporation front reaches the top of the loop.
+ Since then. the plasma coutiuues to fill up," Since then, the plasma continues to fill up"
+to the data is represented in red.,to the data is represented in red.
+ The canonical synchrotron component is shown on light blue and the canonical dust in orange., The canonical synchrotron component is shown on light blue and the canonical dust in orange.
+ The extra low-energy cutoff synchrotron component is represented in blue., The extra low-energy cutoff synchrotron component is represented in blue.
+" As above, we obtain a good global fit to the data as shown by the x?/Naof value."," As above, we obtain a good global fit to the data as shown by the $\chi^2 / N_{dof}$ value."
+ The best-fit parameters obtained for the canonical synchrotron and dust components are in good agreement with those presented in Sect., The best-fit parameters obtained for the canonical synchrotron and dust components are in good agreement with those presented in Sect.
+ 3 and on (Strom&Greidanus1992)., \ref{sec:EMspectrum} and on \cite{Strom1992}.
+. The amplitude of the extra synchrotron component is compatible with zero as well as the value of the spectral index of the electron energy distribution., The amplitude of the extra synchrotron component is compatible with zero as well as the value of the spectral index of the electron energy distribution.
+ We therefore conclude that there is no evidence of an extra component in the form of low-energy cutoff synchrotron., We therefore conclude that there is no evidence of an extra component in the form of low-energy cutoff synchrotron.
+" In addition, we observe that the x?/Naof in the millimetric region from 100 to 1000 GHz is worse than the one obtained in Sect 3.3 assuming the canonical model only."," In addition, we observe that the $\chi^2/ N_{dof}$ in the millimetric region from 100 to 1000 GHz is worse than the one obtained in Sect \ref{sect:mmexcess} assuming the canonical model only."
+ The right panel of figure 6 represents the residuals to the best-fit model to the data on the millimetric regime., The right panel of figure \ref{fig:lowsyncfit} represents the residuals to the best-fit model to the data on the millimetric regime.
+ From this we conclude that the fit to the data in this frequency regime is not improved by adding an extra synchrotron component., From this we conclude that the fit to the data in this frequency regime is not improved by adding an extra synchrotron component.
+" We have also performed the analysis of the data set considering the p and v, parameters fixed and set to the values quoted by (Bandieraetal.2002).", We have also performed the analysis of the data set considering the $p$ and $\nu_{c}$ parameters fixed and set to the values quoted by \cite{Bandiera2002}.
+. 'The analysis shows that the amplitude of the extra synchrotron component is compatible with zero and therefore we conclude that the data show no evidence of an extra synchrotron component., The analysis shows that the amplitude of the extra synchrotron component is compatible with zero and therefore we conclude that the data show no evidence of an extra synchrotron component.
+" From the previous section, we conclude that the intensity emission of the Crab nebula at millimetre and submillimetre wavelengths is dominated by the well known synchrotron radiation observed at radio wavelengths."," From the previous section, we conclude that the intensity emission of the Crab nebula at millimetre and submillimetre wavelengths is dominated by the well known synchrotron radiation observed at radio wavelengths."
+ The data show no evidence of an extra synchrotron component and evidence for an extra dust component is not significant., The data show no evidence of an extra synchrotron component and evidence for an extra dust component is not significant.
+" Therefore, we can postulate that the Crab nebula emission at millimetre and submillimetre wavelengths is produced by the same relativistic electrons that"," Therefore, we can postulate that the Crab nebula emission at millimetre and submillimetre wavelengths is produced by the same relativistic electrons that"
+If the long-lived SO model. is. correct. then the QSO clustering results may be consistent with the fit of PILIS models to the QSO luminosity function.,"If the long-lived QSO model is correct, then the QSO clustering results may be consistent with the fit of PLE models to the QSO luminosity function."
+ A physical interpretation of why the QSO space density at L is observed. to remain approximately constant [rom z=2 Oo 2= Vis that the QSOs have a long lifetime anc the uminous Tvpe LQ8Os seen at 2=2.2 dim by a factor of zz30 to become the low-luminosity Seyfert Fs seen at z=0., A physical interpretation of why the QSO space density at $L^*$ is observed to remain approximately constant from $z=2$ to $z=0$ is that the QSOs have a long lifetime and the luminous Type I QSOs seen at $z=2.2$ dim by a factor of $\approx30$ to become the low-luminosity Seyfert I's seen at z=0.
+ Thus this PLE model prediction may be supported w the QSO clustering results., Thus this PLE model prediction may be supported by the QSO clustering results.
+ We have noted that if Alea and Alparg Follow relations such as MpggoxAlb. then he/heorcliceal erowth rates of Mgg [rom PLE and the one-lived mocel are mutually incompatible.," We have noted that if $M_{BH}$ and $M_{DMH}$ follow relations such as $M_{BH}\propto M_{DMH}^{1.8}$, then the growth rates of $M_{BH}$ from PLE and the long-lived model are mutually incompatible."
+ However. a kev eature of the long-lived model is that it implicitly. assumes hat ealaxy/BIL and heo mass growth rates are decouplec so that the DII mass growth rate with z predicted by PLE is easily aecommocdated.," However, a key feature of the long-lived model is that it implicitly assumes that galaxy/BH and halo mass growth rates are decoupled so that the BH mass growth rate with $z$ predicted by PLE is easily accommodated."
+ At fixed ο. PLE also predicts that BL mass growth rates are proportional to QSO luminosity.," At fixed $z$, PLE also predicts that BH mass growth rates are proportional to QSO luminosity."
+ Bu again. since in the context of long-lived models. halo and BL mass growth are naturally decoupled there is no necessary contradiction with the clustering results here for PLE.," But again, since in the context of long-lived models, halo and BH mass growth are naturally decoupled there is no necessary contradiction with the clustering results here for PLE."
+ This decoupling has given. us the freedom to assume a constan initial DII mass for the PLE model so that it is actually in line with the luminosity independent clustering results a high z., This decoupling has given us the freedom to assume a constant initial BH mass for the PLE model so that it is actually in line with the luminosity independent clustering results at high $z$.
+ Phen the higher accretion rate of the more luminous QSOs gives the possibility of generating the Iuminosity-DI reverberation mass correlation found at 2=0. without also generating clustering that is too luminosity dependent. a these redshifts.," Then the higher accretion rate of the more luminous QSOs gives the possibility of generating the luminosity-BH reverberation mass correlation found at $z=0$, without also generating clustering that is too luminosity dependent at these redshifts."
+ In this picture. it is the lack of à relation between halo ancl galaxy/BILE mass that explains the lack of Luminosity dependence of QSO clustering seen in Fig. 1... ," In this picture, it is the lack of a relation between halo and galaxy/BH mass that explains the lack of luminosity dependence of QSO clustering seen in Fig. \ref{fig:xis}, ,"
+rather than the lack of a relation between luminosity ancl BI mass., rather than the lack of a relation between luminosity and BH mass.
+ ltecentlv. Kormencdy.Dender.&Cornell(2011):Ixo-rmendy&Bender(2011). have presented. evidence that nuclear velocity dispersions may not. correlate well with ealaxy circular velocities in cases they have observed. where there is no bulge. just nuclear star-clusters.," Recently, \citet{korm11a,korm11b} have presented evidence that nuclear velocity dispersions may not correlate well with galaxy circular velocities in cases they have observed where there is no bulge, just nuclear star-clusters."
+ They also argue that dynamically measured Alexw do not. correlate well with disc stellar masses., They also argue that dynamically measured $M_{BH}$ do not correlate well with disc stellar masses.
+ Since circular velocities and. disk stellar masses usually correlate well with DM halo masses. they then conclude that DAIL masses don't correlate well with Adeaw.," Since circular velocities and disk stellar masses usually correlate well with DM halo masses, they then conclude that DMH masses don't correlate well with $M_{BH}$."
+ These results have some resonance with the arguments we have presented above based on QSO clustering and also the direct imaging of QSO hosts., These results have some resonance with the arguments we have presented above based on QSO clustering and also the direct imaging of QSO hosts.
+ We have argued hat these results imply either that QSOs have the same DII mass or that the halo and DII mass growth rates are decoupled., We have argued that these results imply either that QSOs have the same BH mass or that the halo and BH mass growth rates are decoupled.
+ La particular. our long-lived PLE model argues or à decoupling of halo and. DII mass at least at 2<2.2 to reconcile the slow evolution of QSO clustering with the faster oreclicted evolution of halo clustering.," In particular, our long-lived PLE model argues for a decoupling of halo and BH mass at least at $z<2.2$ to reconcile the slow evolution of QSO clustering with the faster predicted evolution of halo clustering."
+ In the models of Lidzetal.(2006) it is the QSO luminosity that decouples from he BIL/DALLE masses to explain the clustering results., In the models of \citet{lidz06} it is the QSO luminosity that decouples from the BH/DMH masses to explain the clustering results.
+ But in these scenarios the 11 and DMLU masses are expected to »e strongly. correlated. in contradiction with the new results of Ixormendsy.Bender.&Cornell(2011)..," But in these scenarios the BH and DMH masses are expected to be strongly correlated, in contradiction with the new results of \citet{korm11a}."
+ For clarity. we summarise the arguments for and against the four QSO mocels we have discussed above: lt is clear that. models which predict. well the QSO0 clustering results then tend to have a problem: with the reverberation and UV. bump AlewL relation and vice versa.," For clarity, we summarise the arguments for and against the four QSO models we have discussed above: It is clear that models which predict well the QSO clustering results then tend to have a problem with the reverberation and UV bump $M_{BH}-L$ relation and vice versa."
+" We take the view that it is easier to circumvent the clustering results (by decoupling Mgg and AM4,) than the reverberation mapping results.", We take the view that it is easier to circumvent the clustering results (by decoupling $M_{BH}$ and $M_{halo}$ ) than the reverberation mapping results.
+ Therefore the answer to the question posed in the title may be that QSOs do not all have the same black-hole mass. although they may occupy similar halo masses at fixed redshift.," Therefore the answer to the question posed in the title may be that QSOs do not all have the same black-hole mass, although they may occupy similar halo masses at fixed redshift."
+ Overall. we conclude that a model that is reasonably consistent with both clustering anc the reverberation/UV-bunmp masses is the long-lived PLE model.," Overall, we conclude that a model that is reasonably consistent with both clustering and the reverberation/UV-bump masses is the long-lived PLE model."
+ We thank an anonymous referee. for comments which improved the quality of this paper., We thank an anonymous referee for comments which improved the quality of this paper.
+ US acknowledges financial supportfrom the Institute, US acknowledges financial supportfrom the Institute
+"We have shown it is possible to derive the 3D kinematics of features, the CME apex in this case, in the inner Heliosphere (~2—-250RRo) using STEREO observations.","We have shown it is possible to derive the 3D kinematics of features, the CME apex in this case, in the inner Heliosphere $\sim$ $_{\odot}$ ) using STEREO observations."
+" The 3D kinematics are free from the projection effects of traditional 2D kinematics but may contain artifact from the 3D reconstruction method (e.g, Maloneyetal. 2009)) and other sources."," The 3D kinematics are free from the projection effects of traditional 2D kinematics but may contain artifact from the 3D reconstruction method (e.g, \citealt{Maloney:2009p6617}) ) and other sources."
+" Both of the accelerating events showed two regimes in the velocity profile, a low down (« RR) early rapid acceleration (in comparison to later values), followed by a gradual acceleration far from the Sun (7330 RR)."," Both of the accelerating events showed two regimes in the velocity profile, a low down $<$ $_{\odot}$ ) early rapid acceleration (in comparison to later values), followed by a gradual acceleration far from the Sun $>$ $_{\odot}$ )."
+" The early acceleration is thought to be due to a magnetic driving force, as the solar wind velocity low in the corona"," The early acceleration is thought to be due to a magnetic driving force, as the solar wind velocity low in the corona"
+delata on 19 ancl 16 dillerent svstems. respectively.,"data on 19 and 16 different systems, respectively."
+ In order to take account of zero-point dilferences reported by Dressler (1984). the 7S LCOLIL data have been subedivided into the three constituent runs from which they derive.," In order to take account of zero-point differences reported by Dressler (1984), the 7S LCOHI data have been subdivided into the three constituent runs from which they derive."
+ Alany galaxies have measurements on more than (wo systems., Many galaxies have measurements on more than two systems.
+ Therefore in. order to determine seli-consistent corrections between dillerent svstenms. a sipultaneous fit for all of the olfsets is necessary.," Therefore in order to determine self-consistent corrections between different systems, a simultaneous fit for all of the offsets is necessary."
+ The fit is performed. using velocity dispersion ancl cdata corrected to the Jorgensen et al. (, The fit is performed using velocity dispersion and data corrected to the rgensen et al. (
+1995b). standard physical aperture size of 1.19 h.1 kpe.,1995b) standard physical aperture size of 1.19 $h^{-1}$ kpc.
+ We determine the corrections necessary to bring all svstems into the best possible agreement. with cach other., We determine the corrections necessary to bring all systems into the best possible agreement with each other.
+ We adopt. the fullv-corrected Lick svstem (Davies et 11987) as the stancare and determine the remaining corrections as follows., We adopt the fully-corrected Lick system (Davies et 1987) as the standard and determine the remaining corrections as follows.
+" Let s=log,,(o) and let #7 and & index the nicasurenient. galaxy and system: respectively."," Let $s = \log_{10}(\sigma)$ and let $i$ $j$ and $k$ index the measurement, galaxy and system respectively."
+ We obtain the corrections Aj. needed to bring each svstem intoagreement with Lick. bv minimising a X7 statistic where ορ is the error in s; (assumed to be the same for all galaxies in a given system) and 8; is the errorweighted mean of all corrected. measurements of the same galaxy.," We obtain the corrections $\Delta_k$, needed to bring each system intoagreement with Lick, by minimising a $\chi^2$ statistic where $e_k$ is the error in $s_i$ (assumed to be the same for all galaxies in a given system) and $\overline{s}_j$ is the error-weighted mean of all corrected measurements of the same galaxy."
+ We determine the errors ej for cach svstem by adjusting these so that the reduced x7 is unity. both when the system is included. and. when it is excluded. from the comparisons.," We determine the errors $e_k$ for each system by adjusting these so that the reduced $\chi^2$ is unity, both when the system is included and when it is excluded from the comparisons."
+ This external error. (654) is twpically 10 25 per cent larger than the internal error (6) estimated from repeat measurements on the same system., This external error $e_{\rm ext}$ ) is typically 10 – 25 per cent larger than the internal error $e_{\rm int}$ ) estimated from repeat measurements on the same system.
+ The overlap data set of velocity dispersion measurcments (galaxies with velocity dispersions on more than one svslenm) consists of 1281 measurements for 350 cdillerent galaxies., The overlap data set of velocity dispersion measurements (galaxies with velocity dispersions on more than one system) consists of 1281 measurements for 350 different galaxies.
+ We exclude galaxies with s<2 as these may. be subject to large random. and. systematic errors (Jorgensen et al., We exclude galaxies with $\overline{s} < 2$ as these may be subject to large random and systematic errors rgensen et al.
+ 1995ο)., 1995b).
+ We also exclude incivicual velocity dispersion measurenrents which are inconsistent at the 3.5¢ level with the other measurecl velocity dispersions of the same galaxy., We also exclude individual velocity dispersion measurements which are inconsistent at the $3.5 \sigma$ level with the other measured velocity dispersions of the same galaxy.
+ The veocitv dispersions so excluded are X2634-E1201. ονο s=2.0784). X1656D-136 (CLNTE90 5=2.0888). N386 (IXPNO 5=1.7923). N548 (LICI 5=15500) and VELA-G22 (FOCTP2 «=1.9237).," The velocity dispersions so excluded are A2634-F1201 (EEV93 $s=2.0784$ ), A1656D-136 (INT90 $s=2.0888$ ), N386 (KPNO $s=1.7923$ ), N548 (LICK $s=1.8856$ ) and VELA-G22 (FOCP2 $s=1.9237$ )."
+ The overlap data set of numeasurements (galaxies with measurements on more than one svsteni) consists of 1013 measurements of 270 cilferent ealaxies on 16 svstems (the LC. FOCP2 and EIV94 systems have no ασια).," The overlap data set of measurements (galaxies with measurements on more than one system) consists of 1013 measurements of 270 different galaxies on 16 systems (the LC, FOCP2 and EEV94 systems have no data)."
+ In addition to the galaxies excluded in the velocity dispersion comparison. we also exclude the following data which are inconsistent with other measurements of the same ealaxy at the 3.50 level: NI282 (PAL == (0.0245). N1549 (AZ == 0.264) N4564 (ΕΟΝ93 350) and N6702 (GONZA 243 and TEI94 == 0.288).," In addition to the galaxies excluded in the velocity dispersion comparison, we also exclude the following data which are inconsistent with other measurements of the same galaxy at the $3.5 
+\sigma$ level: N1282 (PAL = 0.0245), N1549 (A2 = 0.342 and JFK = 0.264) N4564 (EEV93 = 0.350) and N6702 (GONZA = 0.243 and TEK94 = 0.288)."
+ ‘Tables G6 and 7 summarise the required corrections to velocity dispersion. ancAles... respectively.," Tables \ref{sigcorrs} and \ref{mgcorrs} summarise the required corrections to velocity dispersion and, respectively."
+ Note that. because of the interdepencdencies between. the. clillerent corrections. the simple pair olfsets of Table 5 are not trivially related to those clerivecl here by simultaneous fits.," Note that, because of the interdependencies between the different corrections, the simple pair offsets of Table \ref{specruntoruntab} are not trivially related to those derived here by simultaneous fits."
+ In Figures 6 and 7 νο illustrate the level to which systematic olfsets are removed. by. the application of the derived: corrections., In Figures \ref{sysmatsigs} and \ref{sysmatmgs} we illustrate the level to which systematic offsets are removed by the application of the derived corrections.
+ The errors. are. determined. by bootstrap resampling the master cata file ancl computing the corrections [rom the resampled file., The errors are determined by bootstrap resampling the master data file and computing the corrections from the resampled file.
+ This procedure allow us to determine not only the error on the correction to cach system but also the correlation between the corrections for cilferent systems., This procedure allow us to determine not only the error on the correction to each system but also the correlation between the corrections for different systems.
+ Using the bootstrap values of these corrections. we can generate mock merged data sets and. so determine for a given. cluster the error in the mean correction.," Using the bootstrap values of these corrections, we can generate mock merged data sets and so determine for a given cluster the error in the mean correction."
+ This is an estimate of the mean systematic error in s. which will generally depend. on the systems. merged for the eluster. their relative proportions and their covariance.," This is an estimate of the mean systematic error in $s$, which will generally depend on the systems merged for the cluster, their relative proportions and their covariance."
+ For the PP sample. we find that for all clusters this error is ~1.5 per cent in o.," For the PP sample, we find that for all clusters this error is $\sim$ 1.5 per cent in $\sigma$ ."
+ This translates to à systematic error of ~2 per cent in distance. or  aat PP.," This translates to a systematic error of $\sim$ 2 per cent in distance, or $\sim$ at PP."
+ In this section. we brielly summarise. the recipe for," In this section, we briefly summarise the recipe for"
+Interfaces between hot aud cold plasmas can occur in astroplvsics where understanding the rate of thermal conduction may be an important. part. of the the astroplivsical phenomenology.,Interfaces between hot and cold plasmas can occur in astrophysics where understanding the rate of thermal conduction may be an important part of the the astrophysical phenomenology.
+ One example occurs in wind blown bubbles (WDBD) of evolved stars where maegneltized hot supersonie outflow shock heats the cooler ambient magnetized interstellar medium., One example occurs in wind blown bubbles (WBB) of evolved stars where magnetized hot supersonic outflow shock heats the cooler ambient magnetized interstellar medium.
+ For such WDBD. there are examples where (he presumed shock heated: bubble is cooler than expected if only radiative cooling is considered (Zhekovetal. (2011))).," For such WBB, there are examples where the presumed shock heated bubble is cooler than expected if only radiative cooling is considered \citet{zhe11}) )."
+ A possible explanation is that heat loss through Che interface of hot bubble into the cold shell via thermal conduction reduces the temperature of the hot bubble, A possible explanation is that heat loss through the interface of hot bubble into the cold shell via thermal conduction reduces the temperature of the hot bubble
+Sco and a systemic mass loss through the outer Lagrangian poiuts.,Sco and a systemic mass loss through the outer Lagrangian points.
+ We also discuss the relation to a progenitor svsten of SNe Ta. Our U Sco model is eraphically shown in Figure 1.., We also discuss the relation to a progenitor system of SNe Ia. Our U Sco model is graphically shown in Figure \ref{uscofig_q35}.
+ Schaefer (1990) aud Schacter Biugwald (1995) observed eclipses of U Sco in the quiescent phase and detezuinced the orbital period (P=1.23056 davs) and the ephemeris(IILJD 2.151.235.777. | 1.2305G6E) at the epoch of mid-eclipse.," Schaefer (1990) and Schaefer Ringwald (1995) observed eclipses of U Sco in the quiescent phase and determined the orbital period $P= 1.23056$ days) and the ephemeris(HJD 2,451,235.777 $+$ $E$ ) at the epoch of mid-eclipse."
+ Thus. the companion is a 1iadn-sequenuce star (MS) which expands to fill its Roche lobe after a most part of the central hydrogen is coustuued.," Thus, the companion is a main-sequence star (MS) which expands to fill its Roche lobe after a most part of the central hydrogen is consumed."
+ We call such a star a slightly evolved” MS.," We call such a star ""a slightly evolved"" MS."
+ The inclination angle of the orbit (i~ SU?) ds a parameter for fitting., The inclination angle of the orbit $i \sim 80\arcdeg$ ) is a parameter for fitting.
+ We have asstuned that (1) AAyp=1.37M... (2) the WD huninosity of where the first term ds the accretion hDnmuuinositv (6.8... Stairfield. Sparks. Shaviv 1988) aud the second term [νου is the intrinsic huninositv of the WD. aud Ryp=0.0032R.. the radius of the 1.27A.. WD. aud (3) a black- photosphereof the WD.," We have assumed that (1) $M_{\rm WD}= 1.37 M_\odot$, (2) the WD luminosity of where the first term is the accretion luminosity (e.g., Starrfield, Sparks, Shaviv 1988) and the second term $L_{\rm WD,0}$ is the intrinsic luminosity of the WD, and $R_{\rm WD}= 0.0032 R_\odot$ the radius of the $1.37 M_\odot$ WD, and (3) a black-body photosphereof the WD."
+ The accretion luminosity is ~του. for a suggested mass accretion rate of HNe3.5410TAL d., The accretion luminosity is $\sim 1700 L_\odot$ for a suggested mass accretion rate of $\dot M_{\rm acc} \sim 2.5 \times 10^{-7} M_\odot$ $^{-1}$.
+ Tere. we assume Lap=0 because the nuclearbhuuinositv is smaller than the accretion Iuuuinositv for this accretion rate. but we have exanuned other two cases of νου=2000 and 1000 L.. and found no significant differences in the distance as shown below.," Here, we assume $L_{\rm WD,0}=0$ because the nuclearluminosity is smaller than the accretion luminosity for this accretion rate, but we have examined other two cases of $L_{\rm WD,0}=2000$ and 4000 $L_\odot$ and found no significant differences in the distance as shown below."
+ We do not consider the liamh-darkening effect for simplicity., We do not consider the limb-darkening effect for simplicity.
+ Tt is assumed that the companion star is svuclronously rotating on a circular orbit and its surface fills the ΠΟ critical Roche lobe as shown in Figure 1.., It is assumed that the companion star is synchronously rotating on a circular orbit and its surface fills the inner critical Roche lobe as shown in Figure \ref{uscofig_q35}.
+ We ucelect both thelimb-darkening effect and the eravity-darkenius effect of the companion star for simplicity., We neglect both thelimb-darkening effect and the gravity-darkening effect of the companion star for simplicity.
+ Here. we assume a irradiation efficiencv of the companion star Clips= 0.5).," Here, we assume a irradiation efficiency of the companion star $\eta_{\rm ir,MS}=0.5$ )."
+ We have exiunined the dependence of the distance on the radiation efficiency (16. fu=0.25 and 1.0) but fouud no sigui&cant differences iu the distanceas shown below.," We have examined the dependence of the distance on the irradiation efficiency (i.e., $\eta_{\rm ir,MS}=0.25$ and 1.0) but found no significant differences in the distanceas shown below."
+ The nou-nradiated plhotospheric teniperature Duais of the companion star is a parameter for fitting.," The non-irradiated photospheric temperature $T_{\rm ph, MS}$ of the companion star is a parameter for fitting."
+ The mass of the secondary is assumed to be Ais=LOA..., The mass of the secondary is assumed to be $M_{\rm MS}= 1.5 M_\odot$.
+ The size of the accretion disk is a parameter for fitting and defined as where a is a nunerieal factor iudicatiug the size of the aceretion disk. aud Hj the effective radiusof the inner critical Roche lobe for the WD component (c.g. Eeeleton 1983).," The size of the accretion disk is a parameter for fitting and defined as where $\alpha$ is a numerical factor indicating the size of the accretion disk, and $R_1^*$ the effective radiusof the inner critical Roche lobe for the WD component (e.g., Eggleton 1983)."
+ We also asstune that the accretion disk Is axisvnunetrie and has a thickness giveu by where / ds the height of the surface from the equatorial plane. z the distance ou the equatorial plane from the center of the WD. » the power of the surface shape. and Jj a numerical factor showing the degree of thickness aud also a paraiueter for fitting.," We also assume that the accretion disk is axisymmetric and has a thickness given by where $h$ is the height of the surface from the equatorial plane, $\varpi$ the distance on the equatorial plane from the center of the WD, $\nu$ the power of the surface shape, and $\beta$ a numerical factor showing the degree of thickness and also a parameter for fitting."
+" We adopt a ze-squared law (v=2) sinplv to miünüc the flarine-p effect of the accretion disk van (0,0), Schandl. Meyer-Iofineister. Meyer 1997). aud have exanuned the dependence of the distance on the power (p=1.25 and 3.0) without finding ασ siguificaut differences as shown below."," We adopt a $\varpi$ -squared law $\nu=2$ ) simply to mimic the flaring-up effect of the accretion disk rim (e.g., Schandl, Meyer-Hofmeister, Meyer 1997), and have examined the dependence of the distance on the power $\nu=1.25$ and 3.0) without finding any significant differences as shown below."
+" The surface of the accretion disk also absorbs photos frou, the WD photosphere and recnüts with a black-body spectrum at a local temperature.", The surface of the accretion disk also absorbs photons from the WD photosphere and reemits with a black-body spectrum at a local temperature.
+ We assume a irradiation efficiencv of the companion star. Lo. Hipk=0.5 (e.g. Schaucdletal.19973).," We assume a irradiation efficiency of the companion star, i.e., $\eta_{\rm ir,DK}=0.5$ (e.g., \cite{sch97}) )."
+ We have examined other two cases of ion=0.25 and 1.0. and fouud. no sienificant differences in the distance as shown below.," We have examined other two cases of $\eta_{\rm ir,DK}=0.25$ and 1.0, and found no significant differences in the distance as shown below."
+ The uon-iradiated temperature of the disk surface is assumed to be determined by the viscous heating of the standard accretion disk model., The non-irradiated temperature of the disk surface is assumed to be determined by the viscous heating of the standard accretion disk model.
+ Then. the disk surface temperature 1s Given by where rz the distance from the WD center. aud cos0 the incident angle of the surface (c.g.. Schaudletal.1997)).," Then, the disk surface temperature is given by where $r$ the distance from the WD center, and $\cos\theta$ the incident angle of the surface (e.g., \cite{sch97}) )."
+ The temperature of the disk vim is assumed to be 3000 K. Figure 2. shows the observational points (open circles) by Schaefer (1990) together with our calculated 5 light curve (thick solid line) for Ace—2.5«10TAL |l , The temperature of the disk rim is assumed to be 3000 K. Figure \ref{bmag_bv_color_paper} shows the observational points (open circles) by Schaefer (1990) together with our calculated $B$ light curve (thick solid line) for $\dot M_{\rm acc}= 2.5 \times 10^{-7} M_\odot$ $^{-1}$.
+"To fit our theoretical light. curves with the observational points. we calculate Blight curves by changing the paralmcters of a=0 —LO by0.1 step. 5=0.05. 0.50 by 0.05 step. Toiars=3500 SOOO IS by 100 IS step. and —Ἰτο B5"" by 1"" step and seek for the best fit model."," To fit our theoretical light curves with the observational points, we calculate $B$ light curves by changing the parameters of $\alpha=0.5$ —1.0 by0.1 step, $\beta=0.05$ —0.50 by 0.05 step, $T_{\rm ph, MS}= 3500$ —8000 K by 100 K step, and $i=75$ $85\arcdeg$ by $1\arcdeg$ step and seek for the best fit model."
+ The best fit parameters obtained are shown in Figure 2 (see also Table 1))., The best fit parameters obtained are shown in Figure \ref{bmag_bv_color_paper} (see also Table \ref{tbl-1}) ).
+ There are five different contributions to the D-lieht (Lp) in the system: the white dwarf (Lp). the nou-radiated portious of the accretion disk (Lp2) aud the donor star (CLpa4). aud the iradiated portious of the accretion disk (55) aud the donor star (Lps).," There are five different contributions to the $B$ -light $L_B$ ) in the system: the white dwarf $L_{B1}$ ), the non-irradiated portions of the accretion disk $L_{B2}$ ) and the donor star $L_{B3}$ ), and the irradiated portions of the accretion disk $L_{B4}$ ) and the donor star $L_{B5}$ )."
+" In order ο show cach contribution. we have added two light curves iu Figure 2.. that is. a non-irraciation case of the ACDIX Gli.pi=O0. dash-dotted). aud a non-nvradiation case of he MS (s=0. dashed).The light frou the WD is completely blocked by the accretion disk riu. thus having ιο contribution. Lp,=0."," In order to show each contribution, we have added two light curves in Figure \ref{bmag_bv_color_paper}, that is, a non-irradiation case of the ACDK $\eta_{\rm ir, DK}=0$, dash-dotted), and a non-irradiation case of the MS $\eta_{\rm ir, MS}=0$, dashed).The light from the WD is completely blocked by the accretion disk rim, thus having no contribution, $L_{B1}=0$."
+ The depth of the primary eclipse. 1.5 mae. means Lp;=0.25Lp because the ACDIS is completely occulted by the MS.," The depth of the primary eclipse, 1.5 mag, means $L_{B3}= 0.25 L_{B}$ because the ACDK is completely occulted by the MS."
+" The ditfereuce of 1 mag )tween the thick solid. auc dash-dotted lines. iudicates Lp,=0.60Lp.", The difference of 1 mag between the thick solid and dash-dotted lines indicates $L_{B4}=0.60 L_B$.
+ The difference of 0.1 mae between the hick solid and dashed lines icicates £p;= 0.10p., The difference of 0.1 mag between the thick solid and dashed lines indicates $L_{B5} = 0.10 L_B$ .
+" Thus. we obtain cach contribution: Lp,= 0. Lp»= 0.05Lp. Lp, =025Lp. Lp,= 060Lp. aud Γης= 0.101)."," Thus, we obtain each contribution: $L_{B1} = 0$ , $L_{B2} = 0.05 L_B$ , $L_{B3} = 0.25 L_B$ , $L_{B4} = 0.60 L_B$ , and $L_{B5}= 0.10 L_B$ ."
+" Then we calculate the theoretical color index (BW), for these best fit models.", Then we calculate the theoretical color index $(B-V)_c$ for these best fit models.
+" Tere. we explain only tle case of Ma,=2.5«10TAL. |."," Here, we explain only the case of $\dot M_{\rm acc}= 2.5 \times 10^{-7} M_\odot$ $^{-1}$ ."
+ By fitting. we obtain the apparent distance modulus of mg= 16.71. which corresponds to the distance of ¢@=22 |kpe without absorption (lp= 0).," By fitting, we obtain the apparent distance modulus of $m_{B, 0}= 16.71$ , which corresponds to the distance of $d= 22$ kpc without absorption $A_B=0$ )."
+ On the other hand. we obtained a rather blue color iudex of (5V)=0.0) outside eclipses.," On the other hand, we obtained a rather blue color index of $(B-V)_c= 0.0$ outside eclipses."
+ Together with the observed color of, Together with the observed color of
+the theoretical framework of the LGRB-SN association is known as the Collapsar model (AlacFaclven&Woosley1999:MacFadvenοἱal.2001).,"the theoretical framework of the LGRB-SN association is known as the Collapsar model \citep{MacWoos99,MacFadyen01}."
+.. According to this model. following the core collapse of a massive star. a bipolar jet is launched at the center of the star.," According to this model, following the core collapse of a massive star, a bipolar jet is launched at the center of the star."
+ The jet diills through the stellar envelope ancl breaks oul of the surface belore producing the observed eamma-rays., The jet drills through the stellar envelope and breaks out of the surface before producing the observed gamma-rays.
+ However. although this model is supported indirectly. bx the association. Co (his date we could not identify a clear direct observational imprint of the jet-envelope interaction. (hus there is no direct confirmation vet of the Collapsar moclel.," However, although this model is supported indirectly by the LGRB-SN association, to this date we could not identify a clear direct observational imprint of the jet-envelope interaction, thus there is no direct confirmation yet of the Collapsar model."
+ In this letter we analvze the expected duration distribution of the prompt GRB emission from Collapsars., In this letter we analyze the expected duration distribution of the prompt GRB emission from Collapsars.
+ We show that under very general conditions the time that the jet spends drilling through the star leads to a plateau in the duration distribution at (mes much shorter than the breakout time of the jet., We show that under very general conditions the time that the jet spends drilling through the star leads to a plateau in the duration distribution at times much shorter than the breakout time of the jet.
+ We examine the duration distribution of all major GRB satellites and find an extended plateau in all of them over a duration range that is expected for reasonable jet-star parameters., We examine the duration distribution of all major GRB satellites and find an extended plateau in all of them over a duration range that is expected for reasonable jet-star parameters.
+ We interpret these plateaus as the first identified imprint of the Collapsar model., We interpret these plateaus as the first identified imprint of the Collapsar model.
+ Under (his interpretation (these findings (1) supports the hypothesis of compact stellar progenitors. Gi) imply the existence of a large population of chocked jets that [ail to break out of the progenitor star and (11) enable us to determine the transition duration Chat statistically separates Collapsars from non-Collapsars bursts. and (o show that this time is individual for each satellite.," Under this interpretation these findings (i) supports the hypothesis of compact stellar progenitors, (ii) imply the existence of a large population of chocked jets that fail to break out of the progenitor star and (iii) enable us to determine the transition duration that statistically separates Collapsars from non-Collapsars bursts, and to show that this time is individual for each satellite."
+ The separation between Collapsars from is quantified and discussed in length in Brombereetal.(2011c)., The separation between Collapsars from non-Collapsars is quantified and discussed in length in \cite{B11c}.
+. The letter is built accordinglv: In section 2 we briefly review (he propagation of a jet inside a star. and the duration of the prompt. GRB emission.," The letter is built accordingly: In section 2 we briefly review the propagation of a jet inside a star, and the duration of the prompt GRB emission."
+ In section 3 we analvze the expected duration distribution of Collapsars., In section 3 we analyze the expected duration distribution of Collapsars.
+ We compare the expected distribution with the observed one in section 4. and discuss our findings in section 5.," We compare the expected distribution with the observed one in section 4, and discuss our findings in section 5."
+" As the jet. propagates in the stellar envelope. it pushes the matter surrounding it. creating a ""head"" of shocked matter at its Iront."," As the jet propagates in the stellar envelope, it pushes the matter surrounding it, creating a “head"" of shocked matter at its front."
+ For tvpical Iuminosities the head remains sub relativistie across most of the star (Matzner2003:Zhangetal.Morsony2007:Alizuta&Alov2009:Brombergοἱal. 201Ib).. even though the jet is ejected at relativistic velocities.," For typical luminosities the head remains sub relativistic across most of the star \citep{Matzner03,Zhang03,Morsony07,Mizuta09,B11a}, even though the jet is ejected at relativistic velocities."
+ The jet breaks out of the stellar surface after (Brombergetal.2011a):: where νεο is (lie isotropic equivalent jet Iuminositv. 8 is the jet half opening angle ancl we have used (vpical values for à long GRD.," The jet breaks out of the stellar surface after \citep{B11b}: where $ L_{iso}$ is the isotropic equivalent jet luminosity, $\theta$ is the jet half opening angle and we have used typical values for a long GRB."
+" 2, and M, are the radius and the mass of the", $ R_{*}$ and $M_{*}$ are the radius and the mass of the
+ (<1 MM.) z—5 .. formation process over time or on the possible role of the environment., $< 1$ $_\odot$ $z \simeq 5$ $_\odot$ formation process over time or on the possible role of the environment.
+ In order to cover a wide range of metallicities and initial densities one is forced to compare stellar systems that formed a Hubble time apart. such as the YCs and GCs mentioned above.," In order to cover a wide range of metallicities and initial densities one is forced to compare stellar systems that formed a Hubble time apart, such as the YCs and GCs mentioned above."
+ Therefore. the universality of the IMF over time and location remains as yet an unresolved issue.," Therefore, the universality of the IMF over time and location remains as yet an unresolved issue."
+ On the other hand. even though the two-body relaxation process that governs the cluster’s dynamical evolution and that leads to preferential loss of low-mass stars is today rather well understood (e.g. Spitzer 1987). it is in general not possible to roll back the effects of dynamics and derive the IMF from the present day MF. since we cannot trace back the trajectories of stars that have escaped the cluster.," On the other hand, even though the two-body relaxation process that governs the cluster's dynamical evolution and that leads to preferential loss of low-mass stars is today rather well understood (e.g. Spitzer 1987), it is in general not possible to roll back the effects of dynamics and derive the IMF from the present day MF, since we cannot trace back the trajectories of stars that have escaped the cluster."
+ It is however possible and statistically meaningful to study the evolution of the stellar MF on a global scale by looking at the differences between clusters at different evolutionary stages., It is however possible and statistically meaningful to study the evolution of the stellar MF on a global scale by looking at the differences between clusters at different evolutionary stages.
+ This requires a homogeneous sample of high quality observations of Galactic (open and globular) clusters. treated in a uniform way and with reliable errors.," This requires a homogeneous sample of high quality observations of Galactic (open and globular) clusters, treated in a uniform way and with reliable errors."
+ Early in this decade. this type of homogeneous study became possible for GCs (Paresce De Marchi 2000). mostly thanks to the Hubble Space Telescope.," Early in this decade, this type of homogeneous study became possible for GCs (Paresce De Marchi 2000), mostly thanks to the Hubble Space Telescope."
+ In the meanwhile. high quality data have become available for YCs as well. mostly from wide field ground based surveys. thereby making this study possible on a global seale.," In the meanwhile, high quality data have become available for YCs as well, mostly from wide field ground based surveys, thereby making this study possible on a global scale."
+ Since our goal is to detect and quantify changes in the shape of the MF in different environments. in. particular in GCs and YCs. we need a functional form for the distributionof," Since our goal is to detect and quantify changes in the shape of the MF in different environments, in particular in GCs and YCs, we need a functional form for the distributionof"
+and 6.8 x [077em-.. respectively.,"and 6.8 $\times$ $^{22}$, respectively."
+ For the (—9755. 22722) peak. we fixed 7 = 15 K because of low signal-to-noise of one of the three lines used in the fit.," For the $-$ 2) peak, we fixed $T$ = 15 K because of low signal-to-noise of one of the three lines used in the fit."
+" Λομ.οη for both peaks 00""7)) and (—9755. 22722) were calculated from the 3 mm interferometric map only and Ny, from the SCUBA 850 um dust continuum: the interferometric observations were convolved to the SCUBA beam size )).", $N_{\rm CH_3OH}$ for both peaks ) and $-$ 2) were calculated from the 3 mm interferometric map only and $N_{\rm H_2}$ from the SCUBA 850 $\mu$ m dust continuum; the interferometric observations were convolved to the SCUBA beam size ).
+ found several ratios. from some CH3;OH-E transitions. to be ealibration-independent tracers of density in. the | and 3 mm bands.," found several ratios, from some $_3$ $E$ transitions, to be calibration-independent tracers of density in the 1 and 3 mm bands."
+" We have used their LVG calculations and have plotted the following integrated intensity line ratios: 59—49/5-)4-1. δι54,/5.,—54.,. and 5,—4)/59-4o. In a density (7-77) plane."," We have used their LVG calculations and have plotted the following integrated intensity line ratios: $5_{0} \rightarrow 4_{0}/5_{-1} \rightarrow 4_{-1}$, $5_{1} \rightarrow 4_{1}/5_{-1} \rightarrow 4_{-1}$, and $5_{1} \rightarrow 4_{1}/5_{0} \rightarrow 4_{0}$, in a temperature-density $T$ $n$ ) plane."
+ In Fig. 7..," In Fig. \ref{Fcolden},"
+" the observed ratios. towards the 755.0"")) position. are represented by the black dash-dotted line. blue solid line. and pink dashed line."," the observed ratios, towards the ) position, are represented by the black dash-dotted line, blue solid line, and pink dashed line."
+ The 2;—1; quartet of lines in the 3 mm band were not used in the analysis because they are blended., The $2_{k} \rightarrow 1_{k}$ quartet of lines in the 3 mm band were not used in the analysis because they are blended.
+ Given the results for CH;OH column densities in. the LTE analysis. the exploredLVG models for the excitation of methanol run for temperatures in the range of 10-200 K. for densities 100—10* em.. and for two CH3OH column densities per line width 10) and 10? p.," Given the results for $_3$ OH column densities in the LTE analysis, the exploredLVG models for the excitation of methanol run for temperatures in the range of 10–200 K, for densities $^5$ $^8$, and for two $_3$ OH column densities per line width $^{14}$ and $^{15}$ )."
+ For an average line width of 4s7!.. this corresponds to Ncu.og-c=4x104 and + x10Pem. respectively.," For an average line width of 4, this corresponds to $N_{\rm CH_3OH-E} = 4 \times 10^{14}$ and 4 $\times 10^{15}$, respectively."
+ Depending on the column density. two possible solutions are found: at Ncu.ogcp=4x104 em. the three integrated intensity line ratios intercept at 7= 17 K and à = 10!οι”.," Depending on the column density, two possible solutions are found: at $N_{\rm CH_3OH-E} = 4 \times 10^{14}$ , the three integrated intensity line ratios intercept at $T =$ 17 K and $n$ = $^7$."
+ On the other hand. at Ney-ou_p=4x10?em. the interception is at T = 60 K and 1=2x10° (see Fig. 7)).," On the other hand, at $N_{\rm CH_3OH-E} = 4 \times 10^{15}$, the interception is at $T$ = 60 K and $n = 2 \times 10^6$ (see Fig. \ref{Fcolden}) )."
+ The methanol column density inferred in this case is consistent with the LTE analysis within a factor of 3. but the intensities in this model are one order of magnitude higher than the observed line intensities.," The methanol column density inferred in this case is consistent with the LTE analysis within a factor of 3, but the intensities in this model are one order of magnitude higher than the observed line intensities."
+" In other words. to fit the observations with the Nceu.og-u=4X10"" model. we need a beam filling factor of 0.1 resulting in à more compact source (with a size of 0755 corresponding to 0.009 pe or 1800 AU). even smaller than the one modelled by and?."," In other words, to fit the observations with the $N_{\rm CH_3OH-E} = 4 \times 10^{15}$ model, we need a beam filling factor of 0.1 resulting in a more compact source (with a size of 5 corresponding to 0.009 pc or 1800 AU), even smaller than the one modelled by and."
+. To gain more insight into the possible 7-7 degeneracy. in Figure 8.. we show n vs. Ney.on plots. at different temperatures (10-70 K). where the color-coding for the ratios (along with their Io values) are the same as in Fig. 7:," To gain more insight into the possible $T$ $n$ degeneracy, in Figure \ref{FnN}, , we show $n$ vs. $N_{\rm CH_3OH}$ plots, at different temperatures (10-70 K), where the color-coding for the ratios (along with their $\sigma$ values) are the same as in Fig. \ref{Fcolden};"
+ additionally. we show in each panel the peak main beam brightness temperature of the 54-ρε transition (green dotted line).," additionally, we show in each panel the peak main beam brightness temperature of the $_{0} \rightarrow 4_{0}E$ transition (green dotted line)."
+ The parameters obtained in the regions in the left side of the green dotted line imply beam filling factors greater than 1. which is not possible.," The parameters obtained in the regions in the left side of the green dotted line imply beam filling factors greater than 1, which is not possible."
+ In each panel. the areas of interest are those that lie in the right side of the green line and. at the same time. are delimited by the integrated intensity ratios (taking into account their Loc values).," In each panel, the areas of interest are those that lie in the right side of the green line and, at the same time, are delimited by the integrated intensity ratios (taking into account their $\sigma$ values)."
+ We see that. for T in the range 30-70 K. the 7 remains constant up to Neu.opg ~ 4x 107em™.. while for the very low T (= 10 K). there is no solution.," We see that, for $T$ in the range 30-70 K, the $n$ remains constant up to $N_{\rm CH_3OH}$ $\sim$ 4 $\times$ $^{14}$, while for the very low $T$ (= 10 K), there is no solution."
+ Given our density estimate from the dust continuum. we view the low 7. high z solutions with a large filling factor as unlikely.," Given our density estimate from the dust continuum, we view the low $T$ , high $n$ solutions with a large filling factor as unlikely."
+ Therefore. although we cannot strongly constrain the physicalparameters of the gas emitting methanol from the LVG analysis because ofthe limited number of observed transitions.," Therefore, although we cannot strongly constrain the physicalparameters of the gas emitting methanol from the LVG analysis because ofthe limited number of observed transitions,"
+Auger Collaboration if the latter is finally confirmed.,Auger Collaboration if the latter is finally confirmed.
+ This is reasonable. as only a few tens of ACGNs from the Vérron-Cetty catalogue are located at a distance of D<10Mpc. but it is important to keep in mind that ΝΟ σοι [rom larger distances can be heavily suppressed.," This is reasonable, as only a few tens of AGNs from the V{\'e}rron-Cetty catalogue are located at a distance of $D<10$ Mpc, but it is important to keep in mind that CNO UHECRs from larger distances can be heavily suppressed."
+ We have explored the possibility of the correlation of nearby ealaxies [rom the catalogue of ? with ultra high energy cosmic rav events from the Pierre Auger Collaboration., We have explored the possibility of the correlation of nearby galaxies from the catalogue of \citet{2004AJ....127.2031K} with ultra high energy cosmic ray events from the Pierre Auger Collaboration.
+ This analysis is highly motivated in the light that a mixed CNO (as opposed to pure proton) composition of ULLECRs. which is the preferred scenario indicated by current data. could not survive the GZlIx cutoll unless the sources are located in the Local Volume.," This analysis is highly motivated in the light that a mixed CNO (as opposed to pure proton) composition of UHECRs, which is the preferred scenario indicated by current data, could not survive the GZK cutoff unless the sources are located in the Local Volume."
+ Therefore. any correlation with these nearby sources would explain both composition and correlation at the same time.," Therefore, any correlation with these nearby sources would explain both composition and correlation at the same time."
+ The analvsis of the data concludes a 2.60 correlation between both. samples., The analysis of the data concludes a $2.6\sigma$ correlation between both samples.
+ Phere are 5 out. of 27 correlating events. while 0.5 should be expected on average in the case of isotropic Hux.," There are 5 out of 27 correlating events, while 0.5 should be expected on average in the case of isotropic flux."
+ Due to the small number of correlating events. we cannot specify the scan parameters a priori. thus making our results sensitive to the negative inipact of a posteriori anisotropy searches.," Due to the small number of correlating events, we cannot specify the scan parameters a priori, thus making our results sensitive to the negative impact of a posteriori anisotropy searches."
+ The probability of chance correlation due to isotropic flux after penalizing or scans Is Pace=0.96.10.7. rejecting the hypothesis of isotropy of arrival clirections of ULLECRs at confidence evel.," The probability of chance correlation due to isotropic flux after penalizing for scans is $P_{\rmn{chance}}=0.96\times 10^{-2}$, rejecting the hypothesis of isotropy of arrival directions of UHECRs at confidence level."
+ For comparison. the same result. before penalizing for scans is —1510.I which is similar to the ~2«10.! result (with scan parameters set a priori) from the correlation ound by 7. with the ACGNs in the Vérron-Cetty catalogue. and stronger than the revised. value of ~610 when hey include the new CLEC. events from 1 September 2007 o 31 March. 209 (2)..," For comparison, the same result before penalizing for scans is $\sim 1\times 10^{-4}$ which is similar to the $\sim 2\times 10^{-4}$ result (with scan parameters set a priori) from the correlation found by \citet{2008APh....29..188T} with the AGNs in the Vérron-Cetty catalogue, and stronger than the revised value of $\sim 6\times 10^{-3}$ when they include the new UHECR events from 1 September 2007 to 31 March 2009 \citep{2009arXiv0906.2347T}."
+ This is an indication that there is no clear source of the ULLECTIU. and finding a lower value of the probability of chance correlation does not. guarantee hat this is actually the preferred. source as compared. to other sources which present similar values of Pucs.," This is an indication that there is no clear source of the UHECRs, and finding a lower value of the probability of chance correlation does not guarantee that this is actually the preferred source as compared to other sources which present similar values of $P_{\rmn{chance}}$."
+ Fhis was already noticed. even in carly papers such as 2.. and in any case it seems unlikely to get a 3.3m detection (?)..," This was already noticed even in early papers such as \citet{2001JETPL..74..445T}, and in any case it seems unlikely to get a $>3.3\sigma$ detection \citep{2008JCAP...05..006K}."
+ Interestingly. most of the objects in the catalogue of ? which correlate with Auger events. are already. included in the AGN catalogue. which has around 30 objects with D« 10Mpc.," Interestingly, most of the objects in the catalogue of \citet{2004AJ....127.2031K} which correlate with Auger events, are already included in the AGN catalogue, which has around 30 objects with $D<10$ Mpc."
+ This may indicate that the relation between nearby galaxies and VHECRSs found in our work could. be inherited from the ? result., This may indicate that the relation between nearby galaxies and UHECRs found in our work could be inherited from the \citet{2008APh....29..188T} result.
+" In any case. the so-called AGN hypothesis has proven to be controversial (ον, 2)). stressed bv the deficit. of events [rom the direction. of the Virgo supercluster."," In any case, the so-called AGN hypothesis has proven to be controversial \citealt{2008JETPL..87..461G}, \citealt{2009ApJ...693.1261M}) ), stressed by the deficit of events from the direction of the Virgo supercluster."
+ On the other hand. the correlation with the distribution of luminous matter seems to be well established (?2).. without making any reference to any particular source which can accelerate particles up to these energies.," On the other hand, the correlation with the distribution of luminous matter seems to be well established \citep{2008JCAP...05..006K}, without making any reference to any particular source which can accelerate particles up to these energies."
+ The small number of correlating events. with nearby galaxies might also point to the possibility of the absence of VLECK. accelerators in our local environment. although there is no definitive answer vet.," The small number of correlating events with nearby galaxies might also point to the possibility of the absence of UHECR accelerators in our local environment, although there is no definitive answer yet."
+ Surprisingly. most of the objects in the ? catalogue which correlate with the ULDECT sample (with the exception of NGCS300) are located in a small region in the sky. with Supergalactic coordinates 150«SCL165° and. 10«SOW57.," Surprisingly, most of the objects in the \citet{2004AJ....127.2031K} catalogue which correlate with the UHECR sample (with the exception of NGC300) are located in a small region in the sky, with Supergalactic coordinates $150^{\circ}<SGL<165^{\circ}$ and $-10^{\circ}<SGB<-5^{\circ}$."
+ phis region is near the Supergalactic plane. pointing to preferential arrival directions from. nearby galaxies located in a high density. region.," This region is near the Supergalactic plane, pointing to preferential arrival directions from nearby galaxies located in a high density region."
+ Future measurements and hopefully the next data release with a total of 58 events as announced in ? will help to clarify the origin of nearby ancl clistant UIIECIHUs. which will allow to παν more in detail the acceleration physies behind The authors acknowledge support. from the Spanish MEC under. grant PNAYA 2005-07789.," Future measurements and hopefully the next data release with a total of 58 events as announced in \citet{2009arXiv0906.2347T} will help to clarify the origin of nearby and distant UHECRs, which will allow to study more in detail the acceleration physics behind The authors acknowledge support from the Spanish MEC under grant PNAYA 2005-07789."
+ AJC. thanks support of the ALEC through Spanish grant. FPL AP2005-1826., A.J.C. thanks support of the MEC through Spanish grant FPU AP2005-1826.
+ We especially want to thank Vasiliki Pavlidou from CalTech Astronomy Department for useful comments on this paper., We especially want to thank Vasiliki Pavlidou from CalTech Astronomy Department for useful comments on this paper.
+ This work made extensive use of the x4600 machine at CICA (Spain)., This work made extensive use of the x4600 machine at CICA (Spain).
+ We thank Claudio Arjona and Ana Silva at CLOA for making this machine available to us., We thank Claudio Arjona and Ana Silva at CICA for making this machine available to us.
+&ood agreement with Schmidtetal.(1992) within the error limit.,good agreement with \citet{schmidt} within the error limit.
+ 3oth the program and the stanclare stars were observe during the same night., Both the program and the standard stars were observed during the same night.
+ LID) 25433 and LLD 19820 (= CC Cas) were used as a standard polarized stars. while LID 21447 anc G 191-B2B (= HIP. 23692) served as standard unpolarizec stars.," HD 25433 and HD 19820 (= CC Cas) were used as a standard polarized stars, while HD 21447 and G 191-B2B (= HIP 23692) served as standard unpolarized stars."
+ The results are listed in Tables 3.. 4. anc 5.. where P is the fraction of the total light in the linearly. polarize condition and @ is the position angle of polarization plane to the equatorial plane.," The results are listed in Tables \ref{table:polar1}, , \ref{table:polar2} and \ref{table:polar3}, where $P$ is the fraction of the total light in the linearly polarized condition and $\theta$ is the position angle of polarization plane to the equatorial plane."
+ Lo is denoted. by the normalizes Stokes” parameter g(=(Q1). when the hall wave plates fast axis is alignede to the reference axis (a= 0°).," It is denoted by the normalized Stokes' parameter ${\it q}\ (=Q/I)$, when the half wave plate's fast axis is aligned to the reference axis $\alpha = 0^\circ$ )."
+" Similarly. the normalized Stokes! parameter :(=Cf). when the half wave plate is at 22.5""."," Similarly, the normalized Stokes' parameter ${\it u}\ (=U/I)$, when the half wave plate is at $^\circ$."
+ For further details on used equipmen and the method of observations. refer to Mecdhietal. (2007).," For further details on used equipment and the method of observations, refer to \citet{medhi}. ."
+. By dotted. lines on the lighteurve. we indicate. the phases (giτο0.13 for 2007 October 19 and 20 respectively). when our polarimetric observations were conducted (see Fie. 6)).," By dotted lines on the lightcurve we indicate the phases $\varphi_1 = 0.16, \varphi_2 = 0.13$ for 2007 October 19 and 20 respectively), when our polarimetric observations were conducted (see Fig. \ref{photometry2007}) )."
+ FR. €ne was in maximal brightness during that time., FR Cnc was in maximal brightness during that time.
+ Polarization in FR Cne could be magnetic in origin., Polarization in FR Cnc could be magnetic in origin.
+ The degree of polarization depends nonlinearly on the size of magnetic regions (see Lluovelin&Saar 199]: Saar&Lluovelin 1993)).," The degree of polarization depends nonlinearly on the size of magnetic regions (see \citeauthor{polar1} \citeyear{polar1}; ; \citeauthor{polar2}
+ \citeyear{polar2}) )."
+" ""μονο authors have also calculated a grid ol expected degrees of polarization in CDVIH band for stars with temperature from 4000 to 7000 Ix and log g from 2.0 to 4.5.", These authors have also calculated a grid of expected degrees of polarization in $UBVRI$ band for stars with temperature from 4000 to 7000 K and log $g$ from 2.0 to 4.5.
+" We have used their results to compare our observec values of polarization for FR πο,", We have used their results to compare our observed values of polarization for FR Cnc.
+ Fig., Fig.
+ 8. represents the degree of polarization. for FR Cnc in BYR bands., \ref{fig:model} represents the degree of polarization for FR Cnc in $BVR$ bands.
+" The maximum possible degree of polarization for the total spo area of 24 per cent is derived from the calculations of Saar&Lluovelin(1993) for the star corresponding to the spectra tvpe of and characteristic magnetic field of 2.7 kG. ‘These values are over-plotted and represented by a solid line in Fig. δι, ", The maximum possible degree of polarization for the total spot area of 24 per cent is derived from the calculations of \citet{polar2} for the star corresponding to the spectral type of and characteristic magnetic field of 2.7 kG. These values are over-plotted and represented by a solid line in Fig. \ref{fig:model}. .
+Phe observed polarization in the £3 band is in goo agreement with the theoretical values expected for Zeeman polarization model., The observed polarization in the $B$ band is in good agreement with the theoretical values expected for Zeeman polarization model.
+ Llowever. the observed polarization in V and A? bands slightly exceeds the theoretical values.," However, the observed polarization in $V$ and $R$ bands slightly exceeds the theoretical values."
+ Alodel values for. K2V-KTV and IN21V-IN2V. spectral types appear to be even lower and certainly do not mateh the polarization in any of observed bands., Model values for K2V-K7V and K2IV-K2V spectral types appear to be even lower and certainly do not match the polarization in any of observed bands.
+ This is also observed in some other voung spotted stars (MS Ser. LO HHya. VY Ari: see Alekseey 2003)) and. probably. due to the presence of a supplementary source of linoar polarization.," This is also observed in some other young spotted stars (MS Ser, LQ Hya, VY Ari; see \citeauthor{alekseev}
+ \citeyear{alekseev}) ) and probably due to the presence of a supplementary source of linear polarization."
+" The preclieted values of polarization due to Thompson and Ravici¢h scattering from inhomogeneous regions are not enough to explain the observed polarization excess CEhompson and Ravleigh scattering for the assumedspectraltvpe supposed to beof order of 10."" and. 10*per cent respectively: Saar& Lluovelin 1993))."," The predicted values of polarization due to Thompson and Rayleigh scattering from inhomogeneous regions are not enough to explain the observed polarization excess (Thompson and Rayleigh scattering for the assumedspectraltype supposed to beof order of $10^{-7}$ and $10^{-4}$per cent respectively; \citeauthor{polar2}
+ \citeyear{polar2}) )."
+ The mechanism which can produce additional linear polarization is probably, The mechanism which can produce additional linear polarization is probably
+in (he hierarchical picture). the remarkable homogeneity in the properties of cluster galaxies observed by Christlein&Zabludoll(2003) and DeProprisetal.(2003) would result naturally.,"in the hierarchical picture), the remarkable homogeneity in the properties of cluster galaxies observed by \cite{christlein03} and \cite{depropris03} would result naturally."
+ We would like to thank the anonvmous referee for his informative report., We would like to thank the anonymous referee for his informative report.
+ AR and RDP acknowledge funding through the Ulx Particle Physies and Astrophysics Research Council (PPARC)., AR and RDP acknowledge funding through the UK Particle Physics and Astrophysics Research Council (PPARC).
+have gradually revealed their morphology in detail Cl'avlorCGuvon.Sanders.&Stockton 2006)..,"have gradually revealed their morphology in detail \citep{Ta96,Ba97,Pe01,MM01,Ma01,Ha02,Du03,Ve06,Gu06}."
+ They found. not only the elliptical. (whic ris equivalent to the bulge. classified as the spheroicl) component but also the prominent. disk component in the host of lower luminosity QSOs and raclio-quiet. QSOs (Dunlopetal.2003).," They found not only the elliptical (which is equivalent to the bulge, classified as the spheroid) component but also the prominent disk component in the host of lower luminosity QSOs and radio-quiet QSOs \citep{Du03}."
+. Also. we now know that QSO host galaxies are gas-rich (which is not normal for elliptical galaxies) (Evansetal.2001.2006:Scoville and that starbursts occur in the host galaxies (Llaaswewa," Also, we now know that QSO host galaxies are gas-rich (which is not normal for elliptical galaxies) \citep{Ev01,Ev06,Sc03} and that starbursts occur in the host galaxies \citep{Ha03,Ba06,Sc06,Ma07,Ne07,Sh07,Lu08}."
+ his letter. to clarifv the SAIBLDE growth. mechanism lal saislies the observational results. we investigated the UuΑΗΣΗ masshost starburst connection. taking into account 10 hos galaxy morphology for Sevfert galaxies and. QSOs.," In this letter, to clarify the SMBH growth mechanism that satisfies the observational results, we investigated the SMBH mass–host starburst connection, taking into account the host galaxy morphology for Seyfert galaxies and QSOs."
+" Throughout the paper. we adopted. 44,=SOkms!Mpe.7. Oy,—0.3. and Oy—0."," Throughout the paper, we adopted $H_{0} = 80\, \rm km\,s^{-1} Mpc^{-1}$, $\Omega_{M}=0.3$, and $\Omega_{\Lambda}=0.7$."
+ To investigate the clleet of starbursts ancl host galaxy morphology on the final SAIBLL mass. we selected ACGNs based. οἳ estimated SMDIL mass. PALL emission. and morphology.," To investigate the effect of starbursts and host galaxy morphology on the final SMBH mass, we selected AGNs based on estimated SMBH mass, PAH emission, and morphology."
+ The samples in this paper are Palomar-Cireen (DG) OSOs. scleced at. B band to have blue (2B color (Schmidt Green 1983).," The samples in this paper are Palomar-Green (PG) QSOs, selected at $B$ band to have blue $U-B$ color (Schmidt Green 1983)."
+ Since the B- ancl U-band emission of PG-OSOs are dominated by the AGN. there are no biases [or the presence of starbursts.," Since the B- and U-band emission of PG-QSOs are dominated by the AGN, there are no biases for the presence of starbursts."
+ Sevlert galaxies in the CLA (lluchra&Bure1992). and 12 jin (Rush.Malkan.&Spinoglio1993). samples. selected on the basis of their host galaxy magnitudes and £2245 12 jum dFHuxes. respecively.," Seyfert galaxies in the CfA \citep{HB92} and 12 $\mu$ m \citep{Ru93} samples, selected on the basis of their host galaxy magnitudes and $IRAS$ 12 $\mu$ m fluxes, respectively."
+ These samples do not also include some biases for the presence of starbursts., These samples do not also include some biases for the presence of starbursts.
+ ‘To estimate SAIBLD mass. Adgg. of PG-QSOs and. Sevfer 1. galaxies. we used a method based on the reasonable assumption that the motion of ionized gas clouds.arouim the SMDLL is dominated by the gravitational force. ancl tha the clouds within the broad-line region (BLAIR) are virializec (c.g..Peterson&Wandel 1999)...," To estimate SMBH mass, $M_{\rm BH}$, of PG-QSOs and Seyfert 1 galaxies, we used a method based on the reasonable assumption that the motion of ionized gas cloudsaround the SMBH is dominated by the gravitational force, and that the clouds within the broad-line region (BLR) are virialized \citep[e.g.,][]{PW99}. ."
+ Phe velocity. dispersion Uo can be estimated. from the EWIIM. of 13 broad line emission ο=feywgsy by assuming the isotropic spherica virial coellicicnt. f=V3/2 (Netzer 1990).," The velocity dispersion $v$ can be estimated from the FWHM of $\beta$ broad line emission $v=fv_{\rm FWHM}$ by assuming the isotropic spherical virial coefficient, $f=\sqrt{3}/2$ (Netzer 1990)."
+ We selector the LL3 because LL? lines racdiate by simple photoionization mechanisms and that their line profiles rellect. gravitationa potential of the LL3 emission region., We selected the $\beta$ because $\beta$ lines radiate by simple photoionization mechanisms and that their line profiles reflect gravitational potential of the $\beta$ emission region.
+ Adopting an empirica relationship (Ixaspietal.2000) between the size of the DBLItand the rest-frame optical continuum. Iuminositv. ALX(5100:1). and using reverberation mapping. we obtain the following formula:," Adopting an empirical relationship \citep{Ka00} between the size of the BLRand the rest-frame optical continuum luminosity, $\lambda L_{\lambda} (5100 \AA)$ , and using reverberation mapping, we obtain the following formula:"
+limit is delined by a 12-Ciyr isochrone with [Fe/H] = —0.3.,limit is defined by a 12-Gyr isochrone with [Fe/H] = $-0.3$.
+ For the bluer stars. a fixed boundary in AM at a given color will allow a larger range of evolved turuoll stars to be included iun the sample as [Fe/H] decreases.," For the bluer stars, a fixed boundary in $M_V$ at a given color will allow a larger range of evolved turnoff stars to be included in the sample as [Fe/H] decreases."
+ Thus. the shift in the average A4 of the sample above the unevolved main sequence at a given color will increase as [Fe/H] decreases.," Thus, the shift in the average $M_V$ of the sample above the unevolved main sequence at a given color will increase as [Fe/H] decreases."
+ Qualitatively. this should steepen the slope of the mean main sequence with decreasing [Fe/H] while compressing the range in Ady ata eiven color for a given range in [Fe/H]. leadiug to a smaller derived ratio of .NAA with [Fe/H].," Qualitatively, this should steepen the slope of the mean main sequence with decreasing [Fe/H] while compressing the range in $M_V$ at a given color for a given range in [Fe/H], leading to a smaller derived ratio of $\Delta M_V$ with [Fe/H]."
+ As a second test οἱ the revised main sequence relation. we have adjusted all parallax. stars between [Fe/H] = —0.05 aud 0.05 to an assumed [Fe/H] = 0.0 aud calculated residuals relative to the adjusted Hyacles relation shifted to [Fe/H] = 0.0.," As a second test of the revised main sequence relation, we have adjusted all parallax stars between [Fe/H] = $-0.05$ and 0.05 to an assumed [Fe/H] = 0.0 and calculated residuals relative to the adjusted Hyades relation shifted to [Fe/H] = 0.0."
+ Excluding all stars with absolute residuals ereater than 0.5 mag. for B—V. greater than 0.72. 130 stars produce A:Vy = 0.038. + 0.005 (s.e.imi.).," Excluding all stars with absolute residuals greater than 0.5 mag, for $B-V$ greater than 0.72, 130 stars produce $\Delta M_V$ = +0.038 $\pm$ 0.005 (s.e.m.)."
+ Placing the color cut at 8—V = 0.90 reduces the offset to +0.033 + 0.008 (s.e.m.), Placing the color cut at $B-V$ = 0.90 reduces the offset to +0.033 $\pm$ 0.008 (s.e.m.)
+ for mi stars., for 57 stars.
+ The conclusion is that application of the mean relation at solar metallicity to a cluster with le same properties as the nearby solar field stars will leac to an apparent modulus that potentially is too stuall by 0.03+) mae., The conclusion is that application of the mean relation at solar metallicity to a cluster with the same properties as the nearby solar field stars will lead to an apparent modulus that potentially is too small by 0.03 mag.
+ With this caveat in mind. we can now derive the apparent modulus or je. well-stuclied ope: ‘Luster. MOT.," With this caveat in mind, we can now derive the apparent modulus for the well-studied open cluster, M67."
+ Unlike the field) and. Hyades stars iucluded tu the current study. M6T is recddened. aud iol ied directly to the spectroscopic scale or theTycho-2 BV. system.," Unlike the field and Hyades stars included in the current study, M67 is reddened and not tied directly to the spectroscopic scale or the $BV$ system."
+ Within the extensive literatire dealing with the fancdaimental properties o{ M67. the two most relevant receut studies are ancl Pasquiuietal.(2008)..," Within the extensive literature dealing with the fundamental properties of M67, the two most relevant recent studies are \citet{an07} and \citet{pa08}."
+ The former builds upon the approach of Pinsouneaultet2001) to construct semiempirical isochroues ini uinltiple colors to simultaneously defiue the redcdeuiug. metallicity.and distance for a number of open clusters. including MOT.," The former builds upon the approach of \citet{pi03,pi04} to construct semiempirical isochrones in multiple colors to simultaneously define the reddening, metallicity,and distance for a number of open clusters, including M67."
+ The final valtles derived for M67 are ECB—V) = 0.012. [Fe/H] = —0.02 αιda Gu—M) = 9.71.," The final values derived for M67 are $(B-V)$ = 0.042, [Fe/H] = $-0.02$ and $(m-M)$ = 9.74."
+ The latter paper uses high-resolution spectra and line-clepth-ratios. Ha lines. aud. Li measures tο ideutifv stars witliu M67 that resemble solar analogs.," The latter paper uses high-resolution spectra and line-depth-ratios, $\alpha$ lines, and Li measures to identify stars within M67 that resemble solar analogs."
+ With a derived solar coor of B—V = 0.619. E(B—V) = 0.011 and [Fe/H] = 0.01. Pasquinietal.(23008). find G2—AL) = 9.78.," With a derived solar color of $B-V$ = 0.649, $(B-V)$ = 0.041 and [Fe/H] = +0.01, \citet{pa08} find $(m-M)$ = 9.78."
+ we fix E(B—V) at 0.011 aud set [Fe/H] — 0.00. the two investigations iuply Ga—AL) = 9.76 aud 9.77. 'espectively.," If we fix $(B-V)$ at 0.041 and set [Fe/H] = 0.00, the two investigations imply $(m-M)$ = 9.76 and 9.77, respectively."
+ The last piece of the puzzle requires a link between the adopted 51' photometry for À67 (Sandeuist2001) aud that of the trausfo1uedTycho-2 colors., The last piece of the puzzle requires a link between the adopted $BV$ photometry for M67 \citep{sa04} and that of the transformed colors.
+ The ouly di‘ect comparison is that ol Joueretal.(2008) where the Hyacles. 167. aud trausforimed Tycho-2 data are coupled throieh SAAO observations. all ultimately tied to SAAO E-reeion standards.," The only direct comparison is that of \citet{jon08} where the Hyades, M67, and transformed data are coupled through SAAO observations, all ultimately tied to SAAO E-region standards."
+ Withi1 the uncertainties. he uultiple comparisons iuply that the color olfsets required to trauslorm tle Hyades BV data to the Tyrho-2/SAAO system are similar to those applicable to he Saudquist.(001) BV. data of M67 tied to the Johuso1 system through the standards of and Landolt(1992).," Within the uncertainties, the multiple comparisons imply that the color offsets required to transform the \citet{joh55} Hyades $BV$ data to the /SAAO system are similar to those applicable to the \citet{sa04} $BV$ data of M67 tied to the Johnson system through the standards of \citet{ste00} and \citet{lan92}."
+. We therefore have aclclec 0.009 mag to the B—V indices of M67., We therefore have added 0.009 mag to the $B-V$ indices of M67.
+ For V. a direct link between the M67 pliXometry of Sandquist(2001) and the Tycho-2/85AAO system is unavailable.," For $V$, a direct link between the M67 photometry of \citet{sa04} and the /SAAO system is unavailable."
+ Staudards for the N167 CCD calibration are from to the systems of and Laudolt (1992): comparisous with other M67 sources indicate that the V magnitudes are a match to this standard system at the 0.001 mag level., Standards for the M67 CCD calibration are from to the systems of \citet{ste00} and \citet{lan92}; ; comparisons with other M67 sources indicate that the \citet{sa04} $V$ magnitudes are a match to this standard system at the $\pm$ 0.004 mag level.
+ By default. no adjustiuent is mace to the M67 V. photometry.," By default, no adjustment is made to the M67 $V$ photometry."
+"trend of the density profile to have the same value as the mean matter density in the Universe, but there is also a dependence on the coefficients b; and be.","trend of the density profile to have the same value as the mean matter density in the Universe, but there is also a dependence on the coefficients $b_{1}$ and $b_{2}$."
+ In Figure 8 we can also see that this dependence makes our approximation slightly shifted by a constant value bi+bo with respect to the data from cosmological simulations., In Figure 8 we can also see that this dependence makes our approximation slightly shifted by a constant value $b_1+b_2$ with respect to the data from cosmological simulations.
+" This means that the O(1/r) terms in our approximation are still important in the range from 10 to ~30R,;. just before entering the asymptotic regime where the density profile tends to the constant mean matter density p.", This means that the $\mathcal{O}(1/r)$ terms in our approximation are still important in the range from 10 to $\sim 30R_{\rmn{vir}}$ just before entering the asymptotic regime where the density profile tends to the constant mean matter density $\bar{\rho}$.
+" This shows the limitations of our approximation: while it provides a reasonable fit up to 10 virial radii, the goodness of the fit is not so good at larger distances."," This shows the limitations of our approximation: while it provides a reasonable fit up to 10 virial radii, the goodness of the fit is not so good at larger distances."
+" This is mainly due to the simplicity of the approximation, which is one of the premises in our model."," This is mainly due to the simplicity of the approximation, which is one of the premises in our model."
+" Of course, we can"," Of course, we can"
+satellites.,satellites.
+ We compared the observational clata available on the irregular satellites with those on the minor boclics matching the cvnamical range of the solutions to the inverse capture problem. and we found that. within the errors. there is enough match to strengthen our dynamical Finally. we used the data supplied by Cassini to tes the viability. of the collisional capture process in. noncatastrophic impact scenarios and to constrain the impac configuration and the impactor mass by comparing our results with the information available on Jason crater.," We compared the observational data available on the irregular satellites with those on the minor bodies matching the dynamical range of the solutions to the inverse capture problem, and we found that, within the errors, there is enough match to strengthen our dynamical Finally, we used the data supplied by Cassini to test the viability of the collisional capture process in non--catastrophic impact scenarios and to constrain the impact configuration and the impactor mass by comparing our results with the information available on Jason crater."
+ We found that highly oblique impacts could have produces Jason crater and. supplied the right order of magnitude of specific impulse. the requirement on obliquity being imposec by the high specific impulse needed to capture Phoebe.," We found that highly oblique impacts could have produced Jason crater and supplied the right order of magnitude of specific impulse, the requirement on obliquity being imposed by the high specific impulse needed to capture Phoebe."
+" The range of values of the impact angle associated with our calcul:ions is characterised by a low probability. making a ""sinele collision” scenario to appear implausible."," The range of values of the impact angle associated with our calculations is characterised by a low probability, making a “single collision” scenario to appear implausible."
+ Llowever. in interpreting the results exposed we must take into accoun the following fact.," However, in interpreting the results exposed we must take into account the following fact."
+ Phoebe's surface indicates that. the satellite had an extremely intense collisional history. with a least 7 craters bigger than 50Am and 2 (Jason and IEurytus) of the order of 100&in.," Phoebe's surface indicates that the satellite had an extremely intense collisional history, with at least $7$ craters bigger than $50\,km$ and $2$ (Jason and Eurytus) of the order of $100\,km$."
+ As a consequence. the dynamica evolution of the satellite likely suffered more than one major change in the past and the formation of Eurvtus crater coulc be linked to Phoebe's capture as well.," As a consequence, the dynamical evolution of the satellite likely suffered more than one major change in the past and the formation of Eurytus crater could be linked to Phoebe's capture as well."
+ Since we lack detaile information on lIurvtus morphology. we cannot refine further the model.," Since we lack detailed information on Eurytus' morphology, we cannot refine further the model."
+" However. à. 7two collisions"" scenario would relax the constraints on the impact angle. thus making the whole scenario more plausible."," However, a “two collisions” scenario would relax the constraints on the impact angle, thus making the whole scenario more plausible."
+ Moreover. the results we reported. in Paper bL indicate that. Phoche’s elliciency in collisionally removing bodies on retrograde orbits is significant only in a limited region of phase space (i.e. the one that would. be occupied. by possible collisional fragments).," Moreover, the results we reported in Paper I indicate that Phoebe's efficiency in collisionally removing bodies on retrograde orbits is significant only in a limited region of phase space (i.e. the one that would be occupied by possible collisional fragments)."
+" Since we do not observe retrograde satellites in. Phoebe's gap (Le. the region comprised between 11.2210""Am and 14.97QAm [rom Saturn. see Paper 1) and. Phoebe's sweeping cHeet cannot be invoked to explain this fact. the cHicieney of the capture process should be low enough to explain the absence of other retrograde satellites."," Since we do not observe retrograde satellites in Phoebe's gap (i.e. the region comprised between $11.22\times10^6\,km$ and $14.97\times10^6\,km$ from Saturn, see Paper I) and Phoebe's sweeping effect cannot be invoked to explain this fact, the efficiency of the capture process should be low enough to explain the absence of other retrograde satellites."
+" “Phe low frequency associated to the needed impact parameters in the ""single collision” scenario could help explaining why Phoebe is the only retrograde satellite residing in this orbital region.", The low frequency associated to the needed impact parameters in the “single collision” scenario could help explaining why Phoebe is the only retrograde satellite residing in this orbital region.
+ All the authors wish to thank the referee. Anthony It. Dobrovolskis. for help and suggestions to improve this paper.," All the authors wish to thank the referee, Anthony R. Dobrovolskis, for help and suggestions to improve this paper."
+ DT. wishes to thank Alessandro Morbicelli. Torrence Johnson and Luke Dones for fruitful discussions and useful comments on the results and their interpretation.," D.T. wishes to thank Alessandro Morbidelli, Torrence Johnson and Luke Dones for fruitful discussions and useful comments on the results and their interpretation."
+test the hypothesis whether the actual proportion p; exceeds the ent off value on the basis of an one sided confidence interval for py based on the [ast bootstrap.,test the hypothesis whether the actual proportion ${p}_{k}$ exceeds the cut off value on the basis of an one sided confidence interval for ${p}_{k}$ based on the fast bootstrap.
+ The target of the mixture model method developed by Qui Tambhane (2007) is to divide » observations into A(«n) elusters so that the within cluster variations are very small and any (wpe of dissimilarity must occur between clusters., The target of the mixture model method developed by Qui Tamhane (2007) is to divide $n$ observations into $K (<n)$ clusters so that the within cluster variations are very small and any type of dissimilarity must occur between clusters.
+ In (his work Qui Tamhane (2007) have assumed Ix to be known., In this work Qui Tamhane (2007) have assumed K to be known.
+ In our analysis. we have computed the value of Ix by Ix-means clustering (AlacQueen 1967) and a statistical technique developed by Sugar James (2003).," In our analysis, we have computed the value of K by K-means clustering (MacQueen 1967) and a statistical technique developed by Sugar James (2003)."
+ By using this algorithm we have first. determined the structures of subpopulations (clusters) for varving number of clusters taking Ix = 1. 2. 3. 4 ete.," By using this algorithm we have first determined the structures of subpopulations (clusters) for varying number of clusters taking K = 1, 2, 3, 4 etc."
+" For each such cluster formation we have computed the values of a distance measure dy=(l/pymnin,E|(ey—cxPGreyeK)] which is defined as the distance of the wry vector (values of the parameters) from the center ey. (which is estimated as mean value). p is (he order of the ry, vector."," For each such cluster formation we have computed the values of a distance measure $d_{K} = (1/p) min_{x} E[( x_{K}
+- c_{K} )^{'}(x_{K} - c_{K})]$ which is defined as the distance of the $x_{K}$ vector (values of the parameters) from the center $c_{K}$ (which is estimated as mean value), p is the order of the $x_{K}$ vector."
+ Then the aleorithm for determining the optimum number of clusters is as follows (Sugar James 2003)., Then the algorithm for determining the optimum number of clusters is as follows (Sugar James 2003).
+ Let us denote by d. the estimate of dy at the A point., Let us denote by $d_{K}^{\bf'}$ the estimate of $d_{K}$ at the $K^{th}$ point.
+ Then dis is the minimum achievable distortion associated with fitting IX. centers to the data., Then $d_{K}^{\bf'}$ is the minimum achievable distortion associated with fitting K centers to the data.
+ A natural wav of choosing the number of clusters is to plot dj. versus Ix and look for the resulting distortion curve., A natural way of choosing the number of clusters is to plot $d_{K}^{\bf'}$ versus K and look for the resulting distortion curve.
+ This curve is alwavs monotonic decreasing., This curve is always monotonic decreasing.
+ Initially one would expect much smaller drops ie. a levelling off for IX greater than the true number of clusters because past Chis point adding more centers simply partitions within groups rather than between groups., Initially one would expect much smaller drops i.e. a levelling off for K greater than the true number of clusters because past this point adding more centers simply partitions within groups rather than between groups.
+" According to Sugar James (2003) lor a large number of items the distortion curve when transformed lo an appropriate negative power (p/2 in our case). will exhibit a sharp “jump” (if we plot IX versus transIormed d,.)."," According to Sugar James (2003) for a large number of items the distortion curve when transformed to an appropriate negative power $p/2$ in our case), will exhibit a sharp ""jump"" (if we plot K versus transformed $d_{K}^{\bf'}$ )."
+ Then we have calculated the jumps in the transformed, Then we have calculated the jumps in the transformed
+"function of stars with 7-band flags ΑΟΕΕΟΕ gs approximated as If stars with 7-band flag ""H"" are also included. this becomes As outlined above.","function of stars with $I$ -band flags “A+C+E+F+G” is approximated as If stars with $I$ -band flag “H” are also included, this becomes As outlined above."
+ A-band magnitudes were to Hipparcos stars based on high-quality 2MASS and IRC data., $K$ -band magnitudes were to Hipparcos stars based on high-quality 2MASS and IRC data.
+ The completeness function is approximated as For fainter objects. a (reliable) /’-magnitude is available for almost all objects because of the all-sky nature of the 2MASS survey. but for brighter stars this is clearly not the case.," The completeness function is approximated as For fainter objects, a (reliable) $K$ -magnitude is available for almost all objects because of the all-sky nature of the 2MASS survey, but for brighter stars this is clearly not the case."
+ As outlined above. spectroscopic [Fe/H] values have been retrieved using Vizier and recent literature.," As outlined above, spectroscopic [Fe/H] values have been retrieved using Vizier and recent literature."
+ The completeness function is approximated as The error on the {Πρ magnitude is small but has nevertheless been taken into account., The completeness function is approximated as The error on the $Hp$ magnitude is small but has nevertheless been taken into account.
+ The error is Gaussian distributed with the sigma value As explained in the explanatory supplement to the Hipparcos catalogue. stars with an ecliptic latitude more than 47 degree were observed more often. and this resulted in lower parallax errors on average.," The error is Gaussian distributed with the sigma value As explained in the explanatory supplement to the Hipparcos catalogue, stars with an ecliptic latitude more than 47 degree were observed more often, and this resulted in lower parallax errors on average."
+ When 3}2 tthe error on the parallax (in mas) is in the mean with dispersion 0.075. and else with dispersion 0.063.," When $\beta >$ the error on the parallax (in mas) is in the mean with dispersion 0.075, and else with dispersion 0.063."
+ The minimum error on the parallax is In the simulation a parallax error is generated using the mean relation and the Gaussian dispersion., The minimum error on the parallax is In the simulation a parallax error is generated using the mean relation and the Gaussian dispersion.
+ If this parallax error is above the minimum allowed value the star is retained., If this parallax error is above the minimum allowed value the star is retained.
+" The ""measured"" parallax is based on the true parallax (from the true distance) and the Gaussian distributed parallax error.", The “measured” parallax is based on the true parallax (from the true distance) and the Gaussian distributed parallax error.
+ To begin with. à number of input parameters need to be set: The consecutive steps in the model are the following: Table 2 lists the observed absolute magnitudes depending on various selections.," To begin with, a number of input parameters need to be set: The consecutive steps in the model are the following: Table \ref{Tab-Res} lists the observed absolute magnitudes depending on various selections."
+ These selections largely follow the selections made by previous works às quoted in the Introduction., These selections largely follow the selections made by previous works as quoted in the Introduction.
+ In all cases. only stars where Hipparcos flag H30 (the goodness-of-fit) is «5. x>0. and single stars Gsoln = 5 as listed in the revised Hipparcos catalogue) are considered.," In all cases, only stars where Hipparcos flag H30 (the goodness-of-fit) is $<$ 5, $\pi >0$, and single stars (isoln = 5 as listed in the revised Hipparcos catalogue) are considered."
+ Following earlier work. stars are selected inside the box (all colours and magnitudes are understood to be de-reddened according to the model described in Sect. 3.3))," Following earlier work, stars are selected inside the box (all colours and magnitudes are understood to be de-reddened according to the model described in Sect. \ref{S-red}) )"
+" 0.75 <(V L25. 1.6«M, 1.2."," 0.75 $< (V-I) <$ 1.25, $-1.6 < M_{\rm I} <$ 1.2."
+ The data is binned in absolute, The data is binned in absolute
+"does not afford a better fit, although it shows that the properties of these two stacks follow closely the ones of the full sample.","does not afford a better fit, although it shows that the properties of these two stacks follow closely the ones of the full sample."
+" The SED fitting parameters are correlated with each other, and in particular, there is a well-know age-metallicity relation that can be seen easily as the axis of the 2-D credible region in the age-metallicity plane of Fig. 3.."," The SED fitting parameters are correlated with each other, and in particular, there is a well-know age-metallicity relation that can be seen easily as the axis of the 2-D credible region in the age-metallicity plane of Fig. \ref{fig:Prob}."
+ Our estimates of metallicity depend on the SPS templates we use and (although weakly) on the flat prior on log that we assume., Our estimates of metallicity depend on the SPS templates we use and (although weakly) on the flat prior on $\log Z$ that we assume.
+" We think that the latter is physically Zmotivated, and we have no reason not to trust the models; however, it is possible that the observed difference in metallicity between LAEs at the two redshift is overestimated."," We think that the latter is physically motivated, and we have no reason not to trust the models; however, it is possible that the observed difference in metallicity between LAEs at the two redshift is overestimated."
+" If the true metallicity at z=3.1 were higher, and the true metallicity at z=2.1 were lower, this might help reconcile the gap in the physical properties of these galaxies."," If the true metallicity at $z = 3.1$ were higher, and the true metallicity at $z = 2.1$ were lower, this might help reconcile the gap in the physical properties of these galaxies."
+" To test this hypothesis, we run our SED fitting code on the full samples at the two redshifts, holding the metallicity fixed at Z=0.2Z5; this value is suggested by the spectroscopic analysis of three LAEs at z~2.3 by Finkelsteinetal.(2011c).."," To test this hypothesis, we run our SED fitting code on the full samples at the two redshifts, holding the metallicity fixed at $Z = 0.2 Z_\odot$; this value is suggested by the spectroscopic analysis of three LAEs at $z \sim 2.3$ by \cite{2011ApJ...729..140F}."
+ Results are shown in Fig., Results are shown in Fig.
+ 4 and listed in Table 1.., \ref{fig:Syst} and listed in Table \ref{tab:resSED}.
+" As expected, the estimates for the mean ages of the stellar population shift in the desired direction, but remain incompatible at several sigma level."," As expected, the estimates for the mean ages of the stellar population shift in the desired direction, but remain incompatible at several sigma level."
+" This result seems to exclude metallicity as the primary source of the gap in the age of LAEs, at least in this stacking analysis."," This result seems to exclude metallicity as the primary source of the gap in the age of LAEs, at least in this stacking analysis."
+ The assumption of a particular functional form for the star formation history of LAEs can also influence the determination of ages., The assumption of a particular functional form for the star formation history of LAEs can also influence the determination of ages.
+" We investigate this issue by repeating the SED fitting procedure on the full samples using exponentially increasing and decreasing star formation histories (SFHs), v(t)οςe**/7."," We investigate this issue by repeating the SED fitting procedure on the full samples using exponentially increasing and decreasing star formation histories (SFHs), $\psi(t) \propto e^{\pm t/\tau}$."
+" By sampling in 1/r, both increasing and decreasing SFHs are explored as part of a contiguous parameter space, as explained in Acquavivaetal."," By sampling in $1/\tau$, both increasing and decreasing SFHs are explored as part of a contiguous parameter space, as explained in \cite{2011ApJ...737...47A}."
+ This parametrization is also able to capture (2011)..constant or starburst SFHs for appropriate values of r/age., This parametrization is also able to capture constant or starburst SFHs for appropriate values of $\tau$ /age.
+ The results of this SED fitting run (where the metallicities were held fixed at the best-fit value for both samples) are shown in Fig. 4;, The results of this SED fitting run (where the metallicities were held fixed at the best-fit value for both samples) are shown in Fig. \ref{fig:Syst};
+ in both cases the estimates of ages do not change significantly., in both cases the estimates of ages do not change significantly.
+" The quality of the fit does not improve significantly in either case by allowing this larger range of star formation histories, as can be seen in Table 1.."," The quality of the fit does not improve significantly in either case by allowing this larger range of star formation histories, as can be seen in Table \ref{tab:resSED}."
+" The data favor a nearly constant SFH at both redshifts, with a best-fit value of T around 4 Gyr for both samples."," The data favor a nearly constant SFH at both redshifts, with a best-fit value of $\tau$ around 4 Gyr for both samples."
+" Gull had also performed the SED fitting using these SFHs for the z—2.1 sample, albeit with a different algorithm, and had found similar results."," Gu11 had also performed the SED fitting using these SFHs for the $z=2.1$ sample, albeit with a different algorithm, and had found similar results."
+ A possible explanation of the difference in the physical nature of LAEs at z=3.1 and z=2.1 might be the presence of stellar components of different ages., A possible explanation of the difference in the physical nature of LAEs at $z = 3.1$ and $z = 2.1$ might be the presence of stellar components of different ages.
+" We test this hypothesis by performing the SED fitting on the full samples using two stellar populations, fixing the age of the second stellar population at 1 Gyr, and introducing the stellar mass of the second stellar population as a free SED fitting parameter."," We test this hypothesis by performing the SED fitting on the full samples using two stellar populations, fixing the age of the second stellar population at 1 Gyr, and introducing the stellar mass of the second stellar population as a free SED fitting parameter."
+" We show the results of this test for a single-burst model that would solely contain old stars, although we have also considered constant or linearly increasing SFHs for the second stellar population, as suggested by some recent observations Finkelsteinetal.2011;Lee2010 and"," We show the results of this test for a single-burst model that would solely contain old stars, although we have also considered constant or linearly increasing SFHs for the second stellar population, as suggested by some recent observations \citealt{2011ApJ...742..108F, 2010PhDT.......132L} and"
+The redshift distributions of the NEC*. NC“NEC and samples are compared in 77.,"The redshift distributions of the NEC*, 8C-NEC and samples are compared in 7."
+ As can be seen there are no strong ce»ndenceies. Of the redshift distribution on the sample selection method., As can be seen there are no strong dependencies of the redshift distribution on the sample selection method.
+ The largest change is seen at .ο9) where three of the high-z CSS sources drop out of both the qos and the SC-NEC samples. anclone extended: z2» object (8C1736|650) comes in to both the NEC* and SC-NEC samples.," The largest change is seen at $z>2$, where three of the $z$ CSS sources drop out of both the NEC* and the 8C-NEC samples, andone extended $z>2$ object (8C1736+650) comes in to both the NEC* and 8C-NEC samples."
+ For comparison with high radio luminosity objects we have used similar selection criteria to define a bright sample (3C*, For comparison with high radio luminosity objects we have used similar selection criteria to define a bright sample (3C*).
+ This is described in the Appendix., This is described in the Appendix.
+"wr>| averaged power transferred to electrons. Q=[ή<—eE,f,>dy due to the presence of a an harmonic electric field fluctuation Ej=cos(Kjz—wr) Is This can easily be verified from Eq.(1) and for a Maxwellian distribution it is equivalent to Eq.(26).","$\omega t\gg1$ averaged power transferred to electrons, $Q=\int v_{\parallel} <-eE_{\parallel}f_{1}>dv_{\parallel}$ due to the presence of a an harmonic electric field fluctuation $E_{\parallel}= \cos(k_\parallel z -\omega t)$ is This can easily be verified from Eq.(1) and for a Maxwellian distribution it is equivalent to Eq.(26)."
+ Using the relation between Ej and B_. Q can finally be expressed in term of the magnetic energy. Up=|Be|/9π. The coefficient of proportionality between Q and Up. which has the dimension of the inverse of a time. is nothing else than the Landau damping rate.," Using the relation between $E_{\parallel}$ and $B_{\perp}$, $Q$ can finally be expressed in term of the magnetic energy, $U_{B}=\mid B^{2}_{\perp}\mid/8 \pi$, The coefficient of proportionality between $Q$ and $U_{B}$, which has the dimension of the inverse of a time, is nothing else than the Landau damping rate."
+ The Landau damping rate is now directly obtained. without any reference to its physical meaning. from the complex dispersion relation.," The Landau damping rate is now directly obtained, without any reference to its physical meaning, from the complex dispersion relation."
+ The kinetic dispersion relation is obtained from: VAj=ico/Ucit+ZolaEy. aZo(a)]£j and Ej=thy+LLiWAy/c.," The kinetic dispersion relation is obtained from: $\nabla_{\perp}^{2}A_{\parallel}=i\omega/(k_{\parallel}^{2}\lambda_{De}^{2}c)[1+\alpha Z_{0}(\alpha)]E_{\parallel}$, $\rho_{s}^{2}\nabla_{\perp}^{2}\phi=i/k_{\parallel}[1+\alpha Z_{0}(\alpha)]E_{\parallel}$ and $E_{\parallel}=-ik_{\parallel}\phi +i\omega A_{\parallel}/c$."
+ It is This is the general complex dispersion relation for the dispersive Alfven wave., It is This is the general complex dispersion relation for the dispersive Alfven wave.
+ In the limit a<I. it reads Its real part corresponds to the frequency of the kinetic Alfven wave which was also derived from fluid theory above.," In the limit $\alpha \ll 1$, it reads Its real part corresponds to the frequency of the kinetic Alfven wave which was also derived from fluid theory above."
+ Its imaginary part. which corresponds to the Landaudamping rate [see also Eq.(28)] and reads Most calculations above were finalized in the limit a«I. in the opposite limit of a>>| one obtains the frequency and damping rate of the inertial Alfven wave. which has its parallel electric field balanced by the electron inertia in Ohm’s law.," Its imaginary part, which corresponds to the Landaudamping rate [see also Eq.(28)] and reads Most calculations above were finalized in the limit $\alpha \ll1$, in the opposite limit of $\alpha \gg 1$ one obtains the frequency and damping rate of the inertial Alfven wave, which has its parallel electric field balanced by the electron inertia in Ohm's law."
+" For frequency w~ΛΙ. oV4/Ve. hence the kinetic Alfven wave regime corresponds to v4/v;;«I and the inertial Alfven wave regime to v4/v,,2»1."," For frequency $\omega \sim k_{\parallel}v_{A}$, $\alpha\sim v_{A}/v_{te}$, hence the kinetic Alfven wave regime corresponds to $v_{A}/v_{te}\ll 1$ and the inertial Alfven wave regime to $v_{A}/v_{te}\gg 1$."
+" In the following we continue to focus on the warm plasma regime corresponding to |>>B,nzfni.", In the following we continue to focus on the warm plasma regime corresponding to $1\gg \beta_{e}\gg m_{e}/m_{i}$.
+ Phase mixing of a shear Alfven wave packet can be considered in the framework of an eikonal description: with w=μι., Phase mixing of a shear Alfven wave packet can be considered in the framework of an eikonal description: with $\omega=\pm k_{\parallel}v_{A}$.
+ These are the characteristics of the equation In the latter equation e- are the amplitudes of the wave-packets corresponding to w=Kv and y(K) is a wave-number dependent damping rate., These are the characteristics of the wave-kinetic equation In the latter equation $e_{\pm}$ are the amplitudes of the wave-packets corresponding to $\omega=\pm k_{\parallel}v_{A}$ and $\gamma(\mathbf{k})$ is a wave-number dependent damping rate.
+ Following the trajectory of a in phase space (x.K). its amplitude evolves according to The latter equation is integrated to give The principle of phase-mixing is simple: for any damping rate y which ts an increasing function of & any mechanism producing an increase in & as a function of time results also ina smaller damping time scale.," Following the trajectory of a in phase space $(\mathbf{x},\mathbf{k})$, its amplitude evolves according to : The latter equation is integrated to give The principle of phase-mixing is simple: for any damping rate $\gamma$ which is an increasing function of $k$, any mechanism producing an increase in $k$ as a function of time results also in a smaller damping time scale."
+ This is precisely the situation when the Alfven wave packet propagates along field lines with a transverse variation of the Alfven speed: the wave packet is sheared., This is precisely the situation when the Alfven wave packet propagates along field lines with a transverse variation of the Alfven speed: the wave packet is sheared.
+ In this case. say γα)=—Vxz. z being the unit vector in the parallel direction and x the transverse coordinate. then ∖∖⇁≣⋔⇂↴⋝⇁∐⊜⊓∏⋯⋯∏⊽∙∕∖∶⋁⋅∖∕∠≓⋅∠≓⇂↴⊜⋯∶↔↾↾∣↴⊖∁∣⊺≏∐⋪∐∁↾⊖∣⋪≣⋋⊓∁ ∣⊜∏∶↔↾⋔⋂↑↴↾∣↴⊖∏⋪⇣⋯⋋∖⇁⊜∣⋪⋋⊜≣∏∣↴∪⋯⋂∶↔⊺⊖∏⊜⇈⋝⇁∐⋯↿∣∖⊳⊔∶∣∖⊳⊔≼↾∶⋯⊳⊤⋯⋋ ," In this case, say $\mathbf{v}_{A}(x)=-v_{A}'x\mathbf{z}$, $\mathbf{z}$ being the unit vector in the parallel direction and $x$ the transverse coordinate, then with by definition $
+v_{A}'=v_{A}/L_{\perp}$, $L_{\perp}$ being the characteristic length of the transverse inhomogeneity and $k_{\parallel}=k_{\parallel}(t=0)$."
+means that A_ increases linearly with time due to differential advection of the wave packets along the field lines. t.e. where we have taken κο=0)O without loss of generality.," This means that $k_{\perp}$ increases linearly with time due to differential advection of the wave packets along the field lines, i.e. where we have taken $k_{\perp}(t=0)=0$ without loss of generality."
+ For a resistive MHD Ohms? law. Ej=Jj the following results are well known.," For a resistive MHD Ohms' law, $E_{\parallel}=\eta J_{\parallel}$ the following results are well known."
+ The damping rate is y(k)=nee+kin-DG+ ky). this is the fourier transform of the operator responsible for magnetic diffusion in the induction equation.," The damping rate is $\gamma(\mathbf{k})=\eta c(k_{\perp}^{2}+k_{\parallel}^{2})/4\pi=D_{m}(k_{\perp}^{2}+k_{\parallel}^{2})$ , this is the fourier transform of the operator responsible for magnetic diffusion in the induction equation."
+ Hence which in the limit ¢>>vl yields Since. z=v4f. we also2 have for an Alfven wave excited at z=0 with frequency ω.," Hence which in the limit $t\gg v_{A}'^{-1}$ yields Since, $z=v_{A}t$, we also have for an Alfven wave excited at $z=0$ with frequency $\omega$."
+ In a collisionless plasma. when the dissipation is provided by electron Landau damping. with damping rate Vivkkpy /2vie. the equivalent expressions are: and also for an Alfven wave excited at z=0 with frequency c.," In a collisionless plasma, when the dissipation is provided by electron Landau damping, with damping rate $\gamma(\mathbf{k})=\sqrt{\pi}v^{2}_{A}k_{\parallel}k_{\perp}^{2}\rho_{s}^{2}/2v_{te}$ , the equivalent expressions are: and also for an Alfven wave excited at $z=0$ with frequency $\omega$ ."
+ Hence. the phase mixing time scale is and the phase mixinguu length scale isNotice that the scaling of the phase mixing length scale with the frequency c» in. the. spatial problem is different from that of resistive MHD phase mixing since," Hence, the phase mixing time scale is and the phase mixing length scale isNotice that the scaling of the phase mixing length scale with the frequency $\omega$ in the spatial problem is different from that of resistive MHD phase mixing since"
+Gamma-ray bursts (GRBs) are short and intense pulses of soft -rays.,Gamma-ray bursts (GRBs) are short and intense pulses of soft $\gamma$ -rays.
+ In this work we study their and their cosmic rate., In this work we study their and their cosmic rate.
+ These functions are essential to understand the nature of GRBs and to determine their progenitors., These functions are essential to understand the nature of GRBs and to determine their progenitors.
+ They may also shed light on the still mysterious physics of the central engine., They may also shed light on the still mysterious physics of the central engine.
+ The and the cosmic GRB rate are observationally entangled as the observed rate is a convolution of the with the cosmic rate., The and the cosmic GRB rate are observationally entangled as the observed rate is a convolution of the with the cosmic rate.
+" For this reason, almost every work before this paper have made an a-priori assumptions on at least one of these functions."," For this reason, almost every work before this paper have made an a-priori assumptions on at least one of these functions."
+ At first - having no motivation to believe otherwise - the rate was assumed to be constant., At first - having no motivation to believe otherwise - the rate was assumed to be constant.
+ The simplest form for the was a standard candle i.c. a constant luminosity., The simplest form for the was a standard candle i.e. a constant luminosity.
+ These, These
+as the number of fields is increased but the window ellects Es»oil this.,as the number of fields is increased but the window effects spoil this.
+ Phe window used is in theory circularlv svmmoetric out as it is on a lattice it necessarily has some scquareness in un., The window used is in theory circularly symmetric but as it is on a lattice it necessarily has some squareness in it.
+ Phe important point is that we get different results for 1ο Gaussian and fields., The important point is that we get different results for the Gaussian and fields.
+ Now we test the statistic on fields with many squares., Now we test the statistic on fields with many squares.
+" vc""ugures 5 and 6 show the Ist and 2nd Fourier transforms for Lr) and. 100 squares in one field.", Figures \ref{10square1} and \ref{100square1} show the 1st and 2nd Fourier transforms for 10 and 100 squares in one field.
+ We see that as the number of squares increases the result is not so distinctive., We see that as the number of squares increases the result is not so distinctive.
+ For 100 squares the signal is barely distinguishable from that of a Gaussian(see Figure 3))., For 100 squares the signal is barely distinguishable from that of a Gaussian(see Figure \ref{gauss1}) ).
+ In fact in real space it is obviously not Gaussian. suggesting this is not the best way of detecting non-Gaussianity.," In fact in real space it is obviously not Gaussian, suggesting this is not the best way of detecting non-Gaussianity."
+ Figure 7 shows the field in real space., Figure \ref{pic100} shows the field in real space.
+ The important test however. is to see if we can detect when it is superimposed on a Gaussian background so it is not visible in real space.," The important test however, is to see if we can detect when it is superimposed on a Gaussian background so it is not visible in real space."
+ Figure S shows a real field which is the sum of a field containing one square and a Ciaussian field., Figure \ref{pic} shows a real field which is the sum of a field containing one square and a Gaussian field.
+" ""Phough in real space we would not guess there was present. in fact in Fourier space many rings look cillerent [rom those produced by Gaussian fields."," Though in real space we would not guess there was present, in fact in Fourier space many rings look different from those produced by Gaussian fields."
+ An example. of one is shown in Figure 9.., An example of one is shown in Figure \ref{gsq10}.
+ This is the result of averaging over 10 such fields. as in Section ??.. this time with a," This is the result of averaging over 10 such fields, as in Section \ref{small}, , this time with a"
+hieher. ancl the number of particles residing in dense clumps is larger al higher resolutions.,"higher, and the number of particles residing in dense clumps is larger at higher resolutions."
+ For the AB runs. the PDFs from the 512? resolution run is almost identical to higher resolution runs. indicating convergence.," For the AB runs, the PDFs from the $512^2$ resolution run is almost identical to higher resolution runs, indicating convergence."
+ For the BA runs. the small number of dense clumps in the simulation box makes convergence more difficult bv averaging over a finite (me period.," For the BA runs, the small number of dense clumps in the simulation box makes convergence more difficult by averaging over a finite time period."
+ Nevertheless. viewing from the PDF curves. it appears that convergence is finally reached αἱ 024? resolution.," Nevertheless, viewing from the PDF curves, it appears that convergence is finally reached at $1024^2$ resolution."
+ We have also studied the nunerical convergence wilh number of particles in the simulation box., We have also studied the numerical convergence with number of particles in the simulation box.
+ On the right. panels ol Figure 6.. we see that [or both runs AB and BA. the PDFs depend verv weakly on Ny.," On the right panels of Figure \ref{fig:pdf}, we see that for both runs AB and BA, the PDFs depend very weakly on $N_{\rm pc}$."
+ In the AB runs. the five curves almost overlap with each other.," In the AB runs, the five curves almost overlap with each other."
+ In BA runs. however. there is larger scatter.," In BA runs, however, there is larger scatter."
+ Again. such scatter is most likely due to the small number of chimps.," Again, such scatter is most likely due to the small number of clumps."
+" In fact. we have compared the BA run PDFs between short (T.—T,=5000— !) and long (7;—T;=12000!. shown in the ligure) averaging period."," In fact, we have compared the BA run PDFs between short $T_e-T_s=500\Omega^{-1}$ ) and long $T_e-T_s=1200\Omega^{-1}$, shown in the figure) averaging period."
+ The PDFs from averaging converges show betterconvergence!., The PDFs from longer averaging converges show better.
+. We expect the scatter to become smaller if we run BA series for much longer time., We expect the scatter to become smaller if we run BA series for much longer time.
+ In all. we conclude that one can use as small as only one particle per cell to accurately capture the density. distributions of (he solids.," In all, we conclude that one can use as small as only one particle per cell to accurately capture the density distributions of the solids."
+ The streaming instability generates (turbulence which dilfers significantly from (he laminar state described bv the NSIT equilibrium., The streaming instability generates turbulence which differs significantly from the laminar state described by the NSH equilibrium.
+ The turbulence is accompanied by particle concentration. as discussed in the previous subsection.," The turbulence is accompanied by particle concentration, as discussed in the previous subsection."
+ Understanding the turbulent state is important since ib mav stronglv affect the settling process as well as the racial transport of solids in the PPDs (JYOT. Dai&Stone 2010a)).," Understanding the turbulent state is important since it may strongly affect the settling process as well as the radial transport of solids in the PPDs (JY07, \citealp{BaiStone10b}) )."
+" Similar to JYOT. we have computed a number of physical «quantities to characterize (he turbulence properties generated [rom the streaming instability. including: 1) The turbulent. Mach number CM) in three directions. calculated. [rom the rool mean square of gas velocity [Inctuations: 2) Mean radial dift velocity of the particles. normalized by the NSIL radial drift velocitym77/77"": 3) Turbulent diffusion coefficient [or parücles D. caleulatedusing the method in 855.3 ofJYOT: and +) The time and spatial averaged Revnolds stress F=pju,(u,—ty). divided bv the Revnolds stress in the NSII equilibrium (note the quantity we measure is different [rom Fe, in JYOT)."," Similar to JY07, we have computed a number of physical quantities to characterize the turbulence properties generated from the streaming instability, including: 1) The turbulent Mach number $Ma$ ) in three directions, calculated from the root mean square of gas velocity fluctuations; 2) Mean radial drift velocity of the particles, normalized by the NSH radial drift velocity $\overline{v_x}/\overline{v_x}^{NSH}$; 3) Turbulent diffusion coefficient for particles $D$, calculated using the method outlined in 5.3 ofJY07; and 4) The time and spatial averaged Reynolds stress ${\mathcal F}\equiv\rho_gu_x(u_y-v_K)$, divided by the Reynolds stress in the NSH equilibrium (note the quantity we measure is different from ${\mathcal F}_{\mathcal{L},x}$ in JY07)."
+ Properties (1). (2) and (4) are calculated. directly. [rom spatial and (ime averaging [rom (he entire simulation box.," Properties (1), (2) and (4) are calculated directly from spatial and time averaging from the entire simulation box."
+ The lo uncertainties are estimated from the standard deviation in the time sequence., The $1\sigma$ uncertainties are estimated from the standard deviation in the time sequence.
+" In the calculation of the diffusion coefficient. we measure (he standard deviation of the distance traveled bv tracer particles al different time intervals 2,(M).0.(M)."," In the calculation of the diffusion coefficient, we measure the standard deviation of the distance traveled by tracer particles at different time intervals $\sigma_x(\Delta t),
+\sigma_z(\Delta t)$."
+ We then lit (he diffusion coefficient using o2.0M)ο M.," We then fit the diffusion coefficient using $\sigma^2_{x,z}(\Delta t)=2D_{x,z}\Delta t$ ."
+ The maxinnm A? has been chosen to, The maximum $\Delta t$ has been chosen to
+each galaxy for o5; ave listed in column 4.,each galaxy for $\sigma_{\rm BG}^{2}$ are listed in column 4.
+ The error bars of oj; have been computed taking into account only the n() fitting uncertainty., The error bars of $\sigma_{\rm BG}^{2}$ have been computed taking into account only the $n(m)$ fitting uncertainty.
+" Note that in Marin-Franch the error bar of 2c; was determined considering possible variations in the slope of the faint end of the nti) (rom 5.=0.30 to 5=0.40). but in the present studs. these possible varlalions are being considered explicitly by studying the cases 5=0.30. 0.34. and 0.39 separately,"," Note that in \citet{MA02} the error bar of $\sigma_{\rm BG}^{2}$ was determined considering possible variations in the slope of the faint end of the $n(m)$ (from $\gamma=0.30$ to $\gamma=0.40$ ), but in the present study, these possible variations are being considered explicitly by studying the cases $\gamma=0.30$, $0.34$, and $0.39$ separately."
+ Equation Ὁ can now be used to derive ajc for each region., Equation \ref{sigmaIGC} can now be used to derive $\sigma_{\rm IGC}^2$ for each region.
+ The results are given in column 5 of Table 2.., The results are given in column 5 of Table \ref{igc}.
+ The number population of IGC. Nic. can be estimated from eje using the σιιός ratios listed column 6.," The number population of IGC, $N_{\rm IGC}$ , can be estimated from $\sigma_{\rm IGC}^2$ using the $\sigma^{2}_{\rm IGC}/N_{\rm IGC}$ ratios listed column 6."
+ These ratios were computed in (2002)., These ratios were computed in \citet{MA02}.
+. The resulting Nyce ds listed in column 7., The resulting $N_{\rm IGC}$ is listed in column 7.
+" Note that Nige is expresed in IGCs/("")?. while the rest of columns are expresed in pirels."," Note that $N_{\rm IGC}$ is expresed in $IGCs/(\arcsec)^2$, while the rest of columns are expresed in $pixels$."
+ The og; uncertainties have been computed adding in quadrature the uncertainties of 2) and o5., The $\sigma_{\rm IGC}^2$ uncertainties have been computed adding in quadrature the uncertainties of $P_0$ and $\sigma_{\rm BG}^2$.
+ On the other hand. the uncertainties of the ratio oj:;0/.Vic have been neglected.," On the other hand, the uncertainties of the ratio $\sigma^{2}_{\rm IGC}/N_{\rm IGC}$ have been neglected."
+ The reason is that y is assumed to be universal (we assume the same 4? parameters adopted in n-Franch Aparicio 2002. 0=1.40 and m= 27.42).," The reason is that $\varphi$ is assumed to be universal (we assume the same $\varphi$ parameters adopted in n-Franch Aparicio 2002, $\sigma=1.40$ and $m_R^0=27.42$ )."
+ In this sense. the uncertainty in the 2 parameters would introduce a rescaling factor common to all the results. but the relative errors would not change aud the prools of IGCs detection would remain valid.," In this sense, the uncertainty in the $\varphi$ parameters would introduce a rescaling factor common to all the results, but the relative errors would not change and the proofs of IGCs detection would remain valid."
+" The origin of the galaxv-to-galaxy variations in the ratio oj,o/.Njoc relies on the differences of the photometric calibration constants.", The origin of the galaxy-to-galaxy variations in the ratio $\sigma^{2}_{\rm IGC}/N_{\rm IGC}$ relies on the differences of the photometric calibration constants.
+ The values obtained for Vie1 lor the field regions5 associatedwith the 16 5galaxies studied in Marfn-Franch&Aparicio(2002) (dots) and for four additional intergalactic field regions (stars) are plotted in Fig., The values obtained for $N_{\rm IGC}$ for the field regions associatedwith the 16 galaxies studied in \cite{MA02} (dots) and for four additional intergalactic field regions (stars) are plotted in Fig.
+ 3. versus the distance FR to the Coma galaxy cluster center., \ref{gamma} versus the distance $R$ to the Coma galaxy cluster center.
+ Each panel corresponds (o à value of 5., Each panel corresponds to a value of $\gamma$.
+ Regions associated will galaxies were selected in Alarin- witha the aimal of ossamplingT the background., Regions associated with galaxies were selected in \cite{MA02} with the aim of sampling the background.
+ack eThev unare bevondK the GCS of each galaxv and they. consequently sample the intergalactic [field population., They are beyond the GCS of each galaxy and they consequently sample the intergalactic field population.
+ The four intergalactic5 “blank” regionsS are far [rom any 5galaxy., The four intergalactic “blank” regions are far from any galaxy.
+ The fact that the averagehel of Nye1ος for these regionsoO is the same as for the regionse associated with egalaxies confirms that the fielcl is really being samplecl in both cases., The fact that the average of $N_{\rm IGC}$ for these regions is the same as for the regions associated with galaxies confirms that the field is really being sampled in both cases.
+ The most significant. characteristic of the plots in Fie., The most significant characteristic of the plots in Fig.
+ 3. is the flat. distributions οἱ values., \ref{gamma} is the flat distributions of values.
+ The only point perhaps departing from this trend is (hat associated with the central ealaxv. NGC 4874.," The only point perhaps departing from this trend is that associated with the central galaxy, NGC 4874."
+ Furthermore. it can be seen (hat Ngoc is negative in almost all cases lor = 0.39. indicating that this value is unlikely.," Furthermore, it can be seen that $N_{\rm IGC}$ is negative in almost all cases for $\gamma=0.39$ , indicating that this value is unlikely."
+ On the other hand. if 4= 0.30. then /Nqcc is Clearly a positive number. implving that an IGC population shouldexist.," On the other hand, if $\gamma=0.30$ , then $N_{\rm IGC}$ is clearly a positive number, implying that an IGC population shouldexist."
+ This population, This population
+Our scheme is stable if «7ο0.,Our scheme is stable if $\omega^2 \ge 0$.
+ Lt is also desirable for the numerical sound. speed to equal the true sound. speed: w.2 ," It is also desirable for the numerical sound speed to equal the true sound speed: $\omega^2/k^2
+= c_s^2$."
+In SPILL it is typical to use the cubic spline (C'S) kernel eiven by: where .=rfh is the distance from the centre of the kernel in units of the smoothing length., In SPH it is typical to use the cubic spline (CS) kernel given by: where $x=r/h$ is the distance from the centre of the kernel in units of the smoothing length.
+ In Figure L.. we show contours of ος as a function ofwavenumber & and the smoothing length  in units of the inter-particle spacing dr—1. for the CS kernel.," In Figure \ref{fig:stab}, we show contours of $\omega^2 / k^2 / c_s^2$ as a function ofwavenumber $k$ and the smoothing length $h$ in units of the inter-particle spacing $dx=1$, for the CS kernel."
+ We assume an acliabatic equation of state with 7?= 5/3., We assume an adiabatic equation of state with $\gamma = 5/3$ .
+ From left to right the plots show (A.ky.k;)=KG.0:0)&(1.:1.0) and kC1.1.1).," From left to right the plots show $(k_x,k_y,k_z) = k(1,0,0), k(1,1,0)$ and $k(1,1,1)$."
+ The three rows show the longitudinal wave and the two transverse waves for these orientations., The three rows show the longitudinal wave and the two transverse waves for these orientations.
+ From Figure 1. it is clear that the CS kernel in 3D is unstable to longitudinal waves for f22. and very unstable to transverse waves.," From Figure \ref{fig:stab}, it is clear that the CS kernel in 3D is unstable to longitudinal waves for $h \simgt 2$, and very unstable to transverse waves."
+ The unstable longitucinal waves drive the clumping. or tensile instability Chat causes particles to chump on the kernel scale (??:: ?5 Tu 7 ? ?))," The unstable longitudinal waves drive the clumping, or tensile instability that causes particles to clump on the kernel scale \bcite{1981A&A....97..373S}; \bcite{1992MNRAS.257...11T}; \bcite{1994MmSAI..65.1013H}; \bcite{1996PASA...13...97M}; \bcite{2000Monaghan}) )."
+ The clumping instability is à problem because it means that increasing the neighbour number will not give improved sampling of the kernel. and the Eo error will remain large.," The clumping instability is a problem because it means that increasing the neighbour number will not give improved sampling of the kernel, and the $\uE$ error will remain large."
+ Llowever. the situation is dramatically improved if. we add a constant central core to the kernel gradient D.," However, the situation is dramatically improved if we add a constant central core to the kernel gradient $\frac{\partial W}{\partial r}$."
+ ‘This eives a constant force term at the centre of the Ixernel that physically prevents clumping., This gives a constant force term at the centre of the Kernel that physically prevents clumping.
+ We choose a kernel that is maximally similar to theCS kernel. while obeving const. Vr<a. where à is the core size.," We choose a kernel that is maximally similar to theCS kernel, while obeying $\frac{\partial W}{\partial r} = \mathrm{const.}$ $\forall r < \alpha$, where $\alpha$ is the core size."
+ This leads us to Aethe Core Triangle (CL) kernel: where j1l|6a?a. ON=S/r(G4a%6a’| I)]. and the core size is fixed at a=1/3 by the requirement that jS. be continuous.," This leads us to the Core Triangle (CT) kernel: where $\beta = 1 + 6\alpha^2 - 12\alpha^3$, $N
+= 8/[\pi \left(6.4\alpha^5 - 16\alpha^6 + 
+ 1\right)]$ , and the core size is fixed at $\alpha = 1/3$ by the requirement that $\frac{\partial^2 W}{\partial r^2}$ be continuous."
+. Figure 2. shows stability plots for the €T. kernel., Figure \ref{fig:stabct} shows stability plots for the CT kernel.
+ The C'T kernel has greatly improved stability for the longitucinal waves (top row) compared to the CS kernel ancl should eive significantly improved. performance for large neighbour number., The CT kernel has greatly improved stability for the longitudinal waves (top row) compared to the CS kernel and should give significantly improved performance for large neighbour number.
+ We demonstrate this in refsee:kh.., We demonstrate this in \\ref{sec:kh}.
+ Note that for all of the kernels we use in this paper. we consistently apply the kernel for the density estimate and its eracient for the energy and momentum equations.," Note that for all of the kernels we use in this paper, we consistently apply the kernel for the density estimate and its gradient for the energy and momentum equations."
+ For small neighbour number. the central triangle in the CP kernel will degrade the quality of the density estimate.," For small neighbour number, the central triangle in the CT kernel will degrade the quality of the density estimate."
+ Eowever. in this paper we tvpically use large neighbour numbers (2100).," However, in this paper we typically use large neighbour numbers $>100$ )."
+ In this case. very few particles sample the inner regions of the kernel ancl the bias introduced in the density is negligible. (," In this case, very few particles sample the inner regions of the kernel and the bias introduced in the density is negligible. ("
+The quality of the density estimate in ccan be seen in the Sod shock tube test in refsec:socl..),The quality of the density estimate in can be seen in the Sod shock tube test in \\ref{sec:sod}. .)
+ We found in tests that retaining the CS kernel just for the density estimate gives near-identical results., We found in tests that retaining the CS kernel just for the density estimate gives near-identical results.
+ The clumping instabilitv is a result of unstable ongitudinal waves., The clumping instability is a result of unstable longitudinal waves.
+ (X related. instability the. banding instability is a result. of unstable transverse. waves., A related instability – the banding instability – is a result of unstable transverse waves.
+ For Wh the CP and CS kernels. there are broad. bands of instability to transverse waves (see Figures 1. ancl 2)).," For both the CT and CS kernels, there are broad bands of instability to transverse waves (see Figures \ref{fig:stab} and \ref{fig:stabct}) )."
+ Lf he neighbour number is carefully chosen to lic in a stable region. banding will not occur.," If the neighbour number is carefully chosen to lie in a stable region, banding will not occur."
+ Llowever. banding can still x excited at boundaries i£ h changes there. moving into an unstable region.," However, banding can still be excited at boundaries if $h$ changes there, moving into an unstable region."
+ For both the CS and C'T kernels. it is ciflicult to find a suitable neighbour number for which the kernel is stable to all transverse and longitudinal modes.," For both the CS and CT kernels, it is difficult to find a suitable neighbour number for which the kernel is stable to all transverse and longitudinal modes."
+ This suggests hunting for an even more stable kernel., This suggests hunting for an even more stable kernel.
+ A full search is. beyond the scope of this paper., A full search is beyond the scope of this paper.
+ Here. we present a simple class of kernels that improve stability by moving to higher order (?)..," Here, we present a simple class of kernels that improve stability by moving to higher order \citep{1996PhDMorris}."
+" Following ?.. we generalise our CL kernel to order n; to obtain the Following class of kernels that we call the Ligh Order Core-Friangle (HOC""D) kernels: where: and a and No are calculated: numerically for a givenchoice of ni."," Following \citet{2005astro.ph..7472P}, we generalise our CT kernel to order $n_k$ to obtain the following class of kernels that we call the High Order Core-Triangle (HOCT) kernels: where: and $\alpha$ and $N$ are calculated numerically for a givenchoice of $n_k$."
+ Continuity requires that à solves the equation: where ;J and ~ are [ree parameters., Continuity requires that $\alpha$ solves the equation: where $\beta$ and $\gamma$ are free parameters.
+ In this paper we choose 3=0.5.40.75.," In this paper we choose $\beta = 0.5, \gamma = 0.75$."
+ Other choices. and indeed other high-order kernels. may give better results than those presented here.," Other choices, and indeed other high-order kernels, may give better results than those presented here."
+" We tabulate values for 4A.D.P.se and UN as à function n, in Table 1. "," We tabulate values for $A,B,P,Q,\alpha$ and $N$ as a function $n_k$ in Table \ref{tab:hoctkern}. ."
+Notice that the coresize à decreases with nj., Notice that the coresize $\alpha$ decreases with $n_k$ .
+" Stability plots for the HOC kernel (with m,= 4) are given in Figure 3..", Stability plots for the HOCT4 kernel (with $n_k = 4$ ) are given in Figure \ref{fig:stabhoct}. .
+ Notice the improvement over the ς kernel. particularly for the transverse waves.," Notice the improvement over the CT kernel, particularly for the transverse waves."
+ There are two bands where the kernel is fully stable to both longitudinal, There are two bands where the kernel is fully stable to both longitudinal
+Goodman οἱ al.,Goodman et al.
+ 1905. Caselli οἱ al.," 1998, Caselli et al."
+ 2002). small rotational velocity gradients (e.g.. Goodman el al.," 2002), small rotational velocity gradients (e.g., Goodman et al."
+ 1993. Caselli et al.," 1993, Caselli et al."
+ 2002). and subsonic infall motions (e.g.. Tafalla et al.," 2002), and subsonic infall motions (e.g., Tafalla et al."
+ 1998. Lee. Alvers ancl Tafalla 2001). the turbulence alone cannot be responsible for the formation of LMCs.," 1998, Lee, Myers and Tafalla 2001), the turbulence alone cannot be responsible for the formation of LMCs."
+ Furthermore. the associated thermal pressure in the cores dominates the turbulent component by a factor of several (e.g.. Tafalla et al.," Furthermore, the associated thermal pressure in the cores dominates the turbulent component by a factor of several (e.g., Tafalla et al."
+ 2004)., 2004).
+ As another mechanism [for formation of LAICs. we can refer to the work of van Loo. Falle and Hartquist (2007) who examined (he effect of MIID waves on dense cores.," As another mechanism for formation of LMCs, we can refer to the work of van Loo, Falle and Hartquist (2007) who examined the effect of MHD waves on dense cores."
+ They found that short-wavelength waves can plav an important role in the generation of LMCs within a core without. breaking it., They found that short-wavelength waves can play an important role in the generation of LMCs within a core without breaking it.
+ llere. we can encourage ourselves to authenticate the suggestion that the thermal instability (TI) may be a possible favorable scenario for formation of LMCs.," Here, we can encourage ourselves to authenticate the suggestion that the thermal instability (TI) may be a possible favorable scenario for formation of LMCs."
+ This idea can also be supported bv the probes in temperature of the molecular cores in which (he observations show there may be a deviation from uniform isothermal case (e.g.. Harju et al.," This idea can also be supported by the probes in temperature of the molecular cores in which the observations show there may be a deviation from uniform isothermal case (e.g., Harju et al."
+ 2008. Friesen el al.," 2008, Friesen et al."
+ 2009)., 2009).
+ After the pioneer work of Field (1965). there are many. works on development οἱ TI in (he interstellar medium. in both numerical anc analvtical approaches.," After the pioneer work of Field (1965), there are many works on development of TI in the interstellar medium, in both numerical and analytical approaches."
+ The main idea is that the rapid growth of TI leads to a strong density imbalance between the cold dense inhomogeneitv and its low-densitv environment., The main idea is that the rapid growth of TI leads to a strong density imbalance between the cold dense inhomogeneity and its low-density environment.
+ In the molecular clouds. Gilden (1984) caleulated the cooling function to show that molecular gas may. be thermally unstable in environment where CO cooling dominates.," In the molecular clouds, Gilden (1984) calculated the cooling function to show that molecular gas may be thermally unstable in environment where $\mathrm{CO}$ cooling dominates."
+ The numerical calculations of Falle. Ager and Tlartquist (2006) eive the idea that TI may have an important role in formation of internal substructures in the molecular clouds.," The numerical calculations of Falle, Ager and Hartquist (2006) give the idea that TI may have an important role in formation of internal substructures in the molecular clouds."
+ Fukue aud Ixamava (2007) considered the effect of [rietion of two fIuids on the growth of TI., Fukue and Kamaya (2007) considered the effect of ion-neutral friction of two fluids on the growth of TI.
+ They found the TI of the weakly ionized plasma in the magnetic field. even at (he small length scales. could be erown up.," They found the TI of the weakly ionized plasma in the magnetic field, even at the small length scales, could be grown up."
+ Nejad- (2007) applied this idea for a one-dimensional sell-eravitating magnetized molecular slab including the frictional heating due to ion-neutral drift., Nejad-Asghar (2007) applied this idea for a one-dimensional self-gravitating magnetized molecular slab including the frictional heating due to ion-neutral drift.
+ Ie found that the heating of anmbipolar diffusion in outer regions of the slab is more signilicant Chan the average heating rates of cosmic ravs and turbulent motions. tus. isobaric TI could take place in (his area.," He found that the heating of ambipolar diffusion in outer regions of the slab is more significant than the average heating rates of cosmic rays and turbulent motions, thus, isobaric TI could take place in this area."
+ Nejad-Asghar and Molteni (2003) used the (wvo-fIuid smoothed particle hvdrodynanmies to simulate a partially ionized one-dimensional cloud., Nejad-Asghar and Molteni (2008) used the two-fluid smoothed particle hydrodynamics to simulate a partially ionized one-dimensional cloud.
+" Their eained results indicated that the Tl can insist on (he occurrence of density contrast al outer parts of the molecular cloud cores,", Their gained results indicated that the TI can insist on the occurrence of density contrast at outer parts of the molecular cloud cores.
+ Also. the two-dimensional simulations of Nejad-Asghar and Soltani (2009) confirmed (his idea that heating of ambipolar diffusion in the molecular cloud. cores may lead to TI and formation of LMCs.," Also, the two-dimensional simulations of Nejad-Asghar and Soltani (2009) confirmed this idea that heating of ambipolar diffusion in the molecular cloud cores may lead to TI and formation of LMCs."
+ In this paper. we (vv to make an extension to (his idea that TI may be responsible for the formation of LAICs within the molecular cloud cores.," In this paper, we try to make an extension to this idea that TI may be responsible for the formation of LMCs within the molecular cloud cores."
+ For this purpose. we consider a cvlindrical contracting. cloud. as a background. unperturbed. arrange. wilh a power-law drop in its densitv-profile.," For this purpose, we consider a cylindrical contracting cloud as a background unperturbed arrange, with a power-law drop in its density-profile."
+ In section 2. the net cooling Iunction in the molecular clouds is," In section 2, the net cooling function in the molecular clouds is"
+parameters.,parameters.
+" The parameters are energychosen from the WMAP7 CMB+BAO+HSTenergy best-fit, Nph?=0.0226,0.0?0.1123,0.704,0.963, as well as other parameters such as theh scalar n,amplitude A, and the reionization optical depth 7 (Komatsuefal.2009)."," The parameters are chosen from the WMAP7 CMB+BAO+HST best-fit, $\ombh = 0.0226, \omch= 0.1123, h =
+0.704, n_s =0.963$, as well as other parameters such as the scalar amplitude $A_s$ and the reionization optical depth $\tau$ \citep{wmap7}."
+". We modified the publicly available COSMOMC code (Lewis&Bridle2002) for generating the transfer functions, and therefore the power spectrum at different redshifts for these models."," We modified the publicly available COSMOMC code \cite{cosmomc} for generating the transfer functions, and therefore the power spectrum at different redshifts for these models."
+" Figure 1 shows the behavior of the equation of state ofenergy,, and the matter power spectrum at redshift z=0 for these two models."," Figure \ref{fig:ede} shows the behavior of the equation of state of, and the matter power spectrum at redshift $z=0$ for these two models."
+ For comparison we also plot the results for the matter power spectrum when perturbations are not considered., For comparison we also plot the results for the matter power spectrum when perturbations are not considered.
+" We see that when the perturbations are not considered, the EDE matter power spectrum, although there is a suppression the matter this effect is negligible, and for both models powerthe matter spectrum,power spectrum is very close to that for ACDM."," We see that when the perturbations are not considered, the EDE matter power spectrum, although there is a suppression the matter power spectrum, this effect is negligible, and for both models the matter power spectrum is very close to that for $\ld$ CDM."
+" However, when the perturbations are considered, the matter power spectrum is significantly different from ACDM, resulting in a og=0.74 for EDEI, and a og=0.76 for EDE2, whereas the ACDM has a og=0.81."," However, when the perturbations are considered, the matter power spectrum is significantly different from $\ld$ CDM, resulting in a $\sig = 0.74$ for EDE1, and a $\sig = 0.76$ for EDE2, whereas the $\ld$ CDM has a $\sig =0.81$."
+ This behavior is typical of EDE models with rapid and large transition of the equation of state., This behavior is typical of EDE models with rapid and large transition of the equation of state.
+" This translates to a rapid change in the at high scales (low k), thus to enhancement in the transfer function at low k."," This translates to a rapid change in the at high scales (low $k$ ), thus to enhancement in the transfer function at low $k$."
+" When normalized to the CMB scale, this effect shows as an of power at high k, and thereforeup a low og."," When normalized to the CMB scale, this effect shows up as an apparent suppression of power at high $k$, and therefore a low $\sig$."
+ The apparenteffect is suppressionstronger in EDEI lower og) because the difference between the equation of state at present and in the matter dominated era in this case is larger., The effect is stronger in EDE1 lower $\sig$ ) because the difference between the equation of state at present and in the matter dominated era in this case is larger.
+ Further details on how the change in equation of state affects the dark energy perturbations can be found in paper I (Alam2010).., Further details on how the change in equation of state affects the dark energy perturbations can be found in paper I \cite{ede_cmb}.
+ Paper I (Alam2010) explores effects of early dark energy on CMB power spectrum and other current observations; here we focus on forthcoming large-scale structure probes., Paper I \cite{ede_cmb} explores effects of early dark energy on CMB power spectrum and other current observations; here we focus on forthcoming large-scale structure probes.
+ We consider the influence of EDE on cluster counts and SZ power spectrum., We consider the influence of EDE on cluster counts and SZ power spectrum.
+ Redshift evolution of cluster abundance provides a valuable into global dynamics of the universe., Redshift evolution of cluster abundance provides a valuable insight into global dynamics of the universe.
+" This abundance is insightsensitive to both expansion and growth history of the universe, and as such can be used to constrain standard cosmological parameters like og or Qom (see, e.g. Haiman et al."," This abundance is sensitive to both expansion and growth history of the universe, and as such can be used to constrain standard cosmological parameters like $\sigma_8$ or $\omt$ (see, e.g. Haiman et al."
+" 2001), but also to explore possible mechanisms for observed acceleration, like modifying the law of gravity."," 2001), but also to explore possible mechanisms for observed acceleration, like modifying the law of gravity."
+" For a spatially flat cosmologies considered here (Ω,= 0), all-sky number of objects more massive than M,,,;, in a redshift bin is: where the first term is the comoving volume, and the second is the differential mass function."," For a spatially flat cosmologies considered here $\omk \equiv 0$ ), all-sky number of objects more massive than $M_{min}$ in a redshift bin is: where the first term is the comoving volume, and the second is the differential mass function."
+" From numerical simulations, we have a good quantification of the halo mass function."," From numerical simulations, we have a good quantification of the halo mass function."
+ Although its precision is still not at the level needed to make it insignificant for future cosmological parameter studies (Cunha Evrard 2009; Wu et al., Although its precision is still not at the level needed to make it insignificant for future cosmological parameter studies (Cunha Evrard 2009; Wu et al.
+" 2009), that will not be relevant here as we want to explore the relative difference between our EDE models and the fiducial ACDM cosmology."," 2009), that will not be relevant here as we want to explore the relative difference between our EDE models and the fiducial $\Lambda$ CDM cosmology."
+" The important point, on which we rely here, is that the mass function can be presented in auniversal form, where the cosmology dependence comes from the linear power spectrum, and the linear growth factor."," The important point, on which we rely here, is that the mass function can be presented in a form, where the cosmology dependence comes from the linear power spectrum, and the linear growth factor."
+" At a level (often better than that), it is indeed proven to be the case, for many different cosmological models (Jenkins et al."," At a level (often better than that), it is indeed proven to be the case, for many different cosmological models (Jenkins et al."
+ 2001; Linder Jenkins 2003; Jennings et al., 2001; Linder Jenkins 2003; Jennings et al.
+ 2010; Bhattacharya et al., 2010; Bhattacharya et al.
+ 2010) and for a wide range of, 2010) and for a wide range of
+lines of the ionized gas: Le? and. OL) for the August set of observations. and. NU] in other cases.,"lines of the ionized gas: $\beta$ and [OIII] for the August set of observations, and [NII] – in other cases."
+ The processing of spectra was done using the codes we wrote in the LDL environment (see.e.g.Zasovctal. 2008).," The processing of spectra was done using the codes we wrote in the IDL environment \cite[see, e.g.,][]{Zasov2008}."
+. The measurements of the distribution of line-of-sight velocities and stellar velocity dispersion were done by cross-correlating the spectra of galaxies with the spectra of the template star observed in the same night., The measurements of the distribution of line-of-sight velocities and stellar velocity dispersion were done by cross-correlating the spectra of galaxies with the spectra of the template star observed in the same night.
+ The measurement technique is described in our previous papers (Moiseev2001:Zasovetal. 2008).," The measurement technique is described in our previous papers \citep{Moiseev2001,Zasov2008}."
+. We observed. several template: stars belonging to the spectral tvpes HIE Gs HE WS., We observed several template stars belonging to the spectral types III G8 – III K3.
+ For the final measurements we were selecting a template. giving a maximum correlation coefficient.," For the final measurements we were selecting a template, giving a maximum correlation coefficient."
+ Line-of-sight. velocities of ionized gas were measured via the Gaussian approximations of the selected emission line profiles., Line-of-sight velocities of ionized gas were measured via the Gaussian approximations of the selected emission line profiles.
+ In this paper we only brielly discuss the results obtained. mainly confining ourselves. to. the confirmation/disproof of the presence. of a kinematically. isolated. polar component.," In this paper we only briefly discuss the results obtained, mainly confining ourselves to the confirmation/disproof of the presence of a kinematically isolated polar component."
+ La the subsequent. papers we would like to further explore the characteristics of the gas and stars in these galaxies. the composition of their stellar population. rotation curves. ete.," In the subsequent papers we would like to further explore the characteristics of the gas and stars in these galaxies, the composition of their stellar population, rotation curves, etc."
+ For each candidate. we obtained the spectrum. in. two position angles. oriented. along the major axis of the inner body and the possible polar ring.," For each candidate, we obtained the spectrum in two position angles, oriented along the major axis of the inner body and the possible polar ring."
+ SPIC-38 and SPRC-5s are the exceptions. observed only in one slit position.," SPRC-38 and SPRC-58 are the exceptions, observed only in one slit position."
+ Fig., Fig.
+ 6 demonstrates. the distribution of linc-of-sight velocities of the stars and ionized σας (from the brightest. emission lines) along each slit. position.," \ref{fig_spectr1}
+ demonstrates the distribution of line-of-sight velocities of the stars and ionized gas (from the brightest emission lines) along each slit position."
+ The kinematic components. corresponding to the central galaxy and the outer structure are also marked in the left panels on the figure.," The kinematic components, corresponding to the central galaxy and the outer structure are also marked in the left panels on the figure."
+ From now on. we will refer to them as “the central clise” and “the polar ring.," From now on, we will refer to them as “the central disc” and “the polar ring”."
+ Two kinematic components are clearly detected. here., Two kinematic components are clearly detected here.
+ The distribution of linc-of-sight. velocities corresponds to that. expected for a galaxy. in which all the ionized gas rotates in the plane. stronely inclined to the inner stellar body.," The distribution of line-of-sight velocities corresponds to that, expected for a galaxy, in which all the ionized gas rotates in the plane, strongly inclined to the inner stellar body."
+> Along the first section. with M1)=39 (the major axis of the ring) a significant gradient of line-of-sight velocities of ionized σας is observed. while the line-of-ight velocity of stars in the inner region (r«3 aresec) barely varies.," Along the first section with $PA(1)=39\degr$ (the major axis of the ring) a significant gradient of line-of-sight velocities of ionized gas is observed, while the line-of-sight velocity of stars in the inner region $(r<3$ arcsec) barely varies."
+ In the slit. position along the major axis of the central galaxy (D:1(2).= 1387) the situation is reversed here we find a zero gradient of line-of-sight velocities. of gas. ancl a substantial eradient for the stellar component. where a symmetrical rotation curve can clearly be traced.," In the slit position along the major axis of the central galaxy $PA(2)=138\degr$ ) the situation is reversed -- here we find a zero gradient of line-of-sight velocities of gas, and a substantial gradient for the stellar component, where a symmetrical rotation curve can clearly be traced."
+ Interestingly. in the slit. position with ΑΙ)=39 at reο4 aresee on both sides of the nucleus we can see a symmetrical increase in line-ol-sieht velocities of stars. in the same direction with the ionized gas in the ring.," Interestingly, in the slit position with $PA(1)=39\degr$ at $r>3-4$ arcsec on both sides of the nucleus we can see a symmetrical increase in line-of-sight velocities of stars, in the same direction with the ionized gas in the ring."
+ Apparently. what we observe here are the stars of the polar ring.," Apparently, what we observe here are the stars of the polar ring."
+ This galaxy also reveals two kinematic components. but the observed. pattern is more complex than in the previous case.," This galaxy also reveals two kinematic components, but the observed pattern is more complex than in the previous case."
+ The scatter of line-ol-sieht velocities of eas is maximal along the major axis of the ring PAC1)=35°., The scatter of line-of-sight velocities of gas is maximal along the major axis of the ring $PA(1)=35\degr$.
+ Moreover. in the NE half. the ionized gas emission extends almost twice farther than in the opposite direction away from the centre.," Moreover, in the NE half, the ionized gas emission extends almost twice farther than in the opposite direction away from the centre."
+ This slit position almost exactly coincides with the minor axis of the central disc. but line-ofsieht velocities of stars are constant only at r<4 aresec. while at large distances they almost coincide with the velocities of gas in the ring. excep [or the outermost region in the SW part of the section.," This slit position almost exactly coincides with the minor axis of the central disc, but line-of-sight velocities of stars are constant only at $r<4$ arcsec, while at large distances they almost coincide with the velocities of gas in the ring, except for the outermost region in the SW part of the section."
+ I means that just like in SPRC-10. we observe the motion of stars. belonging to the ring at a distance from the centre.," It means that just like in SPRC-10, we observe the motion of stars, belonging to the ring at a distance from the centre."
+ At the same time. line-ol-sieht velocities of ionizecl gas a r=4I aresee to the NIS from the centre almost coincide with 3...," At the same time, line-of-sight velocities of ionized gas at $r=4-12$ arcsec to the NE from the centre almost coincide with $V_{sys}$."
+ This behaviour can be explained by the fact tha a part of the ionizecl gas is located not only in the ring. bu also rotates in the central disc of the galaxy.," This behaviour can be explained by the fact that a part of the ionized gas is located not only in the ring, but also rotates in the central disc of the galaxy."
+ Collisions with gas clouds of the ring ercathy entangle the observed. picture., Collisions with gas clouds of the ring greatly entangle the observed picture.
+ Velocities of gas along the minor axis of the ring (119)=130) noticeably diller from the systemic velocity. and near the nucleus we can even observe the counter-rotation relative to the outer region.," Velocities of gas along the minor axis of the ring $PA(2)=130\degr$ ) noticeably differ from the systemic velocity, and near the nucleus we can even observe the counter-rotation relative to the outer region."
+ This most likely indicates that the polar ring is significantly. curved., This most likely indicates that the polar ring is significantly curved.
+ llere we managed to obtain only one section. — along the major axis of the ring., Here we managed to obtain only one section – along the major axis of the ring.
+ Nonetheless. the observed. picture is very similar to what we have seen in SPBRC-10 and SPBC-14.," Nonetheless, the observed picture is very similar to what we have seen in SPRC-10 and SPRC-14."
+ Namely. the line-of-sight. velocities of stars barely. vary in the centre of (r«6.S aresec). since the slit passes near the minor axis of the central disc.," Namely, the line-of-sight velocities of stars barely vary in the centre of $r<6-8$ arcsec), since the slit passes near the minor axis of the central disc."
+ At large distances. we can see the movement of stars in the ring.," At large distances, we can see the movement of stars in the ring."
+ We note here a strong asvinmetry of the rotation curves of gas in both directions [rom the centre in the SW part it continues to grow. while to the NE from the centre it achieves plateau.," We note here a strong asymmetry of the rotation curves of gas in both directions from the centre – in the SW part it continues to grow, while to the NE from the centre it achieves plateau."
+ There is only one section here as well. revealing a picture close to that. described above for SPIC-39.," There is only one section here as well, revealing a picture close to that, described above for SPRC-39."
+ In the central part. the gradient of line-of-sight velocities of stars is nearly zero.," In the central part, the gradient of line-of-sight velocities of stars is nearly zero."
+ We can see the stars. belonging to the ring at increasing distances from the nucleus.," We can see the stars, belonging to the ring at increasing distances from the nucleus."
+ Unlike in SPIC-39. the rotation curve of gas in the ring is svimetrical here on both sides from the centre.," Unlike in SPRC-39, the rotation curve of gas in the ring is symmetrical here on both sides from the centre."
+ This is the most nearby galaxy we observed., This is the most nearby galaxy we observed.
+" “Phe notes to the catalogue by Zwicky&Zwickv—(1071) indicate that this galaxy ""appears like Saturn.", The notes to the catalogue by \citet{Zwicky1971} indicate that this galaxy “appears like Saturn”.
+ Ehe SDSS colour image clearly demonstrates how the outer highly. inclined ring surrounds the central galaxy (Fig. 3))., The SDSS colour image clearly demonstrates how the outer highly inclined ring surrounds the central galaxy (Fig. \ref{fig_example}) ).
+ The results of spectral observations confirm the presence of the polar ring. which reveals both ionized gas and stars (see the slit position with ΕΤ)= 34°).," The results of spectral observations confirm the presence of the polar ring, which reveals both ionized gas and stars (see the slit position with $PA(1)=34\degr$ )."
+ The curve of Line-of-s1ght. velocities of gas along the major axis of the ring is asvnimetrical ancl considerably more extended in the NE direction from, The curve of line-of-sight velocities of gas along the major axis of the ring is asymmetrical and considerably more extended in the NE direction from
+The recent discovery of extrasolar clant planets orbiting around nearby solartype stars (παντος Butler 19985. 2000) has stimulated renewed iuterest in the theory of planet formation.,"The recent discovery of extrasolar giant planets orbiting around nearby solar–type stars (Marcy Butler 1998, 2000) has stimulated renewed interest in the theory of planet formation."
+" The objects observed so far have masses. A. that are characteristic of giant plaucts (1NE;ZUM,XE ALS). M; denoting a Jupiter mass."," The objects observed so far have masses, $M_p$, that are characteristic of giant planets $0.4 \; {\rm M}_J \; \ls \; M_p \ \ls \; 11 \; {\rm M}_J$ ), $M_J$ denoting a Jupiter mass."
+ The orbital semi-major axes are in the rauge 0101AUSa&2.5AU. aud orbital ecceutricities iu the range 0.0Z«S0.67 (NMarey. Butler 2000).," The orbital semi-major axes are in the range $0.04 \; {\rm AU} \; \ls \; a \; \ls \; 2.5 \; {\rm AU}$, and orbital eccentricities in the range $ 0.0 \; \ls \; e \; \ls \; 0.67$ (Marcy Butler 2000)."
+ Disc-protoplanet interactions have been invoked to explain the preseuce of elaut planets orbiting close to their host stars through inward orbital migration induced through disc-protoplauet tidal interaction ( ce., Disc-protoplanet interactions have been invoked to explain the presence of giant planets orbiting close to their host stars through inward orbital migration induced through disc-protoplanet tidal interaction ( eg.
+ Papaloizou. Terquem Nelson 1999. Lin ct al 2000).," Papaloizou, Terquem Nelson 1999, Lin et al 2000)."
+ Up to now. for the most part extrasolar plaucts appear to be isolated.," Up to now, for the most part extrasolar planets appear to be isolated."
+ However. a few iuultiple systems are known.," However, a few multiple systems are known."
+ The configuration of these may contain iniportaut information about their origin and possible uieration history., The configuration of these may contain important information about their origin and possible migration history.
+" Of special interest is the receutlv discovered system around στο,", Of special interest is the recently discovered system around GJ876.
+ This is found to be close to a 2:1 conunensurability., This is found to be close to a $2:1$ commensurability.
+ Such a configuration is iudecd sugeestive of orbital mueration., Such a configuration is indeed suggestive of orbital migration.
+ Comunensurable satellite systems such as the Galilean satellites ave thought to owe their origin to mueration induced by tidal interaction wit[um the central planet ( eg., Commensurable satellite systems such as the Galilean satellites are thought to owe their origin to migration induced by tidal interaction with the central planet ( eg.
+ Goldreich 1965)., Goldreich 1965).
+ Receut suaulatious of single protoplaucts im the observed inass range (lev 1999. Bryden oct al.," Recent simulations of single protoplanets in the observed mass range (Kley 1999, Bryden et al."
+ 1999. Lubow. Seibert Artvinowicz 1999) interacting with a dise with paraimecters thought to be typical of protoplauctary discs. but constrained to be in circular orbit. indicate eap formation aud upper liasμα limit consistent with the observations.," 1999, Lubow, Seibert Artymowicz 1999) interacting with a disc with parameters thought to be typical of protoplanetary discs, but constrained to be in circular orbit, indicate gap formation and upper mass limit consistent with the observations."
+ Nelson οἳ al. (, Nelson et al. (
+2000) relaxed the assumption of fixed. circular orbits. found isvaurd nueration and that the disc-protoplauct interaction leads to strong eccentricity damping.,"2000) relaxed the assumption of fixed circular orbits, found inward migration and that the disc-protoplanet interaction leads to strong eccentricity damping."
+ Due to accretion onto the central star. an inner cavity was forme: iu the dise interior to which the protoplanet orbited.," Due to accretion onto the central star, an inner cavity was formed in the disc interior to which the protoplanet orbited."
+" Simulations of two planets interacting with a disc have con, performed by EKley (2000). Bryden et al (2000). ane Masset Suclerove (2001)."," Simulations of two planets interacting with a disc have been performed by Kley (2000), Bryden et al (2000), and Masset Snellgrove (2001)."
+ So far duwarcd migration of wo planets locked iuto a 2:1 conunensurability las no con Situated., So far inward migration of two planets locked into a $2:1$ commensurability has not been simulated.
+ However. Bryden e al (2000) found a endency for eap material between the two planets in fixe circular orbits to be cleared. cuding wp interior to the iuncr λαο orbit or exterior to the outer planet orbit.," However, Bryden et al (2000) found a tendency for gap material between the two planets in fixed circular orbits to be cleared, ending up interior to the inner planet orbit or exterior to the outer planet orbit."
+ Taken together the above results sugeest a natura outcome of two protoplauets interacting with a disc is hat they orbit iuterior to an inner disc cavity while the external disc causes inward migration of the outer orit., Taken together the above results suggest a natural outcome of two protoplanets interacting with a disc is that they orbit interior to an inner disc cavity while the external disc causes inward migration of the outer orbit.
+ This catches up the inner orbit leading to the possibilty 6f resonant interaction., This catches up the inner orbit leading to the possibilty of resonant interaction.
+ ]t is the nurpose of this oper to investigate[m] such resonant interaction and whether. for reasonable," It is the purpose of this paper to investigate such resonant interaction and whether, for reasonable"
+the imüupact velocity is suticicuthly bieh.,the impact velocity is sufficiently high.
+ Collisious of muuesized ageregates at low velocities (0.2 1)) were studied in Paper HI (?).., Collisions of mm-sized aggregates at low velocities (0.2 ) were studied in Paper III \citep{WeidlingEtal:2009}.
+ All collisions at this velocity lead to bouncing and compaction of the agerceates. increasing the volume filling factor from ey=0.19 up fo o=0.36.," All collisions at this velocity lead to bouncing and compaction of the aggregates, increasing the volume filling factor from $\phi_0=0.15$ up to $\phi=0.36$."
+ Compaction at higher velocities was studied by 7.hereafterPaperIV.. who measured static and clvnamical dust-agerceate properties to be Huplemented into a smooth particle lydrodvuamics (SPID collision model (secalso7)..," Compaction at higher velocities was studied by \citet[][hereafter Paper IV]{GuettlerEtal:2009}, who measured static and dynamical dust-aggregate properties to be implemented into a smooth particle hydrodynamics (SPH) collision model \citep[seealso][]{GeretshauserEtal:2010}."
+ This model was used το reproduce laboratory experiments on the yagiuentation of ageregates. and in the present paper we will also eive deeper insight into the fragmentation of wotoplauctary dust ageregates. in particular focussing ou iuultiple collisions.," This model was used to reproduce laboratory experiments on the fragmentation of aggregates, and in the present paper we will also give deeper insight into the fragmentation of protoplanetary dust aggregates, in particular focussing on multiple collisions."
+ Fraecientation is a destructive xocess Which. at a first elauce. does not seen to aid in he formation of larger bodies.," Fragmentation is a destructive process which, at a first glance, does not seem to aid in the formation of larger bodies."
+ However. when two dust ageregatesaorof collide at a velocity above the fragmentation laeshold. which is around Lins for mun-sized dust agerceates (?7). both ageregates or at least the weaker one fragment aud the consequences of the fragmentation should be understood.," However, when two dust aggregates collide at a velocity above the fragmentation threshold, which is around 1 for mm-sized dust aggregates \citep{GuettlerEtal:2010}, both aggregates – or at least the weaker one – fragment and the consequences of the fragmentation should be understood."
+ 7? noted that fragmentation can even be helpful iu explaining the strong IR excess in the SEDs of T Tauri stars., \citet{DullemondDominik:2005} noted that fragmentation can even be helpful in explaining the strong IR excess in the SEDs of T Tauri stars.
+ A continuous cycle of growth and fragimeutation could retain sufficiently many small eraius to explain the observations while an efücieut erowth process would uot., A continuous cycle of growth and fragmentation could retain sufficiently many small grains to explain the observations while an efficient growth process would not.
+ Another positive effect of fragmentation was discovered by ? in laboratory experiments. in which they collided cursized ageregates with larecr targets.," Another positive effect of fragmentation was discovered by \citet{WurmEtal:2005b} in laboratory experiments, in which they collided cm-sized aggregates with larger targets."
+ Both cousisted of irregular SiO» dust with a erain size of 0.1 to LO ju and had an overall volue filling factor of o=0.31., Both consisted of irregular $_2$ dust with a grain size of 0.1 to 10 $\mu$ m and had an overall volume filling factor of $\phi=0.34$.
+ The velocity was in the range of 6 to 25 aand they observed that. although the projectile fracimented. a significant amount of mass. ic. up to 0.5 times the projectile mass. was trausferred to the target.," The velocity was in the range of 6 to 25 and they observed that, although the projectile fragmented, a significant amount of mass, i.e. up to 0.5 times the projectile mass, was transferred to the target."
+ ? showed in experiments that this mass transfer is also effücieut in a subsequent collision ou the same spot., \citet{TeiserWurm:2009a} showed in experiments that this mass transfer is also efficient in a subsequent collision on the same spot.
+ Multiple collisions iu a laree nuuber were then studied by ?.. who collided agerceates of approxinatelv 300 4201 in diameter with several targets at a velocity of 7.7," Multiple collisions in a large number were then studied by \citet{TeiserWurm:2009b}, who collided aggregates of approximately 300 $\mu$ m in diameter with several targets at a velocity of 7.7."
+ A stream of particles was falling for 1.5 14 and colliding with targets of different sizes on which a dusty crust developed., A stream of particles was falling for 1.5 m and colliding with targets of different sizes on which a dusty crust developed.
+ For flat targets. they found that a conical structure forms: further growth was onlv uite by the nium anele with respect to the direction of the incomine particles.," For flat targets, they found that a conical structure forms; further growth was only limited by the minimum angle with respect to the direction of the incoming particles."
+ Although the sticking process in a single collision could not be observed. it is natura to assume that the erowth process is similar to the one described by ?..," Although the sticking process in a single collision could not be observed, it is natural to assume that the growth process is similar to the one described by \citet{WurmEtal:2005b}."
+ All these experiments (222) use setups in which the projectiles were shot dowuwards. ic. iu the direction of gravity.," All these experiments \citep{WurmEtal:2005b, TeiserWurm:2009a, TeiserWurm:2009b} used setups in which the projectiles were shot downwards, i.e. in the direction of gravity."
+ Therefore. eravity might lave supported the accretion of mass to the tarect.," Therefore, gravity might have supported the accretion of mass to the target."
+ Especially for the experiments bv ?.. the influence of evavity leads to a re-accretion of rebound particles after jo Mupact. which increases the accretion efficiency.," Especially for the experiments by \citet{TeiserWurm:2009b}, the influence of gravity leads to a re-accretion of rebound particles after the impact, which increases the accretion efficiency."
+ The uthors compare this with a re-accretion induced by ιο drag of the surrounding gas in the protoplanctary disk., The authors compare this with a re-accretion induced by the drag of the surrounding gas in the protoplanetary disk.
+ The simall ejecta couple to the eas aud are carried ck to the target body., The small ejecta couple to the gas and are carried back to the target body.
+" This effect was first observed in experiments bv 77./— and afterwards discussed by. οον, and 7.."," This effect was first observed in experiments by \citet{WurmEtal:2001a, WurmEtal:2001b} and afterwards discussed by \citet{SekiyaTakeda:2003, SekiyaTakeda:2005}, and \citet{WurmEtal:2004}."
+ The re-acerction requires a eas flow towards ιο farget surface., The re-accretion requires a gas flow towards the target surface.
+ Πο]Ποσο re-accretion ds in xinciple oulv valid as long as the mean free path of the eas molecules is lavecr than the size of the tarect body., High-efficiency re-accretion is in principle only valid as long as the mean free path of the gas molecules is larger than the size of the target body.
+ As the experiments showed re-accreation also for higher pressures (or larger targets. respectively). it has been coutroversially discussed ii how far this XUOCcess IS relevant for the growth of kim-sized Even without easons drag.We the lins trausfer du a fraciuenting collision can du principle lead to the formation of larger bodies but the question roniadnus whether the necessary comitions for such an agerceate erowtli a sufücient size ratio between xojectile and target agerceateage (1.6. the projectile needs to be sjeuificautlv smaler than tje farget) and a sufficieutlv large collision velocity (1.6. he velocity has to be aOVE the fragmentation threshokL velocity) can be imet iu a protoplanetary disk.," As the experiments showed re-accreation also for higher pressures (or larger targets, respectively), it has been controversially discussed in how far this process is relevant for the growth of km-sized Even without gas drag, the mass transfer in a fragmenting collision can in principle lead to the formation of larger bodies but the question remains whether the necessary conditions for such an aggregate growth – a sufficient size ratio between projectile and target aggregate (i.e. the projectile needs to be significantly smaller than the target) and a sufficiently large collision velocity (i.e. the velocity has to be above the fragmentation threshold velocity) – can be met in a protoplanetary disk."
+ Mloreover. the efficicney of the nüass-transfer process from xojectile to target. nec:s to be high enough to male this erowth mode duportant for the protoplanetary dust evolition.," Moreover, the efficiency of the mass-transfer process from projectile to target needs to be high enough to make this growth mode important for the protoplanetary dust evolution."
+ Iu a novel approach to unity laloratory aud theoreical work. ? reviewer 19 laboratory experiments on cust-aeerceateOO collisions aud compiled them to the first complete dust-ageregate collisiou model that makes a prediction| for the οιtcome of anv arbitrary dust agerceate collision. iuchuling sticking. bouncing. and fraguentation.," In a novel approach to unify laboratory and theoretical work, \citet{GuettlerEtal:2010} reviewed 19 laboratory experiments on dust-aggregate collisions and compiled them to the first complete dust-aggregate collision model that makes a prediction for the outcome of any arbitrary dust aggregate collision, including sticking, bouncing, and fragmentation."
+ Although he pareuueter space is huge aud so far ouly sparsely COVCLed by laboratory experiments. this work is an imporaut step to include all the widespread experinieutal evidences.," Although the parameter space is huge and so far only sparsely covered by laboratory experiments, this work is an important step to include all the widespread experimental evidences."
+ To flud out. which of the many plhivsical pre)CONSÓNR identified bv ?/— really occur under realistic comitious 1La protoplanctary disk. 7o juuplemeuted tus collision mxlel iuto a Monte-C'arlo erowth model.," To find out, which of the many physical processes identified by \citet{GuettlerEtal:2010} really occur under realistic conditions in a protoplanetary disk, \citet{ZsomEtal:2010a} implemented this collision model into a Monte-Carlo growth model."
+ TheirThe ound tat protoplanetary dust ageroegatesOO GrOW ο a certali size (up to ceutiueYS. depending on the detailed conlitions aud the nebula nodel). at which »»ut their furher erowth is inhibited due to bouncing.," TheirThe found that protoplanetary dust aggregates grow to a certain size (up to centimeters, depending on the detailed conditions and the nebula model), at which point their further growth is inhibited due to bouncing."
+ The siaulatkuns bv? were zoro-dimensional (0D) iu space. Vail for one location in he protoplauctary disk. alc neelecting global particle notion.," The simulations by \citet{ZsomEtal:2010a} were zero-dimensional (0D) in space, valid for one location in the protoplanetary disk, and neglecting global particle motion."
+ The picture changed when 7? included sedimentation. lus performed a simulation iu a verical columm of the disk (1D).," The picture changed when \citet{ZsomEtal:2010b} included sedimentation, thus performed a simulation in a vertical column of the disk (1D)."
+ Due to turbulent mixing iu he vertical direction. the size distribution became wider. which also resulted in larger relaive velocities.," Due to turbulent mixing in the vertical direction, the size distribution became wider, which also resulted in larger relative velocities."
+" A larecr variety iu collision velocities aud size ratios also leads to a arecr variety of collisional outcoues,", A larger variety in collision velocities and size ratios also leads to a larger variety of collisional outcomes.
+ In the simulations o7. also the raenaientation wihi inass transfer (81inhenotationof ?).. as descerbec by ?7.. cune into play aud cau heuce re regarde doasa votential path to larecr xotoplanetarv )odies.," In the simulations by \citet{ZsomEtal:2010b}, also the fragmentation with mass transfer \citep[S4 in the notation of][]{GuettlerEtal:2010}, as described by \citet{WurmEtal:2005b}, came into play and can hence be regarded as a potential path to larger protoplanetary bodies."
+ Iun this paper. we study the growth of a dust-agelomierateao target due to the racsniecntation of many unu-xózed porous projectiles at velocities between 1.5 aud 6.0.," In this paper, we study the growth of a dust-agglomerate target due to the fragmentation of many mm-sized porous projectiles at velocities between 1.5 and 6.0."
+ This velocity rauge is chosen to complement xevious experiments of ? aud it is moreover in the range xecdietec by the simulations of ?.., This velocity range is chosen to complement previous experiments of \citet{WurmEtal:2005b} and it is moreover in the range predicted by the simulations of \citet{ZsomEtal:2010b}.
+ These high-velocity collisions lead to fragmentation with mass transfer aud are slightIv above the fragmentation threshold. velocity., These high-velocity collisions lead to fragmentation with mass transfer and are slightly above the fragmentation threshold velocity.
+ We stuv nultiple collisions because the accretion efficicucy in a single collision is small aud so. to get a sienificant amount of erowtl. a large umber of collisions is required.," We study multiple collisions because the accretion efficiency in a single collision is small and so, to get a significant amount of growth, a large number of collisions is required."
+ A self-consistent body. forms after the first few impacts have coated the target with a laver of dust., A self-consistent body forms after the first few impacts have coated the target with a layer of dust.
+ Iu Sect.??.. we will present the experiucutal setup.," In \ref{section_Setup}, , we will present the experimental setup,"
+observations does not require a non-zero 7j.,observations does not require a non-zero $\eta_r$.
+ The surface flux transport model including the effect of a evele-dependent variation of the Glt angles of sunspot groups reproduces (he major features of the observationally inferred open flix aud (he timing of the polar field reversals from 1913 to 1986 (the period for which we have the tilt angle data)., The surface flux transport model including the effect of a cycle-dependent variation of the tilt angles of sunspot groups reproduces the major features of the observationally inferred open flux and the timing of the polar field reversals from 1913 to 1986 (the period for which we have the tilt angle data).
+ The reversal of polar fields alter strong cveles can be explained bv the observed anti-correlation between the active region tilt angle and evcle aunplitude. so that no additional decay by racial diffusion was required to achieve this result.," The reversal of polar fields after strong cycles can be explained by the observed anti-correlation between the active region tilt angle and cycle amplitude, so that no additional decay by radial diffusion was required to achieve this result."
+ When the observed (ilt angle are used. the polar fiekl maxima [rom the model are correlated. with the strength of the following cvcle.," When the observed tilt angle are used, the polar field maxima from the model are correlated with the strength of the following cycle."
+ This correlation is likely to be present independent of (he parameter choices. provided the model reproduces the minima of the inlerred open flux.," This correlation is likely to be present independent of the parameter choices, provided the model reproduces the minima of the inferred open flux."
+ The correlation suggests that the polar fields ave an important ingredient of the solar dynamo process. which is consistent with Dabcock-Leighton-tvpe models.," The correlation suggests that the polar fields are an important ingredient of the solar dynamo process, which is consistent with Babcock-Leighton-type models."
+ The evele-lo-cvele variation of Jov's law might play a role in the nonlinear modulation of the solar dvnamo., The cycle-to-cycle variation of Joy's law might play a role in the nonlinear modulation of the solar dynamo.
+Giant micropulses from the Vela pulsar at 2.3 GIIz also have apparently power-law statistics al laree Fo that obey the form (6)). this time with »22.85d0.3 2002).. where the error is again estimated bv eve.,"Giant micropulses from the Vela pulsar at 2.3 GHz also have apparently power-law statistics at large $F$ that obey the form \ref{cdf_vela}) ), this time with $\beta \approx 2.85 \pm 0.3$ \citep{krameretal2002}, where the error is again estimated by eye."
+ Converting into the form (2)). as lor D1T06-44. leads to a=6.7x0.6.," Converting into the form \ref{power-law}) ), as for B1706-44, leads to $\alpha = 6.7 \pm 0.6$."
+ Giant micropulses from D0950--08 at 0.43 611 have apparently power-law statistics al large F that obev the form (2)) with index a=6.2£0.5 (Cairnsοἱal.2003€)., Giant micropulses from B0950+08 at 0.43 GHz have apparently power-law statistics at large $F$ that obey the form \ref{power-law}) ) with index $\alpha = 6.2 \pm 0.5$ \citep{cetal2002c}.
+.. These giant micropulses are observed over a relatively large range of pulsar phase compared wilh those for Vela and D1706-44., These giant micropulses are observed over a relatively large range of pulsar phase compared with those for Vela and B1706-44.
+ Several other classes of emissions were identified with approximately power-law statistics., Several other classes of emissions were identified with approximately power-law statistics.
+ However. their indices were substantially less. with values in the range 1-4.," However, their indices were substantially less, with values in the range $1 - 4$."
+ Summarizing. the field distributions observed near 2 GIIz and 0.4 GI for giant micropulses from three pulsars (Vela. D1T06-44. and D09050--03) can be interpreted in terms of with high indices zz6.540.7.," Summarizing, the field distributions observed near 2 GHz and $0.4$ GHz for giant micropulses from three pulsars (Vela, B1706-44, and B0950+08) can be interpreted in terms of power-laws with high indices $\approx 6.5 \pm 0.7$."
+ These limited data provide no evidence for à. varving with observing Irecuency., These limited data provide no evidence for $\alpha$ varying with observing frequency.
+ Giant pulses [rom the Crab pulsar also have apparently power-law statistics at large F:, Giant pulses from the Crab pulsar also have apparently power-law statistics at large $F$ :
+ ofa few variability in the accretion rate are immecdiatelv obvious., of a few variability in the accretion rate are immediately obvious.
+ This variability can be tracked. down to the arrival of individual cool gas blobs or filaments in 16 innermost region., This variability can be tracked down to the arrival of individual cool gas blobs or filaments in the innermost region.
+ Since the density of the hot eas in the vicinity of iis never zero. the accretion rate never decreases to Zero.," Since the density of the hot gas in the vicinity of is never zero, the accretion rate never decreases to zero."
+ It must be stressed that the realVpy further depends on intricate physical details of the inner accretion low tha we cannot resolve here., It must be stressed that the real$\dot{M}_{\rm BH}$ further depends on intricate physical details of the inner accretion flow that we cannot resolve here.
+ Some of the gas entering A—Fh ds unbound and some may become unbound later on as resul of viscous heating in the Dow (?).., Some of the gas entering $R \leq R_{\rm in}$ is unbound and some may become unbound later on as result of viscous heating in the flow \citep{Blandford99}.
+ Pherefore. the acerction rate measured in our simulations is best understood to vielc the outer boundary conditions for the inner accretion Llow. and as such should be a more physically complete estimate of that than the commonly used Bondi accretion rate base on observations of hot. X-ray emitting gas only.," Therefore, the accretion rate measured in our simulations is best understood to yield the outer boundary conditions for the inner accretion flow, and as such should be a more physically complete estimate of that than the commonly used Bondi accretion rate based on observations of hot X-ray emitting gas only."
+ eaach time-step., ach time-step.
+" The cold ""dise found in our simulations is constantly being created but it is also likely to be destroved from. time to time.", The cold `disc' found in our simulations is constantly being created but it is also likely to be destroyed from time to time.
+" Lis mass can be estimated as Ma=AL""EE £a. where εν=(/107. vears and AL is the mass outflow rate of cool stellar winds."," Its mass can be estimated as $M_{\rm disc} = \dot{M}_{\rm cw} t = 10 \msun
+\;(\dot{M}_{\rm cw}/10^{-3} \msun \hbox{year}^{-1})\; t_4$ , where $t_4
+= t/10^4$ years and $\dot{M}_{\rm cw}$ is the mass outflow rate of cool stellar winds."
+ A supernova occurring within the inner 0.5 pe of the Galaxy would easily destroy such a disc., A supernova occurring within the inner $0.5$ pc of the Galaxy would easily destroy such a disc.
+ Phe number of voung early type stars in the ccluster is likely to be in the hundreds. and some have already. reached. the WR. stage (e.g.22)...," The number of young early type stars in the cluster is likely to be in the hundreds, and some have already reached the WR stage \citep[e.g.][]{Paumard01,Genzel03a}."
+ Assuming that the most recent star formation event occurred a few million vears ago. one would then estimate the SN rate in ss inner star cluster to be at least ~107 per vear.," Assuming that the most recent star formation event occurred a few million years ago, one would then estimate the SN rate in s inner star cluster to be at least $\sim
+10^{-4}$ per year."
+" For comparison. the viscous time at a radius of 1"" is LO""10 vears at best (c.g.fig.2in7).."," For comparison, the viscous time at a radius of $1''$ is $10^6-10^7$ years at best \citep[e.g. fig. 2 in][]{NC05}."
+ Furthermore. there is much more cold material on scales of several arcseconds ane bevond that may be plunging on to (7)..," Furthermore, there is much more cold material on scales of several arcseconds and beyond that may be plunging on to \citep{Paumard04}."
+ Arrival of this mass in the inner aresecond coul also destroy the disc., Arrival of this mass in the inner arcsecond could also destroy the disc.
+ Vherefore it is likely that the cool dise of ~ arcsecon scales does not smoothly extend. inside the sub-aresecone region ofA., Therefore it is likely that the cool disc of $\sim$ arcsecond scales does not smoothly extend inside the sub-arcsecond region of.
+".. However. occasionally we observe. coo eas blobs to directly. fall into the capture ""sphere! in the simulation."," However, occasionally we observe cool gas blobs to directly fall into the capture `sphere' in the simulation."
+ Lt is also conceivable that events started by a supernova shell passage. or by an infall of additional coo material. could also leave a remnant in the form of a cole disc.," It is also conceivable that events started by a supernova shell passage, or by an infall of additional cool material, could also leave a remnant in the form of a cold disc."
+ Phermal conduction between the two phases could Leac o evaporation of a cold accretion disc via the 2) mechanism. which is unfortunately model dependent due to our. poor knowledge of the magnetic field geometry and strength.," Thermal conduction between the two phases could lead to evaporation of a cold accretion disc via the \cite{Meyer94}
+ mechanism, which is unfortunately model dependent due to our poor knowledge of the magnetic field geometry and strength."
+ 7) suggested. mainlv based. on the presence of X-ray lues inA®*.. that there is a cool disc at ~0.01 VI” scales and beyond.," \cite{NS03} suggested, mainly based on the presence of X-ray flares in, that there is a cool disc at $\sim 0.01$ $0.1''$ scales and beyond."
+ The required mass of the disc was estimated to be smaller than a Solar mass., The required mass of the disc was estimated to be smaller than a Solar mass.
+ However. observations of NUR Hares (2). put the Lares at no more han a few milli-arescconcds away from the radio position of which is somewhat problematic for this moclel.," However, observations of NIR flares \citep{Genzel03b}
+ put the flares at no more than a few milli-arcseconds away from the radio position of, which is somewhat problematic for this model."
+ Also. the NIR: flare spectra appear to strongly favour a svnchrotron (ie. SMBDIL jet) origin.," Also, the NIR flare spectra appear to strongly favour a synchrotron (i.e. SMBH jet) origin."
+" Further. no eclipses or star ""hrightenings. expected when bright stars approach the disc (2). have been observed so far."," Further, no eclipses or star `brightenings' expected when bright stars approach the disc \citep{Cuadra04} have been observed so far."
+ Summarising. there is currently no observational motivation to favour the existence of such a dise inA... although it is also cillicult to rule out its presence (unlessthediscextendstoSMDLLhori-zon: ?)..," Summarising, there is currently no observational motivation to favour the existence of such a disc in, although it is also difficult to rule out its presence \citep[unless the disc extends to the SMBH horizon;][]{Falcke97}."
+ Concluding. we suggest that there cannot be a smooth transition of the larger scale cool disc into the inner sub-arcsecond regions of A.," Concluding, we suggest that there cannot be a smooth transition of the larger scale cool disc into the inner sub-arcsecond regions of ."
+ Nevertheless. the issue of the current existence of a cool disc or its periodic appearance and disappearance is an open subject for future work.," Nevertheless, the issue of the current existence of a cool disc or its periodic appearance and disappearance is an open subject for future work."
+been cjecked. interactively.,been checked interactively.
+ We used a probability threshok of 99.91% (about 3.5 7)., We used a probability threshold of 99.94 (about 3.5 $\sigma$ ).
+ Soirees detected in reeicàs of radius L. 6 or δ arcinin aronud targets (depending on the target brightuess) have been excluded. from the sauple. which iucludes a total oLL? sources down Τεpa 5-10 keV flux of ~5«101H.," Sources detected in regions of radius 4, 6 or 8 arcmin around targets (depending on the target brightness) have been excluded from the sample, which includes a total of 147 sources down to a 5-10 keV flux of $\sim5\times10^{-14}$."
+ Coun rates were converted to fluxes using a fixed conversion actor equal to S10! (D-]5-10 keV flux) per one “3 MECS count” (1.5-10 keV)., Count rates were converted to fluxes using a fixed conversion factor equal to $\times10^{-11}$ (5-10 keV flux) per one “3 MECS count” (4.5-10 keV).
+" Tus factor is appropriat for a power law spectrum with ag=0.6. but due to he narrow bad it is no strongly scusitive to the spectral shape: for arg0, and 0.5 it is S ancl 74 «101 respectivev."," This factor is appropriate for a power law spectrum with $\alpha_E=0.6$, but due to the narrow band it is not strongly sensitive to the spectral shape: for $\alpha_E=0.4$ and 0.8 it is 8.1 and 7.6 $\times10^{-11}$ respectively."
+" The skvcovcrage varies+ from. ~E1 dee?J : 5«101l to 6.G. 50 are Lal deg? at fuxes of 1000,3« BaAC >LO respectively."," The skycoverage varies from $\sim1$ $^2$ at $5\times10^{-14}$ to 6.6, 50 and 84 $^2$ at fluxes of $10^{-13}, 3\times10^{-13}$ and $>10^{-12}$ respectively."
+ After correcting for this skyvcoverage we fiud 16.94: ?OCG.L} sources > at Fsπρωι{δν10li (Fiore ct a., After correcting for this skycoverage we find $\pm$ 3.0(6.4) sources $^{-2}$ at $F_{5-10 keV}=4.8\times10^{-14}$ (Fiore et al.
+ 2000a. Coma strict al.," 2000a, Comastri et al."
+ 2000)., 2000).
+ First quotec errors are the las atistical confidence nuανα]., First quoted errors are the $1\sigma$ statistical confidence interval.
+" Exrors in brackets include svsteniaic uucertaimties. due to the ack of kjowledee of tie real spectrum of the faint sources,"," Errors in brackets include systematic uncertainties, due to the lack of knowledge of the real spectrum of the faint sources."
+ This k2N-loes COLTCsponds to a resolved fr:iction of the 5-10 NRB equal to 20-30.. de]5idine on the NRB nornalization (C1nastri this mectinge. Vecchi et al.," This logN-logS corresponds to a resolved fraction of the 5-10 XRB equal to 20-30, depending on the XRB normalization (Comastri this meeting, Vecchi et al."
+ 1999)., 1999).
+ The observed mumrer counts are consisteit. within the errors. with the extrapolations youn the DepSAN (€dom et al 2100} and ASC'A (Cagnoni et al.," The observed number counts are consistent, within the errors, with the extrapolations from the BeppoSAX (Giommi et al 2000) and ASCA (Cagnoni et al."
+ 1998. Ueda et al.," 1998, Ueda et al."
+ 1999. Dellv Coca et al.," 1999, Della Ceca et al."
+ 1999) 2-1t) keV. umber counts. assuni1 alb average powcr law specruni wit pop0.6.," 1999) 2-10 keV number counts, assuming an average power law spectrum with $\alpha_E\sim0.6$."
+ To study the spectral variety of the TELLAS sources we lave calculated for cach source tie softness ratio (S-ID)/(S|IT). S=1.5-15 keV. T=1.5-10 keV. Fieire 1 plots (S-ID/(S|II) as a function of source total 5-10 keV fiux.," To study the spectral variety of the HELLAS sources we have calculated for each source the softness ratio (S-H)/(S+H), S=1.3-4.5 keV, H=4.5-10 keV. Figure 1 plots (S-H)/(S+H) as a function of source total 5-10 keV flux."
+ Sources ταider or close to the 1CLIluni strougback supporting the MECS window have been excluded. from this anavais. because their «ibserved softuess iuav be systematically QWOL lia1 real.," Sources under or close to the berillium strongback supporting the MECS window have been excluded from this analysis, because their observed softness may be systematically lower than real."
+ The uuuber of the remaiming sources in the sample is 126., The number of the remaining sources in the sample is 126.
+ Many «ft the HELLAS ποιασος rave a low (S-ID/(S|ID. indicating a hard spectrum.," Many of the HELLAS sources have a low (S-H)/(S+H), indicating a hard spectrum."
+ Assumiueg eE—0.6. we fud hat 3605)5 of he 12G sources have 1.3-1.5j keV count rates Owey hai that expectoc at confidencee low1 29554(299.7," Assuming $\alpha_E=0.6$, we find that 36(5) of the 126 sources have 1.3-4.5 keV count rates lower than that expected at confidence level $\gs 95\%(\gs99.7\%)$."
+ Large absorbing coluuus desities are |ikely responsible or the Did spectrum of these sources., Large absorbing columns densities are likely responsible for the hard spectrum of these sources.
+ A defici of very hard sources at fluxes 12N1013 Is also evident m flewe 1., A deficit of very hard sources at fluxes $\ls1-2\times10^{-13}$ is also evident in figure 1.
+ This cal be cue to both an astrophysical aud a teclinical reason., This can be due to both an astrophysical and a technical reason.
+ The first is a rcdshitt effect: the oserved softuess ratio of sources with simular mmtriusie absorbiu e“ol deusitv increases with he redshift. as the observed. cut-off energy LOVES toward lower energies.," The first is a redshift effect: the observed softness ratio of sources with similar intrinsic absorbing column density increases with the redshift, as the observed cut-off energy moves toward lower energies."
+" The secud is the MECS reduced sensitivity to haze ΒΟΛΟΥ""S due to tlje rapid increase of the vignetting of the telescopes with the energy and with the «taxis angle."," The second is the MECS reduced sensitivity to hard sources, due to the rapid increase of the vignetting of the telescopes with the energy and with the off-axis angle."
+ To qiautify the latter effect we have computed loci of eqal keV couif rate for a eiven flux. shown bv dashed lines in figure 1.," To quantify the latter effect we have computed loci of equal 4.5-10 keV count rate for a given flux, shown by dashed lines in figure 1."
+ The strong curvature o ‘these lines oward low values of (S-IT)/(5|IT) indicate that a large part of the defici of faint hard sources is probably due to this effect., The strong curvature of these lines toward low values of (S-H)/(S+H) indicate that a large part of the deficit of faint hard sources is probably due to this effect.
+ The curves are bet, The curves are bent
+"photons, are always significant in long II) GRBs (Norris et al.","photons, are always significant in long (type II) GRBs (Norris et al."
+ 2000; Gehrels et al., 2000; Gehrels et al.
+ 2006; (typeLiang et al., 2006; Liang et al.
+ 2006; Zhang et al., 2006; Zhang et al.
+" 2009), but are typically negligible for short GRBs Bonnell 2006; Zhang et al."," 2009), but are typically negligible for short (type I) GRBs (Norris Bonnell 2006; Zhang et al."
+ (typ, 2009).
+"eIn I)order to get (Norrishigh signal-to-noise ratio, we only 2009).select the brightest part (Episode I and II, as shown in Fig."," In order to get high signal-to-noise ratio, we only select the brightest part (Episode I and II, as shown in Fig."
+ 3 and listed in Table 2) of each sub-burst to calculate lags., \ref{fig:multilc} and listed in Table 2) of each sub-burst to calculate lags.
+" For the first sub-burst, we extracted 32ms-binned light curves in the following four BAT energy bands: 15—25 keV, 25—50 keV, 50—100 keV and 100—150 keV and 64ms-binned light curvesin the following threeWIND bands: 25—95 keV, 95—380 keV and 380—1435 keV. Then, using the CCF (cross-correlation function; Norris et al."," For the first sub-burst, we extracted 32ms-binned light curves in the following four BAT energy bands: $15-25$ keV, $25-50$ keV, $50-100$ keV and $100-150$ keV and 64ms-binned light curvesin the following three bands: $25-95$ keV, $95-380$ keV and $380-1435$ keV. Then, using the CCF (cross-correlation function; Norris et al."
+" 2000, Ukwatta et al."," 2000, Ukwatta et al."
+" 2010) method, we calculate the lags between any two light curves in the neighboring and next-to-neighbor energy bands within each instrument in Episode I. The uncertainty of lags are estimated by Monte Carlo simulation (see e.g., Peterson et al."," 2010) method, we calculate the lags between any two light curves in the neighboring and next-to-neighbor energy bands within each instrument in Episode I. The uncertainty of lags are estimated by Monte Carlo simulation (see e.g., Peterson et al."
+" 1998, Ukwatta et al."," 1998, Ukwatta et al."
+ 2010) and are illustrated in Fig., 2010) and are illustrated in Fig.
+ 10 and 11.., \ref{fig:ccf} and \ref{fig:ccfsim}.
+" For the second sub-burst, we extracted 32ms-binned light curves in the same four BAT energy bands as mentioned above and three XRT energy bands: 0.3—1 keV, 1—4 keV and 4—10 keV. Then using the same method we calculated the spectral lags between these energy bands."," For the second sub-burst, we extracted 32ms-binned light curves in the same four BAT energy bands as mentioned above and three XRT energy bands: $0.3-1$ keV, $1-4$ keV and $4-10$ keV. Then using the same method we calculated the spectral lags between these energy bands."
+ Our results are shown in Table 2., Our results are shown in Table 2.
+ Some lags are not well-constrained possibly due to low signal- levels and the combination of multiple pulses., Some lags are not well-constrained possibly due to low signal-to-noise levels and the combination of multiple pulses.
+ Yet the typical values of 201+52 ms between 25—50 keV and 50—100 keV for the first sub-burst is similar with other long GRBs et al., Yet the typical values of $201\pm52$ ms between $25-50$ keV and $50-100$ keV for the first sub-burst is similar with other long (type II) GRBs (Zhang et al.
+ 2009)., 2009).
+" In Fig 12,, we plot the (typeluminosity-lagII) (Zhangdiagram by assuming the double burst is at redshift 0.1—7.0."," In Fig \ref{fig:lplag}, we plot the luminosity-lag diagram by assuming the double burst is at redshift $0.1-7.0$."
+" We found that at the average redshift (z~ 2) of Swift GRBs, GRB 110709B falls into the *main track"" of typical long/Type II GRBs in the luminosity-lag diagram."," We found that at the average redshift $z\sim 2$ ) of GRBs, GRB 110709B falls into the “main track"" of typical long/Type II GRBs in the luminosity-lag diagram."
+differences in the stellar populations coutributing to the flux iu a ]oxel,differences in the stellar populations contributing to the flux in a pixel.
+ After applying the pixel-based. two-dimensional extinction correction described in Paper 1. we defiued six. different pixel groups iu this diagnostic diagram. raneine frou pixels for which both acad ccolors iudicate that their fluxes are domiuated by very voune. lassive stellar populations. to pixels that appear to saluple οuly light from evolved stellar populations.," After applying the pixel-based, two-dimensional extinction correction described in Paper I, we defined six different pixel groups in this diagnostic diagram, ranging from pixels for which both and colors indicate that their fluxes are dominated by very young, massive stellar populations, to pixels that appear to sample only light from evolved stellar populations."
+ We demonstrated that it was not possible to meaninefully define such pixel eroups extinction correction., We demonstrated that it was not possible to meaningfully define such pixel groups extinction correction.
+ We then showx«l that pixels that belong to a given pixel eroup form well-defined. contiguous regions in a pixel coordinate lap. revealing svsteniatic spatial variations in the dominant stellar populations that would not be readily discοποίο without a two-dimensional correction or extinction.," We then showed that pixels that belong to a given pixel group form well-defined, contiguous regions in a pixel coordinate map, revealing systematic spatial variations in the dominant stellar populations that would not be readily discernible without a two-dimensional correction for extinction."
+ We were able to report the presence of a xeviouslv unrecognized stellar bar. for example.," We were able to report the presence of a previously unrecognized stellar bar, for example."
+ Our pixel-based two-dimensional inethod to correct for extinction. based on only aand V-baud mages (adding C improves fidelity). has he advantage that it is applicable to anv galaxy. that is sienificautly resolved at rest-frame ~3.6 aan that it is not restricted to a few incliviclual sielitlines within a galaxw. nor to the werv nearest ealaxies.," Our pixel-based two-dimensional method to correct for extinction, based on only and $V$ -band images (adding $U$ improves fidelity), has the advantage that it is applicable to any galaxy that is significantly resolved at rest-frame $\sim$ and that it is not restricted to a few individual sightlines within a galaxy, nor to the very nearest galaxies."
+ Athough secuunely simple aud crude. we lave demonstrated that application of this method allowed us to uncover relatively detailed spatial information on the nature of sellar populations aud ou the star formation history within 00959. despite sienificaut spatial variations 1- he attenuation by intervening dust.," Although seemingly simple and crude, we have demonstrated that application of this method allowed us to uncover relatively detailed spatial information on the nature of stellar populations and on the star formation history within 0959, despite significant spatial variations in the attenuation by intervening dust."
+" While the spatial resolution in the UM study was limited by the NUV PSF offuture, (EWIIND to linear scales of ~250 ppc. in the near conibinatiou of //NTEC3TL and will provide much higher resolution. = (111) at rest-frame UV through mid-IR wavelengths."," While the spatial resolution in the present study was limited by the NUV PSF of 3 (FWHM) to linear scales of $\sim$ pc, in the near future, combination of /WFC3 and will provide much higher resolution $\simeq$ 11) at rest-frame UV through mid-IR wavelengths."
+ These instruments will thus allow us to study the distribution of dust oxtinetion and the underlving stellar populations at high fidelity and at much finer spatial resolution within nearby ealaxies. as well as at similar scales of a few πάτος pe within more distant galaxies.," These instruments will thus allow us to study the distribution of dust extinction and the underlying stellar populations at high fidelity and at much finer spatial resolution within nearby galaxies, as well as at similar scales of a few hundred pc within more distant galaxies."
+ Since galaxies at + 2 11 are iuportant building blocks of the galaxies that we sce today at 7200. studvine these hieher redshitt objects with proper two-dimensional extinction estimates will be of sienificaut iniportaunce to reveal the nature of galaxy assembly aud evolution.," Since galaxies at $z$ $\gtrsim$ 1 are important building blocks of the galaxies that we see today at $z$ $\simeq$ 0, studying these higher redshift objects with proper two-dimensional extinction estimates will be of significant importance to reveal the nature of galaxy assembly and evolution."
+optical imaging requires the removal of continuum emission. which can be done through the subtraction of appropriately-sealed 2-band images.,"optical imaging requires the removal of continuum emission, which can be done through the subtraction of appropriately-scaled $R$ -band images."
+ “Phere are. uncertainties in this process which may lead the observer to question whether there is He emission. or whether the apparent. emission is à result of poor continuum subtraction. particularly from poorly subtracted stars or other image artifacts.," There are uncertainties in this process which may lead the observer to question whether there is $\alpha$ emission, or whether the apparent emission is a result of poor continuum subtraction, particularly from poorly subtracted stars or other image artifacts."
+ Most of these can be rejected. from. visual inspection of the continuunrsubtractec lla image (an example of a stellar residual is shown towards the lower-right hand corner of the example image presented in Fig. l1.," Most of these can be rejected from visual inspection of the continuum-subtracted $\alpha$ image (an example of a stellar residual is shown towards the lower-right hand corner of the example image presented in Fig. \ref{imagefig},"
+ with characteristic positive ancl negative zones). ancl through their low equivalent widths.," with characteristic positive and negative zones), and through their low equivalent widths."
+" However. at faint magnitudes and for stars with very extreme colours this process is less clear-cut,"," However, at faint magnitudes and for stars with very extreme colours this process is less clear-cut."
+ There is also the possibility ofother. shorter wavelength emission. lines from intermediate-redshift ealaxies. lor example the A5007 line. being redshifted into the passhand of the Lla filter.," There is also the possibility of other, shorter wavelength emission lines from intermediate-redshift galaxies, for example the $\lambda$ line, being redshifted into the passband of the $\alpha$ filter."
+A Adclitionally. the narrow-band Glter widths. typically 1%. allow Lla emission to be detected in à range which is much wider than the velocities of the satellites with reference to their hosts.," Additionally, the narrow-band filter widths, typically $\%$, allow $\alpha$ emission to be detected in a range which is much wider than the velocities of the satellites with reference to their hosts."
+ This means that even if the line detected is truly Ho. any apparent satellites could casily be separated by ~1000  from the central galaxy. and thus not be truly associated companions.," This means that even if the line detected is truly $\alpha$, any apparent satellites could easily be separated by $\sim1000$ $^{-1}$ from the central galaxy, and thus not be truly associated companions."
+ The aims of this paper are to analyse spectroscopically the potential satellites detected. through Lla imaging. to quantify the above uncertainties and ascertain the reliability ofthe Lla imaging method for detecting true companions of field spiral galaxies.," The aims of this paper are to analyse spectroscopically the potential satellites detected through $\alpha$ imaging, to quantify the above uncertainties and ascertain the reliability of the $\alpha$ imaging method for detecting true companions of field spiral galaxies."
+ From an original imaging cata set of 262 field. spiral ealaxies. 73 candidate satellites associated with 56 host ealaxies were discovered of which Ll hosts and 12 candidate satellites are to be. spectroscopically tested and: results presented in the present paper.," From an original imaging data set of 262 field spiral galaxies, 73 candidate satellites associated with 56 host galaxies were discovered of which 11 hosts and 12 candidate satellites are to be spectroscopically tested and results presented in the present paper."
+ The complete set of imaging data are to be analysed. to. determine. the integrated properties and spatial distribution of SE. the SE rates and timescales. ancl their {1ρα and. Lea luminosities.," The complete set of imaging data are to be analysed to determine the integrated properties and spatial distribution of SF, the SF rates and timescales, and their $R$ -band and $\alpha$ luminosities."
+ These properties of the central galaxies. and analvis of the numbers and. properties of their satellites. will be studied in a series of separate papers.," These properties of the central galaxies, and analyis of the numbers and properties of their satellites, will be studied in a series of separate papers."
+ The original imaging data set of nearby spiral galaxies was observed with the Jacobus Ixapteyn Telescope CJINTT) or the Wa Galaxy Survey. Πας (Jamesetal.2004).," The original imaging data set of nearby spiral galaxies was observed with the Jacobus Kapteyn Telescope (JKT) for the $\alpha$ Galaxy Survey, $\alpha$ GS \citep{jam04}."
+. Additional galaxies were observed. with the Isaac Newton Telescope (INTE) for à study of supernova-hosting galaxies Anderson&James2008)., Additional galaxies were observed with the Isaac Newton Telescope (INT) for a study of supernova-hosting galaxies \citep{and08}.
+. Phe recession velocities obtained rom the NASA/IPAC Extragalactic Database (NED) for he central galaxies cliscussed in the present paper lic in the range 1500 - 6000 + and all had. candidate satellites hat were apparent in the continuume-subtracted Ho images.," The recession velocities obtained from the NASA/IPAC Extragalactic Database (NED) for the central galaxies discussed in the present paper lie in the range 1500 - 6000 $^{-1}$, and all had candidate satellites that were apparent in the continuum-subtracted $\alpha$ images."
+ An example galaxy anc satellite (rom thisimaging data set is shown in Fig. 1:, An example galaxy and satellite from thisimaging data set is shown in Fig. \ref{imagefig};
+ here the satellite can been seen in the x»tom left hand corner of both the #-band and Lla images., here the satellite can been seen in the bottom left hand corner of both the $R$ -band and $\alpha$ images.
+ The optical galaxy spectra presented in this paper were observed. between the 4th and Sth January 2009 with the INT on La Palma., The optical galaxy spectra presented in this paper were observed between the 4th and 8th January 2009 with the INT on La Palma.
+ The observations were made using the Intermediate. Dispersion Spectrograph (LDS). with the 1160011. erating. the 235-mm camera and the EEVLO CCD.," The observations were made using the Intermediate Dispersion Spectrograph (IDS), with the R600R grating, the 235-mm camera and the EEV10 CCD."
+ The spectral range was - to ensure the Ia emission line lies within the central unvignetted region of the CCD., The spectral range was - to ensure the $\alpha$ emission line lies within the central unvignetted region of the CCD.
+ The slit width used: was 2.0. areseconcs and was positioned. centrally on the visual peak of cach central and satellite galaxy observed. using the image on the INT acquisition TV. sereen. with a slit position angle of O° (N-S orientation) for all spectra.," The slit width used was 2.0 arcseconds and was positioned centrally on the visual peak of each central and satellite galaxy observed, using the image on the INT acquisition TV screen, with a slit position angle of $^{\circ}$ (N-S orientation) for all spectra."
+ Waveleneth calibration was achieved by taking neon lamp spectra throughout the night alter cach ealaxy observation., Wavelength calibration was achieved by taking neon lamp spectra throughout the night after each galaxy observation.
+ Eleven host ealaxies and twelve candidate satellites were observed., Eleven host galaxies and twelve candidate satellites were observed.
+ Flux calibrations were not required and therefore. standard: stars were not observed., Flux calibrations were not required and therefore standard stars were not observed.
+ Integration. times varied. between 300 and. 1200 seconds., Integration times varied between 300 and 1200 seconds.
+ “Phe frames were debiased ancl Dat-fielded. cosmic ravs and sky lines were removed. and the extracted spectra were wavelenegth-calibrated. using Séerlink software.," The frames were debiased and flat-fielded, cosmic rays and sky lines were removed, and the extracted spectra were wavelength-calibrated, using $Starlink$ software."
+ The galaxy host and satellite spectrum: pairs are shown in Fig. 2.., The galaxy host and satellite spectrum pairs are shown in Fig. \ref{specfig}.
+ In each case. the satellite spectrum is shown as the solid line. and the satellite spectrum is olfset. vertically by an arbitrary amount for ease of comparison.," In each case, the satellite spectrum is shown as the solid line, and the satellite spectrum is offset vertically by an arbitrary amount for ease of comparison."
+ The purpose of these spectra is to confirm the strongest line as Lla by identifying other lines in the spectra. where available: aud to ect recession velocities from. accurate determination of wavelengths of emission. lines.," The purpose of these spectra is to confirm the strongest line as $\alpha$ by identifying other lines in the spectra, where available; and to get recession velocities from accurate determination of wavelengths of emission lines."
+ The latter were measured by fitting a Gaussian profile to each wavelength-calibrated spectrum., The latter were measured by fitting a Gaussian profile to each wavelength-calibrated spectrum.
+ These were converted. to. recession velocities assuming the rest wavelength of Lla to beA., These were converted to recession velocities assuming the rest wavelength of $\alpha$ to be.
+. Uncertainties in the neon are calibration had rms values between andA., Uncertainties in the neon arc calibration had rms values between and.
+. Waveleneth uncertainties were further quantified by repeated: Gaussian fits to cach Ha emission line. the estimated error on each fit varied between and and. where there were multiple spectra of à galaxy. the central wavelength of the peak typically varied by toA.," Wavelength uncertainties were further quantified by repeated Gaussian fits to each $\alpha$ emission line, the estimated error on each fit varied between and and, where there were multiple spectra of a galaxy, the central wavelength of the peak typically varied by to."
+ Uncertainties in incliviclual central wavelengths were taken as the mean error determined from the Gaussian fits. as this was at least a [actor of 2 larger than any other source of uncertainty and it implicitly incorporates the ellects of varving signal-to-noise ratio from pixel to pixel.," Uncertainties in individual central wavelengths were taken as the mean error determined from the Gaussian fits, as this was at least a factor of 2 larger than any other source of uncertainty and it implicitly incorporates the effects of varying signal-to-noise ratio from pixel to pixel."
+ Projected. separations between hosts and. satellite galaxies were calculated by assuming both to lie at. distances corrected for Virgo Infall only. taken from NED.," Projected separations between hosts and satellite galaxies were calculated by assuming both to lie at distances corrected for Virgo Infall only, taken from NED."
+ The results are shown in Table 1.. which lists central galaxy. catalogued names. positions. satellite and. host. projected. separations. host and satellite heliocentric recession. velocities ancl velocity cilferences (with. the velocities in columns 6 to 8 all being derived from the data presented here). and. in the final column estimates of the central galaxies’ circular rotation velocities (discussed below).," The results are shown in Table \ref{table2}, which lists central galaxy catalogued names, positions, satellite and host projected separations, host and satellite heliocentric recession velocities and velocity differences (with the velocities in columns 6 to 8 all being derived from the data presented here), and in the final column estimates of the central galaxies' circular rotation velocities (discussed below)."
+ Phe recession velocities measured from our spectroscopy (col., The recession velocities measured from our spectroscopy (col.
+ 6) can be seen to be in good agreement with the heliocentric recession velocities taken from NED (col., 6) can be seen to be in good agreement with the heliocentric recession velocities taken from NED (col.
+ 5). with maximum cilferences within afew tens of kms +.," 5), with maximum differences within a few tens of km $^{-1}$ ."
+ While these differences are larger than the formal errors in some cases. they are in every case smaller," While these differences are larger than the formal errors in some cases, they are in every case smaller"
+This unresolved radio source lies outside the ield of the optical images.,This unresolved radio source lies outside the field of the optical images.
+ This unresolved radio source lies outside the ield of the optical images., This unresolved radio source lies outside the field of the optical images.
+ «o optical counterpart is seen at the ocation of this radio source., No optical counterpart is seen at the location of this radio source.
+ This unresolved radio source lies outside the ield of the optical images., This unresolved radio source lies outside the field of the optical images.
+ A low redshift) spiral ὃνgalaxy is. clearly he source of this extended: radio emission., A low redshift spiral galaxy is clearly the source of this extended radio emission.
+ Although the ealaxv has not been observed. spectroscopically. its bright magnitude suggests a low redshift. implving that the radio emission. is likely το be associated. with on-going star formation.," Although the galaxy has not been observed spectroscopically, its bright magnitude suggests a low redshift, implying that the radio emission is likely to be associated with on-going star formation."
+ This unresolved radio source Lies outside the field of the optical images., This unresolved radio source lies outside the field of the optical images.
+ This radio source is possibly resolved. and lies at the very edge of the optical image., This radio source is possibly resolved and lies at the very edge of the optical image.
+ An optical galaxy yartially seen at the edgee of the imageὃν appears to. be associated. with it., An optical galaxy partially seen at the edge of the image appears to be associated with it.
+ Extending nearly half an areminute on the sky. this radio source is clearly brightest at its centre and xobablv has EIULE structure. although an FRU structure with a bright radio core cannot be excluded.," Extending nearly half an arcminute on the sky, this radio source is clearly brightest at its centre and probably has I structure, although an II structure with a bright radio core cannot be excluded."
+ The central radio region is highly polarised. with the electric field vectors ving parallel to the extended: radio structure.," The central radio region is highly polarised, with the electric field vectors lying parallel to the extended radio structure."
+ The radio source lies outside the field of the optical images., The radio source lies outside the field of the optical images.
+ This unresolved radio source lies outside the ieldo 16 optical images., This unresolved radio source lies outside the field of the optical images.
+ This unresolved radio source lies outside the ieldo 16 optical images., This unresolved radio source lies outside the field of the optical images.
+ This unresolved radio source Ies outside the ield of the optical images., This unresolved radio source lies outside the field of the optical images.
+ A large cdilluse area of radio mission. extending an arcminute or more. is detected at this location.," A large diffuse area of radio emission, extending an arcminute or more, is detected at this location."
+ Lt lies outside of the optical frames. and no luminous nearby ealaxies are seen in the digitized skv catalogues at this location. so its nature remains unclear.," It lies outside of the optical frames, and no luminous nearby galaxies are seen in the digitized sky catalogues at this location, so its nature remains unclear."
+ One possibility again is that this is associated with a cluster relic radio source: although it is found. much further from the centre of the, One possibility again is that this is associated with a cluster relic radio source; although it is found much further from the centre of the
+The fitted parallax signal. with the best-fitting proper notion removed. is shown in Fieure 2..,"The fitted parallax signal, with the best-fitting proper motion removed, is shown in Figure \ref{fig:par_signal}."
+ The \? value of the fit is 3.2. which. with 9 deerces of yeedonm. maplies a >955 --coufideuce result.," The $\chi^2$ value of the fit is 3.2, which, with 9 degrees of freedom, implies a $>95$ -confidence result."
+ A bootstrap data-strippiug analysis. as detailed by Chatterjeeetal. (2009).. confirmed the validity of these results. ceiving uediaui values differing by less than lo from those ound bv the SVD fit. albeit with sliebitlv larger error wars owing to the data stripping.," A bootstrap data-stripping analysis, as detailed by \citet{Cha09}, confirmed the validity of these results, giving median values differing by less than $1\sigma$ from those found by the SVD fit, albeit with slightly larger error bars owing to the data stripping."
+ Our fitted parallax corresponds to a source distance of 2.39+O.1 Lkispe., Our fitted parallax corresponds to a source distance of $2.39\pm0.14$ kpc.
+ Qur fitted source distance is siguificauth closer thanthe best previous distance cstimate of LONE3 kkpe (Jouker&Nelemanus2001)., Our fitted source distance is significantly closer thanthe best previous distance estimate of $4.0^{+2.0}_{-1.2}$ kpc \citep{Jon04}.
+. Towever. that estimate was derived using au assuned interstellar extinction of Ay=3.3. and did not take into account uucertaiuties in the extinction.," However, that estimate was derived using an assumed interstellar extinction of $A_V=3.3$, and did not take into account uncertainties in the extinction."
+ IIvnesetal.(2009) fitted the full iiultiwvaveleugth spectral energy distribution with a template spectrum for a WO IV star and found a formal best fit reddening of Ay-=L0., \citet{Hyn09} fitted the full multiwavelength spectral energy distribution with a template spectrum for a K0 IV star and found a formal best fit reddening of $A_V=4.0$.
+ Indeed. we find that inercasing ly to close to LO brings the estimate of Jonker&Nelemanus(2001) down iuto agreement with ow results.," Indeed, we find that increasing $A_V$ to close to 4.0 brings the estimate of \citet{Jon04} down into agreement with our results."
+ This lighlehts the iportauce of using accurate extinction values when estimating source distances., This highlights the importance of using accurate extinction values when estimating source distances.
+ Caven the typical uncertainties iu the extinction towards black hole soft N-rav trausicuts (Jouker&Nelemaus2001).. we couchide that this issuc is likely to affect the majority of such svstems. rendering their distance estimates also uncertain.," Given the typical uncertainties in the extinction towards black hole soft X-ray transients \citep{Jon04}, we conclude that this issue is likely to affect the majority of such systems, rendering their distance estimates also uncertain."
+ Our measured distance is significantly closer than tle LOkkpe comunonly assumed for the source (c.e..Corbeletal. 2008).. reducing its lminosity bv a factor of 2.22.8.," Our measured distance is significantly closer than the kpc commonly assumed for the source \citep[e.g.,][]{Jon04,Gal05,Bra07,Cor08}, reducing its luminosity by a factor of 2.2–2.8."
+ The quiescent kkeV. unabsorbed flux of LOS10!l 78s d. derived using a power-law iodel for the XN-rav spoectiun (Bradleyetal.2007).. nuplies a quiescent Lbhuninositvof ~T« lo ," The quiescent keV unabsorbed flux of $1.08\times10^{-12}$ $^{-2}$ $^{-1}$, derived using a power-law model for the X-ray spectrum \citep{Bra07}, implies a quiescent luminosityof $\sim 7\times10^{32}$ $^{-1}$."
+Mako(1989) aimeasured a maxim flux of Crab withGinga diving the 1989 outburst of VIOL Cre. upline à κο luminosity of 7.0)<107 +. which is of order 0.5Lpaqq for the black hole mass of 12+24M.. derived for οι Cre (Shalibazetal.1991).," \citet{Mak89} measured a maximum flux of Crab with during the 1989 outburst of V404 Cyg, implying a keV luminosity of $7.0\times10^{38}$ $^{-1}$, which is of order $0.5L_{\rm
+Edd}$ for the black hole mass of $12\pm2M_{\odot}$ derived for V404 Cyg \citep{Sha94}."
+. Thus the system was not super-Eddingtou during the 1989 outburst., Thus the system was not super-Eddington during the 1989 outburst.
+ Furtherinore. in reducing the source radio and N-ray hunuduosifies. our new distance will reduce the scatter in the radio/N-rav correlatioBoe (Calloetal.2003).. bringing the points measured fo V[01 Cyve into better aliguiment with those of CN. 339-," Furthermore, in reducing the source radio and X-ray luminosities, our new distance will reduce the scatter in the radio/X-ray correlation \citep{Gal03}, , bringing the points measured for V404 Cyg into better alignment with those of GX 339-4."
+ With a amore accurate distance ancl proper motion for the source. aud the updated Galactic rotation paraiacters (Ro—Sd kkpe. O.=251 1) given by Reidetal.(2009). (butnoteMcMillan&Binney2009)... we can revisit the analvsis of Millez-Jouesetal.(2009). to derive amore accurate peculiar velocity of 39.935.5 ως|.," With a more accurate distance and proper motion for the source, and the updated Galactic rotation parameters $R_{\odot}=8.4$ kpc, $\Theta_{\odot}=254$ $^{-1}$ ) given by \citet[][]{Rei09}
+ \citep[but note][]{McM09}, we can revisit the analysis of \citet{Mil09} to derive a more accurate peculiar velocity of $39.9\pm5.5$ $^{-1}$."
+ This error bar now includes the uncertainties iu both the distance aud in all three space velocity compoucuts., This error bar now includes the uncertainties in both the distance and in all three space velocity components.
+ This new value is significantly lower than the best-fittiug peculium velocity of 3! derived for a Dkkpe distance CMillea-Jonesetal.2009).. aud cau casily be achieved via a Blaaww kick. such that no asvuunetric supernova kick is required.," This new value is significantly lower than the best-fitting peculiar velocity of $^{-1}$ derived for a kpc distance \citep{Mil09}, and can easily be achieved via a Blaauw kick, such that no asymmetric supernova kick is required."
+" The componcut of the »eculiar. velocity in the Galactic Plane is 39.6)kinss+, which. compared to the expected velocity dispersion in he Galactic Plane of (Miguard. 2000) for he likely Fü-E5 progenitor of the donor star. is a 2.10 result. nuplviug a probability of only 0.038 that the eculiar velocity is a result of Galactic velocity dispersion (formoredetails.seeMiller-Jonues.etal.2009)."," The component of the peculiar velocity in the Galactic Plane is $^{-1}$, which, compared to the expected velocity dispersion in the Galactic Plane of $^{-1}$ \citep{Mig00}
+ for the likely F0-F5 progenitor of the donor star, is a $2.1\sigma$ result, implying a probability of only 0.038 that the peculiar velocity is a result of Galactic velocity dispersion \citep[for more
+ details, see][]{Mil09}."
+. We herefore fiud it most probable that the peculiar velocity arises frou a natal supernova kick. the magnitude of which is consistent with recoil due to mass loss (a Blaaiw sick). with any additional asviunietrie kick beiug sinall.," We therefore find it most probable that the peculiar velocity arises from a natal supernova kick, the magnitude of which is consistent with recoil due to mass loss (a Blaauw kick), with any additional asymmetric kick being small."
+ Aliller-Jouesetal.(2009)— placed an upper liuüt of δαν on the source size. correspouding to a physical size of < 3AAU at our new distance.," \citet{Mil09} placed an upper limit of mas on the source size, corresponding to a physical size of $<3$ AU at our new distance."
+ While we did not resolve the source. our lighest-resolution data. taken with a elobal VLBI array at GGIIz. constrain the source size to «0.6 1unas. a plivsical size of < LLAAT.," While we did not resolve the source, our highest-resolution data, taken with a global VLBI array at GHz, constrain the source size to $<0.6$ mas, a physical size of $<1.4$ AU."
+ We can therefore place a lower limit of 10 KK on the brightuess temperature of the source., We can therefore place a lower limit of $10^6$ K on the brightness temperature of the source.
+" From the high brightuess temperature. the observed fat radio spectruni (Galloetal.2005) characteristic of a steady, partially self-absorbed conical outflow (Blaudford&Kóuiel 1979).. and the location of the source on the radio/X-rav correlation of Calloetal.(2003).. at the lieh-huuinositv eud of which jets lave been directly resolved (Stirlingetal.2001).. we inter that the observed radio cluission is likely to arise from compact. steady svuchrotron-cuutting jets."," From the high brightness temperature, the observed flat radio spectrum \citep{Gal05} characteristic of a steady, partially self-absorbed conical outflow \citep{Bla79}, and the location of the source on the radio/X-ray correlation of \citet{Gal03}, at the high-luminosity end of which jets have been directly resolved \citep{Sti01}, we infer that the observed radio emission is likely to arise from compact, steady synchrotron-emitting jets."
+" Our upper limit on the jet size canbe compared to that of the resolved jets in Cveuus X-1 (Stirlingetal. 2001).. with an augular size of nuuas at δια,"," Our upper limit on the jet size canbe compared to that of the resolved jets in Cygnus X-1 \citep{Sti01}, , with an angular size of mas at GHz."
+ Since the jet size scales inversely with observing frequency (Blandford&Ixónigl 1979).. this nuplies Sumas at GGIIz. for a plivsical scale of 10(7/2kpc) AAT.," Since the jet size scales inversely with observing frequency \citep{Bla79}, , this implies mas at GHz, for a physical scale of $10(d/2{\rm kpc})$ AU."
+ Since, Since
+we carry out realistic CMD simulations with and without point sources. considering the characteristics of the Planck experiment at 70 and 100 Gilz and perform the non-Caussianity detection by filtering with the MEINE.,"we carry out realistic CMB simulations with and without point sources, considering the characteristics of the Planck experiment at 70 and 100 GHz and perform the non-Gaussianity detection by filtering with the MHWF."
+ In g4. we obtain the third and fourth order cumulants produced by point sources by means of a fit of the cumulants of wavelet filtered CAIB maps and we also determine the parameters à and zl.," In 4, we obtain the third and fourth order cumulants produced by point sources by means of a fit of the cumulants of wavelet filtered CMB maps and we also determine the parameters $\alpha$ and $A$."
+ Finally in 85.r we draw the main conclusions of the paper.," Finally in 5, we draw the main conclusions of the paper."
+ Skewness and. kurtosis provide us with suitable tests for determining whether certain data come from a Caussian distribution or not., Skewness and kurtosis provide us with suitable tests for determining whether certain data come from a Gaussian distribution or not.
+ Phe skewness. s. is defined as where jes is the third central moment of the distribution and σ its rms deviation. the skewness is zero for a Gaussian distribution.," The skewness, $\textit{s}$ , is defined as where $\mu_3$ is the third central moment of the distribution and $\sigma$ its rms deviation, the skewness is zero for a Gaussian distribution."
+ The kurtosis.&.. is the fourth central moment divided by the squared. variance and its value is 3 in the Gaussian case.," The kurtosis, is the fourth central moment divided by the squared variance and its value is 3 in the Gaussian case."
+ We are interested now in calculating the skewness and kurtosis due to point sources ina wavelet filtered map from these quantities in the original observed map., We are interested now in calculating the skewness and kurtosis due to point sources in a wavelet filtered map from these quantities in the original observed map.
+ Then we consider a pixelized Mat 2D. image including. CMD. Galactic foregrounds. instrumental. noise and extragalactic point sources: these four components are independent anc are added up to form the image: whereas the CAIB is Gaussian ancl the noise is considered as non-uniform but Caussian-clistributec at cach pixel. the other two are non-Gaussian.," Then we consider a pixelized flat 2D image including CMB, Galactic foregrounds, instrumental noise and extragalactic point sources; these four components are independent and are added up to form the image; whereas the CMB is Gaussian and the noise is considered as non-uniform but Gaussian-distributed at each pixel, the other two are non-Gaussian."
+ Skewness ancl kurtosis can be obtained from the cumulants of the probability distribution function (pdf) pr)., Skewness and kurtosis can be obtained from the cumulants of the probability distribution function (pdf) $ p(x)$.
+ The characteristic function O(i) is defined as the Fourier transform of the pdf and the cumulant function. is A(u)=log(ó(u))., The characteristic function $ \phi (u)$ is defined as the Fourier transform of the pdf and the cumulant function is $ K(u)= log(\phi(u))$.
+ The cumulants of order m. νι can be caleulated as derivatives of the cumulant function The second. order cumulant is the variance. the third order cumulant is the third. central moment yr; and the fourth order cumulant is fp30.," The cumulants of order m, $K_m$, can be calculated as derivatives of the cumulant function The second order cumulant is the variance, the third order cumulant is the third central moment $\mu_3$ and the fourth order cumulant is $\mu_4-3\sigma^4$."
+ The advantage of using the eumulants is that when the different components of the data are independent the cumulant of the sum is the sum of the cumulants., The advantage of using the cumulants is that when the different components of the data are independent the cumulant of the sum is the sum of the cumulants.
+ So. the cumulant of any order of the sum of CALB. noise. point sources and other foregrounds is the sum of the cumulants of the components.," So, the cumulant of any order of the sum of CMB, noise, point sources and other foregrounds is the sum of the cumulants of the components."
+ Ina typical CMD experiment. the image is filtered with a Caussian-shaped instrumental beam. this means that the map is convolved with a Gaussian function.," In a typical CMB experiment, the image is filtered with a Gaussian-shaped instrumental beam, this means that the map is convolved with a Gaussian function."
+ This convolution changes the skewness ancl kurtosis of the original map., This convolution changes the skewness and kurtosis of the original map.
+ Alter the convolution. the temperature of a pixel is a linear combination of the temperatures of the neighboring pixels in the unfiltered map.," After the convolution, the temperature of a pixel is a linear combination of the temperatures of the neighboring pixels in the unfiltered map."
+ For point sources. which are not correlated (this is true for the frequencies. 70 ancl LOO Cillz. considered in this work. as can be seen in González-Nuevoetal.(2005))). the eumulants of this linear combination can then be written as a linear combination of the cumulants of the original map.," For point sources, which are not correlated (this is true for the frequencies, 70 and 100 GHz, considered in this work, as can be seen in \cite{gon05}) ), the cumulants of this linear combination can then be written as a linear combination of the cumulants of the original map."
+" If we are filtering with afunction sr). the cumulant of order m. AS.nt of the liltered image can be written as where A,, is the original cumulant of order m in the pixelized. image and 4, the pixel area."," If we are filtering with afunction $\varphi(x)$, the cumulant of order m, $K^f_m$, of the filtered image can be written as where $K_m$ is the original cumulant of order m in the pixelized image and $A_p$ the pixel area."
+" From this formula. we can obtain the cumulants of the Gaussian filtered image AY,nn from those of the original one."," From this formula, we can obtain the cumulants of the Gaussian filtered image $ K_{m}^{g}$ from those of the original one."
+ where / is the pixel size anc 5 the bean. dispersion. see eq. (," where $l$ is the pixel size and $\gamma$ the beam dispersion, see eq. ("
+4).,4).
+ We are going to filter now the image with the first three members of the MIIWE. the whole process amounts to filtering the original image with a convolution of the Gaussian instrumental beam and the correspondingwavelets.," We are going to filter now the image with the first three members of the MHWF, the whole process amounts to filtering the original image with a convolution of the Gaussian instrumental beam and the correspondingwavelets."
+ We could. then apply Eq.12) with qr) this convolution., We could then apply Eq.(12) with $\varphi(\vec{x})$ this convolution.
+" Finally. we can relate the cumulants A,nn of the wavelet filtered image to the same quantities A, in the original pixelized image."," Finally, we can relate the cumulants $ K_{m}^{w} $ of the wavelet filtered image to the same quantities $ K_m$ in the original pixelized image."
+ We write the formula for a &eneral cumulant and a general member of index n of the MIINE in this formula 3=ές. Rois the wavelet scale (see eq.," We write the formula for a general cumulant and a general member of index n of the MHWF in this formula $\beta=R/\gamma$, R is the wavelet scale (see eq."
+ 5) and the coellicients ων are calculated as This last expression can be casily computed. for the first members of the MIINE and for m=3.4.," 5) and the coefficients $C_{mn}$ are calculated as This last expression can be easily computed for the first members of the MHWF and for m=3,4."
+" When we consider the skewness of the total wavelet filtered image (CMD|]point sources|Galactic foregrounds|noise). it can be obtained as the total third order cumulant divided. by a?u with o, the total rnis deviation."," When we consider the skewness of the total wavelet filtered image (CMB+point sources+Galactic foregrounds+noise), it can be obtained as the total third order cumulant divided by $\sigma_w^3$ with $\sigma_w$ the total rms deviation."
+ The contribution of point sources to the total skewness of the wavelet filtered map is where A is the third order cumulant (due to point sources) of the wavelet filtered map., The contribution of point sources to the total skewness of the wavelet filtered map is where $ K^w_3$ is the third order cumulant (due to point sources) of the wavelet filtered map.
+" The contribution of point sources to the total skewness of the Gaussian filtered map is where A7 is the third order cumulant (due to point sources) of the Gaussian filtered map and e, the total rms deviation of the Gaussian filtered maps.", The contribution of point sources to the total skewness of the Gaussian filtered map is where $ K^g_3$ is the third order cumulant (due to point sources) of the Gaussian filtered map and $\sigma_g$ the total rms deviation of the Gaussian filtered maps.
+ Finally. by using Eqs. (," Finally, by using Eqs. ("
+"13)and (14) and the definition of the amplification. À. we can relate ὁμ and s,.","13)and (14) and the definition of the amplification, $\lambda$ , we can relate $\hat{s}_w$ and $\hat{s}_g$ ."
+ We obtain the following formula for the MEINT for the MIIW2. we obtain a similar Formula," We obtain the following formula for the MHW1 for the MHW2, we obtain a similar formula"
+"an increase in X-ray reflection is οιντος, in outburst than in quiescence.",an increase in X-ray reflection is required in outburst than in quiescence.
+ In this paper we present X-ray spectroscopic observations which caught the cwarf nova Z Cam in outburst in 1995 and during a transition from quiescence to outburst in LOOT (see Fie. 1))., In this paper we present X-ray spectroscopic observations which caught the dwarf nova Z Cam in outburst in 1995 and during a transition from quiescence to outburst in 1997 (see Fig. \ref{fig:longoptical}) ).
+ Ixnigge (1997) observed the same outburst with the Hopkins. Ultraviolet Telescope on the Astro-2 mission., Knigge (1997) observed the same outburst with the Hopkins Ultraviolet Telescope on the Astro-2 mission.
+" ""They apply an accretion clisk wind mocel to fit absorption features in the UV spectrum of Z Cam. and conclude that a dense. slow-moving disk wind transition region acts as an absorbing medium with a temperature of 101."," They apply an accretion disk wind model to fit absorption features in the UV spectrum of Z Cam, and conclude that a dense, slow-moving disk wind transition region acts as an absorbing medium with a temperature of $\times10^{4}$."
+ Modelling of disk. winds has been carried out by Proga (1998) who find that radiation driven wines are intrinsically unsteacdy. with large density and velocity variations.," Modelling of disk winds has been carried out by Proga (1998) who find that radiation driven winds are intrinsically unsteady, with large density and velocity variations."
+ The dense low-velocity low. component of the wind is confined to angles below of the equatorial plane., The dense low-velocity flow component of the wind is confined to angles below of the equatorial plane.
+ Z Cam has an inclination of 57-11 (Shafter 1983) ancl so the line of sight to Z Cam should pass through this zone., Z Cam has an inclination of $\pm$ (Shafter 1983) and so the line of sight to Z Cam should pass through this zone.
+ Z Cam has been observed in both outburst and during the optical rise to another outburst using the Japanese satellite., Z Cam has been observed in both outburst and during the optical rise to another outburst using the Japanese satellite.
+ has four X-ray telescopes and detectors: two X-ray. CCD cameras (Solid-state Imaging Spectrometers. or SIS) and two gas scintillation imaging proportional counters (Cas Imaging Spectromoeters. or GIS) CTanaka. Inoue 11ο 1994).," has four X-ray telescopes and detectors: two X-ray CCD cameras (Solid-state Imaging Spectrometers, or SIS) and two gas scintillation imaging proportional counters (Gas Imaging Spectrometers, or GIS) (Tanaka, Inoue Holt 1994)."
+ Phe SIS detectors have an energy range of kkeV whilst the CUS have an energy range of kkeV. As shown in Figs 1. 2.. Z Cam was observed on two occasions: once during outburst between S March. to 12 March. 1995 (JD. 2449783-2449788). and once during a transition as Z Cam optically brightened to another outburst state on 12 April. LOOT (JD 2450550).," The SIS detectors have an energy range of keV whilst the GIS have an energy range of keV. As shown in Figs \ref{fig:longoptical}~ \ref{fig:opticalxray}, Z Cam was observed on two occasions: once during outburst between 8 March to 12 March, 1995 (JD 2449783-2449788), and once during a transition as Z Cam optically brightened to another outburst state on 12 April, 1997 (JD 2450550)."
+ To maximise the amount of data. and to avoid unnecessary gaps in the light curve. loose screening criteria have been applied. compared to the standard REW2 screening.," To maximise the amount of data, and to avoid unnecessary gaps in the light curve, loose screening criteria have been applied compared to the standard REV2 screening."
+ We reject. data. from all instruments in. both observations when the telescope vas cleviated more than away from the source., We reject data from all instruments in both observations when the telescope has deviated more than away from the source.
+ The elevation. angle between the Earth's limb and. satellite »ointing direction has been limited to greater than with he SIS and 5with the GIS. the bright Earth elevation angle is limited to greater. than 15and. the. radiation xekeround monitor is restricted to less than 250 counts vcr second.," The elevation angle between the Earth's limb and satellite pointing direction has been limited to greater than with the SIS and with the GIS, the bright Earth elevation angle is limited to greater than and the radiation background monitor is restricted to less than 250 counts per second."
+" ""The minimum value. of the cut-olf. rigidity of the ecomagnetic field. throughout both observations is approximately 4 GeV/c. which is high enough not to require an additional screening criterion."," The minimum value of the cut-off rigidity of the geomagnetic field throughout both observations is approximately 4 GeV/c, which is high enough not to require an additional screening criterion."
+ Time [filters have been applied to cut out the lew remaining periods of very high background. during the transitional observation. although this amounts to less than 300 seconds in all instruments.," Time filters have been applied to cut out the few remaining periods of very high background during the transitional observation, although this amounts to less than 300 seconds in all instruments."
+ Circular source extraction regions have been used of 6 aremin diameter in the GIS and 4 aremin in the SIS., Circular source extraction regions have been used of 6 arcmin diameter in the GIS and 4 arcmin in the SIS.
+ The background has been taken from the remaining area of the same chip in the SIS. and an annulus of outer radius. 18 aremin around the source for the CIS.," The background has been taken from the remaining area of the same chip in the SIS, and an annulus of outer radius 18 arcmin around the source for the GIS."
+ The outburst observation exposure times after screening are SOkks kks for SISO SISI and kks for each CIS., The outburst observation exposure times after screening are ks ks for SIS0 SIS1 and ks for each GIS.
+ For the transition observation. the exposure times are LS ks 22 ks for the two SIS. and 22 ks 23 ks for the two GIS.," For the transition observation, the exposure times are 18 ks 22 ks for the two SIS, and 22 ks 23 ks for the two GIS."
+ “Phe response matrices for cach SIS have been generated using the POOL SISILMC. and for the CIS the standard response matrices (version 0.8. November 1994) have been used.," The response matrices for each SIS have been generated using the FTOOL SISRMG, and for the GIS the standard response matrices (version 0.8, November 1994) have been used."
+ After this loose screening. the two SIS instruments and two CLS instruments were separately combined in both observations.," After this loose screening, the two SIS instruments and two GIS instruments were separately combined in both observations."
+ The outburst spectra have been grouped to à minimum. of 40 counts per bin and the transition spectra have been grouped to at least 20 counts per bin. allowing X7 fitting.," The outburst spectra have been grouped to a minimum of 40 counts per bin and the transition spectra have been grouped to at least 20 counts per bin, allowing $\chi^{2}$ fitting."
+ Figure 2. shows the optical light curve. the X-ray light curve and the X-ray hardness ratio during the two observations.," Figure \ref{fig:opticalxray} shows the optical light curve, the X-ray light curve and the X-ray hardness ratio during the two observations."
+ The top panel shows the .XAVSO optical observations. and the bottom two panels show the data over a one sixth smaller time range.," The top panel shows the AAVSO optical observations, and the bottom two panels show the data over a one sixth smaller time range."
+ The hardness ratio, The hardness ratio
+regions (ruling out ram pressure stripping as the dominant process) and neither can it proceed faster in higher densities.,regions (ruling out ram pressure stripping as the dominant process) and neither can it proceed faster in higher densities.
+ Our data are consistent with the process that turns spirals red being more likely to happen in high density regions as long as it is possible and proceeds in the same fashion in all environments., Our data are consistent with the process that turns spirals red being more likely to happen in high density regions as long as it is possible and proceeds in the same fashion in all environments.
+" Both gentle interactions, and/or strangulation mechanisms could explain these observations."," Both gentle interactions, and/or strangulation mechanisms could explain these observations."
+" The most recent semi-analytic models for galaxy formation all invoke some form of AGN feedback which inhibits star formation in the most massive, red (early-type) galaxies (Granatoetal.2004;Silk2005;Crotonal.2006;Boweret2006) in order to match the number counts and properties of these objects in the local universe."," The most recent semi-analytic models for galaxy formation all invoke some form of AGN feedback which inhibits star formation in the most massive, red (early-type) galaxies \citep{G04,S05,Cr06,Bower06} in order to match the number counts and properties of these objects in the local universe."
+" Observational evidence now exists to support this hypothesis, showing that the AGN fraction peaks in the region between the blue cloud and the red sequence (Schawinskietal.2007,2009a)."," Observational evidence now exists to support this hypothesis, showing that the AGN fraction peaks in the region between the blue cloud and the red sequence \citep{S07,ScL09}."
+. Could AGN feedback be responsible for the shutdown of star formation in the red spirals?, Could AGN feedback be responsible for the shutdown of star formation in the red spirals?
+" We test this by using emission line diagnostic diagrams (Baldwin Phillips Terlevich 1981, hereafter BPT; Veilleuxetal. 2001)) to probe the dominant source of ionisation in the red and blue spiral galaxies."," We test this by using emission line diagnostic diagrams (Baldwin Phillips Terlevich 1981, hereafter BPT; \citealt{VO87,Ke01,Ke06,CF01}) ) to probe the dominant source of ionisation in the red and blue spiral galaxies."
+" For this study, we restrict our analysis to galaxies with greater than 2σ emission line detection for all four lines used in the BPT diagram (i.e. [OI], Ha, [NII] and Hf)."," For this study, we restrict our analysis to galaxies with greater than $\sigma$ emission line detection for all four lines used in the BPT diagram (i.e. [OIII], $\alpha$, [NII] and $\beta$ )."
+" Likewise, in order to reduce the effect of aperture bias of the SDSS 3” fibre on the BPT diagram, we shift the lower redshift range of our sample up from z>0.03 to z>0.05."," Likewise, in order to reduce the effect of aperture bias of the SDSS $3\arcsec$ fibre on the BPT diagram, we shift the lower redshift range of our sample up from $z>0.03$ to $z>0.05$."
+" This reduces the sample size to 181 (60%)) of the red spirals, and 3462 (66%)) of the blue spirals."," This reduces the sample size to 181 ) of the red spirals, and 3462 ) of the blue spirals."
+ Standard emission line diagnostic diagrams (e.g.Kauff-mannetal.2003b;Schawinski2007) are used to divide galaxies whose budget of ionising photons is dominated by current star formation (i.e. ionisation from OB associations) and those whose inter-stellar medium is excited by other processes.," Standard emission line diagnostic diagrams \citep[e.g.][]{K03b,S07} are used to divide galaxies whose budget of ionising photons is dominated by current star formation (i.e. ionisation from OB associations) and those whose inter-stellar medium is excited by other processes."
+" These other processes include AGN, but also shocks and evolved stellar populations such as post-asymptotic giant branch (pAGB) stars, extreme horizontal branch stars, white dwarfs (Stasiriskaetal.2008;Sarzi 2009)."," These other processes include AGN, but also shocks and evolved stellar populations such as post-asymptotic giant branch (pAGB) stars, extreme horizontal branch stars, white dwarfs \citep{St08,Sarzi09}."
+". Observationally, the area between these two regions is populated by transition or composite objects where the total contribution to the ionising budget from star formation and other processes is roughly comparable."," Observationally, the area between these two regions is populated by transition or composite objects where the total contribution to the ionising budget from star formation and other processes is roughly comparable."
+" Kewleyetal.(2001) give a prescription based on theoretical modelling of starburst galaxies which indicates the region of the diagram which can be explained by starburst emission (to the lower left of the line), while Kauffmannetal.(2003b) made an empirical fit to the observed separation between the populations in SDSS galaxies, which is often used as a lower limit for the location of AGN."," \citet{Ke01} give a prescription based on theoretical modelling of starburst galaxies which indicates the region of the diagram which can be explained by starburst emission (to the lower left of the line), while \citet{K03b} made an empirical fit to the observed separation between the populations in SDSS galaxies, which is often used as a lower limit for the location of AGN."
+" Here as a short-hand, we will call objects to the upper right of the Kewleyetal.(2001) line “AGN” (even if other processes may be more important in some objects), or sometimes ”Seyfert+LINER” (as explained below), objects below the"," Here as a short-hand, we will call objects to the upper right of the \citet{Ke01} line “AGN"" (even if other processes may be more important in some objects), or sometimes ""Seyfert+LINER"" (as explained below), objects below the"
+Active stars are known to show strong emission in the lline cores at 3934.8 and 3969.7 citep[see][foraprofounddiscussion]Linsky1980..,Active stars are known to show strong emission in the line cores at 3934.8 and 3969.7 \\citep[see][for a profound discussion]{Linsky1980}.
+ This is also true forCoRoT-2A;; we show its llines in Fig. 2.., This is also true for; we show its lines in Fig. \ref{fig:CaHK}.
+" In late-type stars the width, Wo, of the emission cores seen in the the llines is mainly sensitive to the value of logg, and, hence,K to the mass and radius of the star, but insensitive to the effective temperature and metallicity."," In late-type stars the width, $W_0$, of the emission cores seen in the the lines is mainly sensitive to the value of $\log g$, and, hence, to the mass and radius of the star, but insensitive to the effective temperature and metallicity."
+" If calibrated appropriately, the width can be used as a rough estimator of the absolute luminosity of a star (?)."," If calibrated appropriately, the width can be used as a rough estimator of the absolute luminosity of a star \citep{WilsonBappu1957}."
+ We used the recent calibration of ? together with CoRoT-2A'ss apparent magnitude corrected for interstellar extinction (see Sect. 4.3)), We used the recent calibration of \citet{Pace2003} together with s apparent magnitude corrected for interstellar extinction (see Sect. \ref{sec:distanceToCorot2}) )
+ to obtain a distance estimate for the ssystem., to obtain a distance estimate for the system.
+" ? find no significant effect of rotational and instrumental broadening within their sample, which comprises stars with vsini<14 s*!."," \citeauthor{Pace2003} find no significant effect of rotational and instrumental broadening within their sample, which comprises stars with $v\sin{i}<14$ ."
+". Neglecting broadening effects, we obtained a distance estimate of 190770 pc."," Neglecting broadening effects, we obtained a distance estimate of $190_{+60}^{-50}$ pc."
+" When both broadening mechanisms were taken into account by a quadratic correction to Wo as described by ?,, the distance estimate reduced to 140130 pc, which is lower but still compatible, given the errors."," When both broadening mechanisms were taken into account by a quadratic correction to $W_0$ as described by \citeauthor{Pace2003}, the distance estimate reduced to $140_{-40}^{+50}$ pc, which is lower but still compatible, given the errors."
+ The Mt. Wilson chromospheric flux index (Sww; is a popular measure of chromospheric activity., The Mt. Wilson chromospheric flux index \citep[$S_{\mathrm{MW}}$ is a popular measure of chromospheric activity.
+" To estimate Sww from our spectra, we used the calibration procedure for UVES spectra described by ?.."," To estimate $S_{\mathrm{MW}}$ from our spectra, we used the calibration procedure for UVES spectra described by \citet{Melo2006}."
+ These authors defined a proper index Sus and determined the following relation between their index and the Mt. Wilson index: (seeEq.1in?).., These authors defined a proper index $S_{\mathrm{US}}$ and determined the following relation between their index and the Mt. Wilson index: \citep[see Eq.~1 in][]{Melo2006}.
+" Inverting this relation, we obtained an Smw index of 0.479."," Inverting this relation, we obtained an $S_{\mathrm{MW}}$ index of $0.479$."
+ We proceeded by correcting our Smw index for the color dependence (?) and subtracted the expected photospheric contribution to the flux in the line cores (?).., We proceeded by correcting our $S_{\mathrm{MW}}$ index for the color dependence \citep{Rutten1984} and subtracted the expected photospheric contribution to the flux in the line cores \citep{Noyes1984}.
+" Using a value of 0.854 for the B-V color of ffrom the SIMBADdatabase’,, we obtained an emission ratio of logRy=—4.458+0.051, which agrees well with the results of ? (—4.471+ 0.0629) and ? (—4.331)."," Using a value of $0.854$ for the B-V color of from the SIMBAD, we obtained an emission ratio of $\log{R'_{\mathrm{HK}}}=-4.458\pm0.051$, which agrees well with the results of \citet{Gillon2010} $-4.471\pm0.0629$ ) and \citet{Knutson2010} $-4.331$ )."
+ The central part of the eemission core does not show the double-horned structure usually observed in the line cores of solar-like stars., The central part of the emission core does not show the double-horned structure usually observed in the line cores of solar-like stars.
+ We found no sign of self-reversal in the central part of the line core., We found no sign of self-reversal in the central part of the line core.
+" According to ?,, strong chromospheric heating through stellar activity leads to a decrease of the wavelength separation between the two Ky peaks of the emission core, so that the central Ks dip is easily obliterated by instrumental, macroscopic, and microscopic broadening effects."," According to \citet{Ayres1979}, strong chromospheric heating through stellar activity leads to a decrease of the wavelength separation between the two $_2$ peaks of the emission core, so that the central $_3$ dip is easily obliterated by instrumental, macroscopic, and microscopic broadening effects."
+" The lack of a detectable self-reversal is, therefore, an indicator of strong chromospheric heating itself."," The lack of a detectable self-reversal is, therefore, an indicator of strong chromospheric heating itself."
+" Because chromospheric heating causes an increase in the wavelength separation of the two K; minima that approximately counterbalances the mutual approach of the Κο peaks, the Wilson-Bappu width, Wo, remains basically insensitive to stellar activity."," Because chromospheric heating causes an increase in the wavelength separation of the two $_1$ minima that approximately counterbalances the mutual approach of the $_2$ peaks, the Wilson-Bappu width, $W_0$, remains basically insensitive to stellar activity."
+" Chromospheric activity can also be measured in theCar infrared triplet (IRT) at 8498, 8542, and 8662 citep[seee.g.,][]Andretta2005, Busa2007.."," Chromospheric activity can also be measured in the infrared triplet (IRT) at 8498, 8542, and 8662 \\citep[see e.g.,][]{Andretta2005, Busa2007}."
+" To correctly assess the chromospheric contribution to the line profile, the line forming process within the photosphere must be accurately modeled."," To correctly assess the chromospheric contribution to the line profile, the line forming process within the photosphere must be accurately modeled."
+ ? proposed the activity indicator Rr defined as the difference between the central line depths (or central depressions) of the observed spectrum and a rotationally-broadened NLTE model., \cite{Andretta2005} proposed the activity indicator $R_{\mathrm{IRT}}$ defined as the difference between the central line depths (or central depressions) of the observed spectrum and a rotationally-broadened NLTE model.
+" We calculated the central depression, ΚΙΝΤ. and AW, i.e., the EW of the residual lineprofile and the corresponding error estimates following the approach detailed in ?.."," We calculated the central depression, $R_{\mathrm{IRT}}$, and $\Delta W_{\mathrm{IRT}}$, i.e., the EW of the residual lineprofile and the corresponding error estimates following the approach detailed in \cite{Busa2007}."
+" In our calculations, we used LTE models synthesized via SME, which,according to ?,, suffice to approximate the photospheric spectrum for main-sequence stars of solar metallicity."," In our calculations, we used LTE models synthesized via SME, which,according to \citeauthor{Andretta2005}, , suffice to approximate the photospheric spectrum for main-sequence stars of solar metallicity."
+ The parameters thus obtained are summarized in Table 3.., The parameters thus obtained are summarized in Table \ref{tab:CoRoT2CaIRT}.
+" The CCa IRT along with our synthetic templates are shown in Fig. 3,,"," The Ca IRT along with our synthetic templates are shown in Fig. \ref{fig:CaIRT}, ,"
+ which also indicates the residuals., which also indicates the residuals.
+" ? found relationship between the logRi, activity index and ΚΙΚΤ, however,a with a large scatter."," \citeauthor{Busa2007} found a relationship between the $\log{R'_\mathrm{\mathrm{HK}}}$ activity index and $R_{\mathrm{IRT}}$, however, with a large scatter."
+" Given our value of logΚμκ=4.46, the relation predicts a value of 0.25 for Rirr, which reasonably agrees with our results."," Given our value of $\log{R'_\mathrm{HK}}=4.46$, the relation predicts a value of $0.25$ for $R_{\mathrm{IRT}}$, which reasonably agrees with our results."
+" Analyzing the emission line cores of the ccomplex, we found narrow Ky (seee.g.,?,fordesignation) and Hy, absorption features close to both central chromospheric emissionpeaks (Fig. 4))."," Analyzing the emission line cores of the complex, we found narrow $_4$ \citep[see e.g.,][for designation]{Reimers1977} and $_4$ absorption features close to both central chromospheric emissionpeaks (Fig. \ref{fig:the_wind_plot}) )."
+" We therefore attribute the features to absorption and determine a barycentric velocity of -15 kms-!, corresponding to a relative velocity of about +40 wwith respect to CoRoT-2A.", We therefore attribute the features to absorption and determine a barycentric velocity of $-15$ corresponding to a relative velocity of about $+40$ with respect to .
+". Separately considering our 24 UVES spectra observed within 6 hours and the archivedUVES spectrum taken about three yearsearlier in 2007, we found no intrinsic variability"," Separately considering our $24$ UVES spectra observed within 6 hours and the archivedUVES spectrum taken about three yearsearlier in 2007, we found no intrinsic variability"
+The relation between the midplane temperature νι and the ellective temperature is given by Llubeny(1990). as with where τς&Xj2 is the optical depth at the midplane.,The relation between the midplane temperature $T_{\rm vis}$ and the effective temperature is given by \citet{Hub90} as with where $\tau = \kappa \Sigma/2$ is the optical depth at the midplane.
+ Eq. (11)), Eq. \ref{eq:taueff}) )
+ is applicable to all cases., is applicable to all cases.
+" For pure MIU active regions without dead zones. we set X—X, and a=o4."," For pure MRI active regions without dead zones, we set $\Sigma = \Sigma_{\rm a}$ and $\alpha = \alpha_{\rm a}$."
+" For the dead zone region with ay=0. we set Na=Xo, so that Eq. (10) "," For the dead zone region with $\alpha_{\rm d} = 0$, we set $\Sigma\alpha = \Sigma_{\rm a}\alpha_{\rm a}$ so that Eq. \ref{eq:tvis}) )"
+becomes the same with Eq. (, becomes the same with Eq. (
+7) of Ciamumie(1996) for 7/YX9 L.,7) of \citet{Gam96} for $\tau \Sigma_{\rm a}/\Sigma \gg 1$.
+A detailed derivation of the first term in the square bracket of the right hand side of Eq. (11)), A detailed derivation of the first term in the square bracket of the right hand side of Eq. \ref{eq:taueff}) )
+ is elven in Wünschetal.(2006) , is given in \citet{Wun06}. .
+The opacity & is represented by Hocsal. where ο is a constant.," The opacity $\kappa$ is represented by $\kappa = \kappa_0 T^{\beta}$, where $\kappa_0$ is a constant."
+ We use & from Stepinski shown in Table 1., We use $\kappa$ from \citet{Ste98} shown in Table 1.
+ Results of the present study do not sensitively depend on opacity law., Results of the present study do not sensitively depend on opacity law.
+" The disc effective temperature irradiated by the central staris given by (Ruden&Pollack1991) where 7, and r; are the photospheric temperature and racius of the central star.", The disc effective temperature irradiated by the central staris given by \citep{Rud91} where $T_{*}$ and $r_{*}$ are the photospheric temperature and radius of the central star.
+" We adopt 7;=4000 Ix and +,—:λ) .", We adopt $T_{*} = 4000$ K and $r_{*} = 3$ $r_{\odot}$.
+ The scale height/ used in Iq. (12)), The scale height$h$ used in Eq. \ref{eq:tc}) )
+ is the photospheric disc scale height and we simply assume it to bethe same with the pressure scale height., is the photospheric disc scale height and we simply assume it to be the same with the pressure scale height.
+ After Llueso&αιοι(2005)... we assume d1Inh/dInr=9/7.," After \citet{Hue05}, we assume $d\ln{h}/d\ln{r} = 9/7$."
+ This is validated as Loxr5 Cad large radii where irraciation from the central star is usually the dominant heat source., This is validated as $T \propto r^{-3/7}$ at large radii where irradiation from the central star is usually the dominant heat source.
+" The midplane temperature. Zi. irracliatecl by the central star is calculated using the vertical temperature structure derived by Malbet&Bertout (see also Alalbet. Lachaume Monin (2001): Llere gro is the averaged cosine of theincident angle of the central star given as with. the zeroth and first. order moments. J""í and í41"". of the incident intensity at the upper most laver: Eq. (13))"," The midplane temperature, $T_{\rm irr}$, irradiated by the central star is calculated using the vertical temperature structure derived by \citet{Mal91}
+ (see also Malbet, Lachaume Monin (2001)): Here $\mu_0$ is the averaged cosine of the incident angle of the central star given as with the zeroth and first order moments, $J^0$ and $H^0$, of the incident intensity at the upper most layer: Eq. \ref{eq:tirr}) )"
+ represents that Zi; of an optically thick disc is lower than Diy by a factor of 22/3. as the upper most lavers emit half of the received eem toward the midplane.," represents that $T_{\rm irr}$ of an optically thick disc is lower than $T_{\rm irr, eff}$ by a factor of $2^{1/4}$, as the upper most layers emit half of the received energy toward the midplane."
+" In the optically thin limit (7 ). the term in the bracket dominates anc we obtain ""n--—qurun which corresponds to the equilibrium irracliatecl by the half hemisphere of the central star without any obstacle."," In the optically thin limit $\tau \rightarrow 0$ ), the third term in the bracket dominates and we obtain $T_{\rm irr}^4 = T_{*}^4(r_{*}/r)^2/8$, which corresponds to the equilibrium temperature irradiated by the half hemisphere of the central star without any obstacle."
+ The temperature at 7) should. be higher by a factor of 25/7. but this dillerence does not alleet our results as the stage with a low surface clensity is very short with photoevaporative winds.," The temperature at $\tau \rightarrow 0$ should be higher by a factor of $2^{1/4}$, but this difference does not affect our results as the stage with a low surface density is very short with photoevaporative winds."
+ The ambient temperature is simply assumed as Zi)=20 Ix. Among various types of photoevaporations. the one induced by X-ravs from the central star is most strong ancl most important for gas disc evolution in the planet formation region (Ercolano.Clarke&Drake2009:Owenctal.2010.2011).. unless clises are close to very massive external stars (Aclamsetal.2004:Mitchell&Stewart2010)..," The ambient temperature is simply assumed as $T_{\rm amb} = 20$ K. Among various types of photoevaporations, the one induced by X-rays from the central star is most strong and most important for gas disc evolution in the planet formation region \citep{Erc09, Owe10, Owe11}, unless discs are close to very massive external stars \citep{Ada04,Mit10}. ."
+ The total mass loss rate from a disc due to X-ray photoevaporation is estimated by Owenetal.(2011) as where ἐς a constant of the order of unity (we fix it to be unity) and Lx is the X-ray luminosity from the central star., The total mass loss rate from a disc due to X-ray photoevaporation is estimated by \citet{Owe11} as where $A$ is a constant of the order of unity (we fix it to be unity) and $L_{\rm X}$ is the X-ray luminosity from the central star.
+ rn sub-solar to e mass protostars. Lx is LOph 104 erg s.+ (Güdeletal..," For sub-solar to solar mass protostars, $L_{\rm X}$ is $10^{29} $ $10^{31}$ erg $^{-1}$ \citep{Gud07}."
+ Numerical simulations of Owenetal.(2010). show m2 H is roughly proportional to r (their Fig., Numerical simulations of \citet{Owe10} show that $\dot{\Sigma}_{\rm w}$ is roughly proportional to $r^{-3/2}$ (their Fig.
+ 13), 13).
+ Thus. we give the mass loss rate per unit pea AS where rin and rxoaa are the inner and outer edges of the range where mass loss occurs.," Thus, we give the mass loss rate per unit area as where $r_{\rm X, in}$ and $r_{\rm X,out}$ are the inner and outer edges of the range where mass loss occurs."
+ We take rx= 1 AU and rx= 70 AU from Owenetal.(2010).," We take $r_{\rm X,in} =$ 1 AU and $r_{\rm X,out} =$ 70 AU from \citet{Owe10}."
+. Outside riu. other mechanisms such as EUW or ΕΝ) photoevaporation (Gortietal.2000) may. work. although these elfects. are not taken into account in the present work.," Outside $r_{\rm X, out}$, other mechanisms such as EUV or FUV photoevaporation \citep{Gor09} may work, although these effects are not taken into account in the present work."
+ Results of disc evolution do not sensitively depend on r-dependence of NS. as similar results are obtained even if we vary the power-law index between -1 and -2," Results of disc evolution do not sensitively depend on $r$ -dependence of $\dot{\Sigma}_{\rm w}$, as similar results are obtained even if we vary the power-law index between -1 and -2."
+ In some test simulations. we also adopt the mass loss due to MBI disc winds (Suzukietal.2010)," In some test simulations, we also adopt the mass loss due to MRI disc winds \citep{Suz10}."
+ Phe mass loss rate of. MIU clise winds at a given radius is proportional to the surface density., The mass loss rate of MRI disc winds at a given radius is proportional to the surface density.
+" We find that gas depletes very. rapidly even near the inner οσο, and that the mass acerction rate onto the central star becomes less than 10.CAL. vet within ] Myr. as shown in Suzukietal.(2010)."," We find that gas depletes very rapidly even near the inner edge, and that the mass accretion rate onto the central star becomes less than $10^{-9} M_{\odot}$ $^{-1}$ within 1 Myr, as shown in \citet{Suz10}."
+.. This. accretion rate looks too low as compared with those for typical discs around classical T Tauri stars (Παναμάetal.1998), This accretion rate looks too low as compared with those for typical discs around classical T Tauri stars \citep{Har98}.
+ Probably. the mass loss rate of MIL disc winds is much smaller. but. this mechanism may still play an important role in disc evolution.," Probably, the mass loss rate of MRI disc winds is much smaller, but this mechanism may still play an important role in disc evolution."
+ We numerically solve Ίσα. (4)), We numerically solve Eq. \ref{eq:sigd}) )
+ using the method described in Bath ancl Pringle (1981). in which the radial grid size isproportional to Vr.," using the method described in Bath and Pringle (1981), in which the radial grid size isproportional to $\sqrt{r}$ ."
+ We use 1000 erids between 0.03 AU and 3000 AU., We use 1000 grids between 0.03 AU and 3000 AU.
+" ""Phese numbers give the inner most grid size of 0.02 AU.", These numbers give the inner most grid size of 0.02 AU.
+ At the inner boundary X is fixed to be zero. and at the outer boundary we adopt the outward mass Iux eiven bv. 3x Xe.," At the inner boundary $\Sigma$ is fixed to be zero, and at the outer boundary we adopt the outward mass flux given by $-3\pi\Sigma\nu$ ."
+ We find thatthe evolution of the total mass is nearly the same even with larger outer boundary, We find thatthe evolution of the total mass is nearly the same even with larger outer boundary
+"the TP-AGB phase, and such models have been extensively used to reconstruct star formation histories in both resolved and unresolved populations.","the TP-AGB phase, and such models have been extensively used to reconstruct star formation histories in both resolved and unresolved populations."
+" Models of the TP-AGB phase, however, are extremely dependent on the assumptions related to mass loss and their predictive power 1s thus limited."," Models of the TP-AGB phase, however, are extremely dependent on the assumptions related to mass loss and their predictive power is thus limited."
+" There has been intriguing work suggesting that the fraction of light emitted by TP-AGB stars in LMC star clusters is very high, of order 40 (Perssonetal.1983)."," There has been intriguing work suggesting that the fraction of light emitted by TP-AGB stars in LMC star clusters is very high, of order 40 \citep{Persson83}."
+". Such a substantial flux, largely redistributed to the far-IR by dust, can be especially important for interpreting intermediate redshift galaxy properties (Conroyctal.2009)."," Such a substantial flux, largely redistributed to the far-IR by dust, can be especially important for interpreting intermediate redshift galaxy properties \citep{Conroy09}."
+. There has been a strong emphasis on updating population synthesis to include the TP-AGB phase (e.g.Bruzual&Charlot2003:Maraston2005).," There has been a strong emphasis on updating population synthesis to include the TP-AGB phase \citep[\eg][]{Bruzual03, Maraston05}."
+. Marastonetal.(2006) demonstrated that modeling of the TP-AGB phase has become a defining characteristic of different stellar population synthesis (SPS) codes while Conroyetal.(2009) identified the substantial uncertainties in TP-AGB properties as a major component of the error budget for galaxy evolution models., \citet{Maraston06} demonstrated that modeling of the TP-AGB phase has become a defining characteristic of different stellar population synthesis (SPS) codes while \citet{Conroy09} identified the substantial uncertainties in TP-AGB properties as a major component of the error budget for galaxy evolution models.
+ Inthis paper we employ a fuel consumption argument to set a firm lower bound on the fraction of light emitted during the TP-AGB phase., In this paper we employ a fuel consumption argument to set a firm lower bound on the fraction of light emitted during the TP-AGB phase.
+" We use stellar interiors models to set the core mass at the onset of the TP-AGB phase, and demonstrate that it is surprisingly insensitive to the choice of input physics."," We use stellar interiors models to set the core mass at the onset of the TP-AGB phase, and demonstrate that it is surprisingly insensitive to the choice of input physics."
+" The nuclear processed core then grows until the envelope is expelled, at which point the final white dwarf mass is set."," The nuclear processed core then grows until the envelope is expelled, at which point the final white dwarf mass is set."
+ The white dwarf initial-final mass relationship (1ΗΜΒ) can in turn be inferred from open clusters., The white dwarf initial-final mass relationship (IFMR) can in turn be inferred from open clusters.
+" We use the difference between the final and starting masses in the TP-AGB phase as a bound on the fuel consumed, and thus the emitted light."," We use the difference between the final and starting masses in the TP-AGB phase as a bound on the fuel consumed, and thus the emitted light."
+" Although it is common for investigators to compare their models to initial-final mass relationships, and fuel consumption arguments have been used to check population synthesis models, the quantitative bounds from the IFMR are typically not used in population synthesis calculations."," Although it is common for investigators to compare their models to initial-final mass relationships, and fuel consumption arguments have been used to check population synthesis models, the quantitative bounds from the IFMR are typically not used in population synthesis calculations."
+ We argue that the additional empirical information encoded there permits a narrower set of possibilities than considered by Conroyetal.(2009)., We argue that the additional empirical information encoded there permits a narrower set of possibilities than considered by \citet{Conroy09}.
+". Processed fuel (especially helium) can be ejected in winds, so the white dwarf data formally sets only a lower bound on the emitted light."," Processed fuel (especially helium) can be ejected in winds, so the white dwarf data formally sets only a lower bound on the emitted light."
+" This raises interesting links between the light emitted in this phase and chemical evolution studies, which we discuss in our conclusions."," This raises interesting links between the light emitted in this phase and chemical evolution studies, which we discuss in our conclusions."
+ The plan of our paper is straightforward., The plan of our paper is straightforward.
+" Our sample and methods are discussed in Section 2, our results are presented in Section 3, and we discuss their broader implications in Section 4."," Our sample and methods are discussed in Section 2, our results are presented in Section 3, and we discuss their broader implications in Section 4."
+ Stellar evolution theory makes a robust prediction for the core mass at the onset of the phase as a function of composition and initial mass., Stellar evolution theory makes a robust prediction for the core mass at the onset of the TP-AGB phase as a function of composition and initial mass.
+ Open clusters provide a laboratory where we can constrain the initial mass given a sample of known stellar remnants., Open clusters provide a laboratory where we can constrain the initial mass given a sample of known stellar remnants.
+ Through fuel, Through fuel
+o—(p 1)/2=0.75.,$\alpha=(p-1)/2=0.75$.
+ Adopting muin=10 MlIz. we have (Longair 19904) The Fitt Alexander (1993) sample used in the preceding section is representative. but. by no means complete.," Adopting $\nu_{\rm min}=10$ MHz, we have (Longair 1994) The Fitt Alexander (1993) sample used in the preceding section is representative, but by no means complete."
+ Using racio luminosity as a proxy for magnetic field strength allows us to use a complete (in both optical and radio) sample of Shabala (2008)., Using radio luminosity as a proxy for magnetic field strength allows us to use a complete (in both optical and radio) sample of Shabala (2008).
+ Synchrotron luminosity for late-type galaxies in the sample is plotted as a function. of stellar mass in Figure 7.., Synchrotron luminosity for late-type galaxies in the sample is plotted as a function of stellar mass in Figure \ref{fig:LradioVsMstars}.
+ Llere. late-type galaxies are defined: as galaxies with concentration indices in the Petrosian r-banel of Co> 0.375.," Here, late-type galaxies are defined as galaxies with concentration indices in the Petrosian $r$ -band of $C>0.375$ ."
+ Stellar masses were obtained from z-band magnitudes by assuming solar-tvpe stars (Shabalaetal.2008: 2004).," Stellar masses were obtained from $z$ -band magnitudes by assuming solar-type stars \cite{SAAR08,NikolicEA04}."
+. Given a disk size (ic. the parameter isk. set to 0.25 in Section 4.1)). the model predicts a single radio luminosity. or each halo. and. hence a given stellar mass.," Given a disk size (i.e. the parameter $\epsilon_{\rm disk}$, set to $0.25$ in Section \ref{sec:diskSizes}) ), the model predicts a single radio luminosity for each halo, and hence a given stellar mass."
+ Figure 7 shows that these agree well with mean observed luminosities at low stellar masses., Figure \ref{fig:LradioVsMstars} shows that these agree well with mean observed luminosities at low stellar masses.
+ However. at higher masses the racio uminosity is overpredicted.," However, at higher masses the radio luminosity is overpredicted."
+ The radio-Ioud. galaxies in the lux-limited sample of Shabala make up the bright tail of the. radio. luminosity. distribution., The radio-loud galaxies in the flux-limited sample of Shabala make up the bright tail of the radio luminosity distribution.
+ The majority of galaxies have racio uminosities below the detection limit. and thus the mean radio luminosity at a given stellar mass was calculated by assuming a Gaussian cistribution with the observed scatter of 0.3 dex.," The majority of galaxies have radio luminosities below the detection limit, and thus the mean radio luminosity at a given stellar mass was calculated by assuming a Gaussian distribution with the observed scatter of $0.3$ dex."
+ By constrast. the Fitt Alexander observations were performed. with much greater sensitivity. and thus no such corrections are necessary.," By constrast, the Fitt Alexander observations were performed with much greater sensitivity, and thus no such corrections are necessary."
+ The peak racio luminosity in Figure 7. corresponds to the D-field peak in Figure 2.. and is related to an increase in gas cooling (and hence the random. component of the magnetic field) at late times in haloes with Mia~107 (Ligures lee and É. and Figure 3cc).," The peak radio luminosity in Figure \ref{fig:LradioVsMstars} corresponds to the $B$ -field peak in Figure \ref{fig:BvsMhalo}, and is related to an increase in gas cooling (and hence the random component of the magnetic field) at late times in haloes with $M_{\rm halo} \sim 10^{13}$ (Figures \ref{fig:effPot}e e and f, and Figure \ref{fig:BcomponentsVsZ}c c)."
+ This late cooling is a result of the way ACN feedback. is implemented. in. our model., This late cooling is a result of the way AGN feedback is implemented in our model.
+ La haloes more massive that ~LOY? supernovae feedback is too weak to quench gas cooling and star formation., In haloes more massive that $\sim 10^{12}$ supernovae feedback is too weak to quench gas cooling and star formation.
+ AGN heating can in principle do this. however only the dense central regions can be heated: by weak AGNs found in haloes with mass <Lol— ," AGN heating can in principle do this, however only the dense central regions can be heated by weak AGNs found in haloes with mass $\leq 10^{13}$."
+These cores have short cooling times. and thus the net ellect of AGN heating is negligible.," These cores have short cooling times, and thus the net effect of AGN heating is negligible."
+ By contrast. in the most massive hosts powerful ACGNs can heat rarefied gas at larger radii. where the cooling times are longer. anc thus ACN heating is more cllective.," By contrast, in the most massive hosts powerful AGNs can heat rarefied gas at larger radii, where the cooling times are longer, and thus AGN heating is more effective."
+ As pointed out by Shabala Alexander (2000). an important feedback ingredient missing in our mioclel is gas ejection by powerful radio sources.," As pointed out by Shabala Alexander (2009), an important feedback ingredient missing in our model is gas ejection by powerful radio sources."
+ In. addition to. shock jeating. jet-inflated radio cocoons sweep up the ambient eas through which they expand. transporting it to larger radii (e.g. Alexander 2002: Basson Alexander 2003).," In addition to shock heating, jet-inflated radio cocoons sweep up the ambient gas through which they expand, transporting it to larger radii (e.g. Alexander 2002; Basson Alexander 2003)."
+ The onger cooling times will then prevent the hot. gas from cooling. decreasing the turbulent energy densities. magnetic ield strengths and. radio luminosities in AGN hosts.," The longer cooling times will then prevent the hot gas from cooling, decreasing the turbulent energy densities, magnetic field strengths and radio luminosities in AGN hosts."
+ A comparison of the present dav cold gas mass function orediceted by our model with the observed LE mass function (Zwaanetal.2003). indicates that the amount of cold gas in massive hosts is overpredicted by 0.51 dex., A comparison of the present day cold gas mass function predicted by our model with the observed I mass function \cite{ZwaanEA03} indicates that the amount of cold gas in massive hosts is overpredicted by $0.5-1$ dex.
+ The ellects of adding ACN gas uplifting to our model on the wedieted: magnetic fields can be estimated. by decreasing he parameter eqs by this factor., The effects of adding AGN gas uplifting to our model on the predicted magnetic fields can be estimated by decreasing the parameter $\epsilon_{\rm grav}$ by this factor.
+" Figure 7. shows that this wovicles a better fit to observations. however the rise in racio. LuminosityMN between Al,10123.107 is still too steep."," Figure \ref{fig:LradioVsMstars} shows that this provides a better fit to observations, however the rise in radio luminosity between $\Mstar\/ = 10^{12}-3 \times 10^{12}$ is still too steep."
+ This is not surprising. since this is oweciselv the range where distributed supernovae feedback cannot stop large amounts of gas [rom cooling. while AGN wating is not vet powerful enough to allect easat large radii with long cooling times.," This is not surprising, since this is precisely the range where distributed supernovae feedback cannot stop large amounts of gas from cooling, while AGN heating is not yet powerful enough to affect gasat large radii with long cooling times."
+ One would expect the gas uplifting o plav a more important. role in these intermecdiate-niass aloes., One would expect the gas uplifting to play a more important role in these intermediate-mass haloes.
+wilh Sevfert nuclei.,with Seyfert nuclei.
+ While (he division of early-type spirals into (wo categories has been uselul. we have to be cautious.," While the division of early-type spirals into two categories has been useful, we have to be cautious."
+ First. it is diffieult to measure the thus of an individual region.," First, it is difficult to measure the flux of an individual region."
+ In ILD99. we measured fluxes for the largest regions manually. using circular apertures.," In HD99, we measured fluxes for the largest regions manually, using circular apertures."
+" However. since (hen we have been using an automated program. ""IHllphot (See Thilker.DraunWallerbos(2000) [or details of the program). to determine the Iuminositv. fanctions Lor star forming regions in spiral galaxies (Ilameecl. Thilker. Devereux preparation),"," However, since then we have been using an automated program, “HIIphot” (See \citet{Thilker2000} for details of the program), to determine the luminosity functions for star forming regions in spiral galaxies (Hameed, Thilker, Devereux )."
+ All of our manual classifications are consistent with HlIIphot. with the exception of NGC 3169 and NGC 7213.," All of our manual classifications are consistent with HIIphot, with the exception of NGC 3169 and NGC 7213."
+ Both of these galaxies host regions with ))10?ergs|. and thus have been classified as Categorye 2 egalaxies in the current paper.," Both of these galaxies host regions with $> 10^{39}~\lum$, and thus have been classified as Category 2 galaxies in the current paper."
+ second. we do not vet. know of anv physical significance of our classification scheme.," Second, we do not yet know of any physical significance of our classification scheme."
+" A few studies. however. have suggested that the physical properties of regions change between LO—10?ergs!. such as the transition from normal “giant” regions to ""supereiant regions (Ixennicutt.Edgar.&Ilodge|1989).. the transition Irom sparse embecdced star clusters to dense embedded. clusters in individual regions 1993).. and the claim of the change from radiation bounded to density bounded regions (Beckmane£al.2000)."," A few studies, however, have suggested that the physical properties of regions change between $10^{38}-10^{39}~\lum$, such as the transition from normal “giant” regions to “supergiant” regions \citep{KEH1989}, the transition from sparse embedded star clusters to dense embedded clusters in individual regions \citep{OC1998}, and the claim of the change from radiation bounded to density bounded regions \citep{Beckman2000}."
+. Whatever the underlving significance of our scheme. it is of sufficient importance to report the detection of a significant fraction of early (wpe spirals will giant regions. which is in contrast to the earlier findings of and Caldwellefaf(1991).," Whatever the underlying significance of our scheme, it is of sufficient importance to report the detection of a significant fraction of early type spirals with giant regions, which is in contrast to the earlier findings of \citet{KEH1989} and \citet{Caldwell1991}."
+. The presence of massive stars can be traced via both and far-infrared emission., The presence of massive stars can be traced via both and far-infrared emission.
+ The relationship between the global and far-infrared (40-120 jan) huminosities is presented in Figure 6 lor 51 Sa-Sab galaxies., The relationship between the global and far-infrared (40-120 $\mu$ m) luminosities is presented in Figure 6 for 51 Sa-Sab galaxies.
+ Open dots represent Category 1 earlv-tvpe spirals. Category 2 galaxies are represented by solid dots. ancl (he (wo unclassified galaxies. NGC 660 and NGC 2146. have been plotted as stars.," Open dots represent Category 1 early-type spirals, Category 2 galaxies are represented by solid dots, and the two unclassified galaxies, NGC 660 and NGC 2146, have been plotted as stars."
+ The far-infrared Iuminositv. L(EIBR). has been calculated using. where S(FIR) = (2.58 Soon + Stoop) and D is the distance expressed in Mpc.," The far-infrared luminosity, L(FIR), has been calculated using, where S(FIR) = (2.58 $_{60\mu m}$ + $_{100\mu m}$ ) and D is the distance expressed in Mpc."
+" The values Suo, and S105 ave ηχος measured by the Infrared. Astronomical Satellite RAS) in units of Jv.", The values $_{60\mu m}$ and $_{100\mu m}$ are fluxes measured by the Infrared Astronomical Satellite (IRAS) in units of Jy.
+ The solid lines in Figure 6 identilv ratios of )) expected for, The solid lines in Figure 6 identify ratios of ) expected for
+is. however. a stronger discriminant against the eiaut GC formation scenario.,"is, however, a stronger discriminant against the giant GC formation scenario."
+ Here theory predicts that au L relation will be imprinted bv the plysics of massive cluster formation., Here theory predicts that an $L$ relation will be imprinted by the physics of massive cluster formation.
+ For merged (ιο a range of sizes are possible at a given huuinosity (Britisσιctal.2011)..," For merged GCs, a range of sizes are possible at a given luminosity \citep{2011A&A...529A.138B}."
+" Another observation frou Figure ὃ is a lint of a gap in the UCD size distribution at nm,~ 1520 pc (see also iniddle panel of Figure 3)). which was previously noticed for the UCDs near the center of a galaxy cluster by Blakeslee&BarberDeCwaaft(2008)."," Another observation from Figure \ref{fig:uber}
+ is a hint of a gap in the UCD size distribution at $r_{\rm h}\sim$ 15–20 pc (see also middle panel of Figure \ref{fig:hist}) ), which was previously noticed for the UCDs near the center of a galaxy cluster by \citet{2008AJ....136.2295B}."
+. Such a feature nueht iaply differcut modes of UCD formation. aud will ierit further exaiuination with larecr samples.," Such a feature might imply different modes of UCD formation, and will merit further examination with larger samples."
+ There also appears to be amagnitude gap. between UCDs aud ECs from My~ 8 ta ~ 9 (also called, There also appears to be a gap between UCDs and ECs from $M_V \sim$ $-8$ to $\sim$ $-9$ (also called
+equations: Note that Equation (13)) (signifying that óf is advected) is equivalent to requiring that the pressure perturbation vanishes (compare Equations (13)) and (A7)).,equations: Note that Equation \ref{eq:adv_df2}) ) (signifying that $\delta f$ is advected) is equivalent to requiring that the pressure perturbation vanishes (compare Equations \ref{eq:adv_df2}) ) and \ref{eq:dp}) )).
+" Thus, when the advected wave is subtracted from the perturbations below the shock, the pressure perturbation is unchanged."," Thus, when the advected wave is subtracted from the perturbations below the shock, the pressure perturbation is unchanged."
+" This system of equations can be solved to obtain the following expression of the advected perturbations as a function of the amplitude of the entropy and vorticity perturbations (described by the variables 0S and 6q): In order to compute purely acoustic modes, the advected wave is then subtracted from the perturbations at the shock to define new boundary values 6facsh and ORacsh? where ὃf. and dhsp are determined bythe boundary conditions at the shock (Equations (A16)) and (A16)))."," This system of equations can be solved to obtain the following expression of the advected perturbations as a function of the amplitude of the entropy and vorticity perturbations (described by the variables $\delta S$ and $\delta q$ ): In order to compute purely acoustic modes, the advected wave is then subtracted from the perturbations at the shock to define new boundary values $\delta f_{\rm acsh}$ and $\delta h_{\rm acsh}$: where $\delta f_\sh$ and $ \delta h_\sh$ are determined bythe boundary conditions at the shock (Equations \ref{eq:dfsh}) ) and \ref{eq:dfsh}) ))."
+" Ofaav and óhaa, are computed using Equations (15)) and (16)) and the values of 6S and dq at the shock (Equations (A)) and (A)))."," $\delta f_{\rm adv}$ and $\delta h_{\rm adv}$ are computed using Equations\ref{eq:fadv}) ) and \ref{eq:hadv}) ), and the values of $\delta S$ and $\delta q$ at the shock (Equations \ref{eq:dSsh}) ) and \ref{eq:dqsh}) ))."
+" Using these modified boundary conditions, the acoustic modes can then be computed numerically as described for example in ?.."," Using these modified boundary conditions, the acoustic modes can then be computed numerically as described for example in \citet{foglizzo07}."
+ Figure 6 compares the growth rate of the most unstable acoustic modes (left panel) and SASI modes (right panel)., Figure \ref{fig:acoustic_growth_rate} compares the growth rate of the most unstable acoustic modes (left panel) and SASI modes (right panel).
+" The difference is striking: the purely acoustic modes are all stable for all shock radii in the explored range ra,=2—100r., while SASI modes have significant growth rates."," The difference is striking: the purely acoustic modes are all stable for all shock radii in the explored range $r_\sh=2-100r_*$, while SASI modes have significant growth rates."
+ This shows that the purely acoustic cycle cannot account for SASI and that the advected wave is essential to the instability., This shows that the purely acoustic cycle cannot account for SASI and that the advected wave is essential to the instability.
+" The most ""unstable"" acoustic mode (actually the least damped) for each | is the lowest frequency one."," The most ""unstable"" acoustic mode (actually the least damped) for each $l$ is the lowest frequency one."
+ The discussion of the purely acoustic frequency spectrum is delayed to Section 5.2.2.., The discussion of the purely acoustic frequency spectrum is delayed to Section \ref{sec:SASI_spectrum}.
+" In order to quantify the effect of non-adiabaticity at the shock on the acoustic modes, we have used boundary conditions which either take into account or neglect the cooling processes in the computation of 6S and óq."," In order to quantify the effect of non-adiabaticity at the shock on the acoustic modes, we have used boundary conditions which either take into account or neglect the cooling processes in the computation of $\delta S$ and $\delta q$."
+" The associated change in the frequency and growth rate of the modes is barely noticeable (less than 296 for any shock radius), confirming the validity of our quasi-adiabatic approximation."," The associated change in the frequency and growth rate of the modes is barely noticeable (less than $2\%$ for any shock radius), confirming the validity of our quasi-adiabatic approximation."
+ The error due to the high frequency approximation can in turn be quantified in the following way., The error due to the high frequency approximation can in turn be quantified in the following way.
+" Noting that there is some arbitrariness in the choice of the variables that should be advected, one can require another variable to be advected and compare the result with the previous calculation."," Noting that there is some arbitrariness in the choice of the variables that should be advected, one can require another variable to be advected and compare the result with the previous calculation."
+" For example requiring the perturbation of radial velocity δυρ/υν to be advected instead of dh, we get a different definition of the advected perturbations: where e; and €2 are two parameters vanishing in the limit wrsp/Usn>> 1, and defined by:"," For example requiring the perturbation of radial velocity $\delta v_r/v_r$ to be advected instead of $\delta h$ , we get a different definition of the advected perturbations: where $\epsilon_1$ and $\epsilon_2$ are two parameters vanishing in the limit $\omega r_\sh/v_\sh \gg 1 $ , and defined by:"
+moments of inertia outlined in this model.,moments of inertia outlined in this model.
+ However. the association of the two reeiues with two different dynaimos is worth exploring further.," However, the association of the two regimes with two different dynamos is worth exploring further."
+ ere we sugsest that coronal saturation docs not require a separate physical iechanisia. but is a manifestation of the differcut dynaimos that are prescut ii stars on the convective and interface sequences.," Here we suggest that coronal saturation does not require a separate physical mechanism, but is a manifestation of the different dynamos that are present in stars on the convective and interface sequences."
+ Iu this picture. slowly rotating stars have ai cyclic. interface-type dynamo with Skumanich-stvle spin-down. while faster rotators have a turbulent dynamo.," In this picture, slowly rotating stars have a cyclic, interface-type dynamo with Skumanich-style spin-down, while faster rotators have a turbulent dynamo."
+ The efüciencies of the two different dynauuos could have sienificautly different dependencics. in terms of PAstellar paralucters. and therefore have differcut forms in the Ry Ro diagram.," The efficiencies of the two different dynamos could have significantly different dependencies, in terms of stellar parameters, and therefore have different forms in the $R_X$ $Ro$ diagram."
+ N-rav saturation would therefore not be an actual saturation m any sense. but a completely different magnetic dvyianaio configuration with completely different depeudenucies.," X-ray saturation would therefore not be an actual saturation in any sense, but a completely different magnetic dynamo configuration with completely different dependencies."
+ Tf this scenario is correct it partly auswers the question raised by 7) a8 to why the bolometric Dpmuiuinositv is the only parameter necessary to determine the N-rav huuinosity in the saturated regime. despite the Rossby uunber being cousidered the fundamental parameter of the efficiency of the stellar dynamo.," If this scenario is correct it partly answers the question raised by \citet{pizz03} as to why the bolometric luminosity is the only parameter necessary to determine the X-ray luminosity in the saturated regime, despite the Rossby number being considered the fundamental parameter of the efficiency of the stellar dynamo."
+ The answer is that there are two different dvnamo coufieurations. a convective dynamo paraneterized by the bolometric huninosity. an interface dynamo parameterized by the Rossby uuuber.," The answer is that there are two different dynamo configurations, a convective dynamo parameterized by the bolometric luminosity, an interface dynamo parameterized by the Rossby number."
+" This does however raise the question of why the transition between the two regimes is apparcutly so smooth in the Ly/Ly., versus Ro diagram.", This does however raise the question of why the transition between the two regimes is apparently so smooth in the $L_X / L_{bol}$ versus $Ro$ diagram.
+ In other words. why is the huninositv ratio for uusaturated stars close to the saturation threshold equal to ~1039," In other words, why is the luminosity ratio for unsaturated stars close to the saturation threshold equal to $\sim 10^{-3}$?"
+ The answer to this question will require a mach better understanding ofthe two cyvuamo configurations prescut in the saturated aud unsaturated reeiues., The answer to this question will require a much better understanding of the two dynamo configurations present in the saturated and unsaturated regimes.
+" At very high rotational velocities the fractional X-rav Duninositv has been observed to decrease below the saturation level (7)... an effect dubbed ""supersaturatiou."," At very high rotational velocities the fractional X-ray luminosity has been observed to decrease below the saturation level \citep{rand96}, an effect dubbed “supersaturation”."
+ This effect has ouly been persuasively observed iu the coronal cussion frou vouug (<100 AIvrs) C aud Is dwarts and is primarily based on projected rotational velocities. not rotation periods. in the vouug clusters IC 2391. IC 2602. aud a Persei (e.g.22?) from which 7). found a decline in X-rav saturation for stars with Μος0.01.," This effect has only been persuasively observed in the coronal emission from young $< 100$ Myrs) G and K dwarfs and is primarily based on projected rotational velocities, not rotation periods, in the young clusters IC 2391, IC 2602, and $\alpha$ Persei \citep[e.g.][]{rand96,pros96,stau97,step01} from which \citet{stau97} found a decline in X-ray saturation for stars with $Ro < 0.01$."
+ Despite this. evidence for supersaturation in lower-nuass stars is scarce: both ο) and ?7) could fud uo evidence in samples of ultra-fast rotating M dwarts.," Despite this, evidence for supersaturation in lower-mass stars is scarce; both \citet{jame00} and \citet{jeff11} could find no evidence in samples of ultra-fast rotating M dwarfs."
+ Like coronal saturation. the phenomenon of supeor-saturation and its plivsical origin is still heavily debated.," Like coronal saturation, the phenomenon of super-saturation and its physical origin is still heavily debated."
+ Figure 5 shows ενω for all stars with saturated X-ray clission as a function of various quantities suggested as parameters influcucing the supersaturation effect., Figure \ref{supersaturation} shows $L_X / L_{bol}$ for all stars with saturated X-ray emission as a function of various quantities suggested as parameters influencing the supersaturation effect.
+ The first two panels of Figure 5. show the effect of rotation period aud Rossby uunuber ou the X-ray Iuninosity ratio. ucither of which show evideuce for supersaturation. or a notable decline iu the huuimositv ratio for either the fastest rotators or the sources with the smallest Rosshy numbers.," The first two panels of Figure \ref{supersaturation} show the effect of rotation period and Rossby number on the X-ray luminosity ratio, neither of which show evidence for supersaturation, or a notable decline in the luminosity ratio for either the fastest rotators or the sources with the smallest Rossby numbers."
+ 7) developed the theory of ?) that the X-ray euüttiug volume of rapid rotators is reduced xia centrifugal stripping., \citet{jame00} developed the theory of \citet{jard99} that the X-ray emitting volume of rapid rotators is reduced via centrifugal stripping.
+ This theory was initially put forward to explain XN-rav saturation. with the increase in coronal density balancing the decrease in volume and leading to saturation. but ?) aud 7) used it to explain super-saturation for stars rotating fast chough that the Keplerian corotation radius becomes very close to the stellar surface.," This theory was initially put forward to explain X-ray saturation, with the increase in coronal density balancing the decrease in volume and leading to saturation, but \citet{jame00} and \citet{jard04} used it to explain super-saturation for stars rotating fast enough that the Keplerian corotation radius becomes very close to the stellar surface."
+ As a function of the stellar radius. the corotation radius is given by ?7) which. using stellu: properties taken from the models of 7). allows the fractional N-raw lunünositv to be calculated as a function of the relative corotation radius. as shown in Figure 5..," As a function of the stellar radius, the corotation radius is given by \citet{jame00}
+ which, using stellar properties taken from the models of \citet{sies00}, allows the fractional X-ray luminosity to be calculated as a function of the relative corotation radius, as shown in Figure \ref{supersaturation}. ."
+ In this sample there appears to be some evidence for a decrease in the typical fractional N-rav luminosity of stars with the lowest corotation radi., In this sample there appears to be some evidence for a decrease in the typical fractional X-ray luminosity of stars with the lowest corotation radii.
+ Based ou the mean Lyζω a decrease is seen for Πορ/ReS dn agreement with that fouud by 7) and ?7)..," Based on the mean $L_X / L_{bol}$ a decrease is seen for $R_{Kepler} / R_\star \lesssim 3$, in agreement with that found by \citet{jame00} and \citet{jeff11}."
+ This effect is still evident if oue cousidors ouly I& stars. aud to a lesser extent C stars. but not M stars. as found by previous authors.," This effect is still evident if one considers only K stars, and to a lesser extent G stars, but not M stars, as found by previous authors."
+ As parzuneterized in Figure 5 the bin containing the fastest rotators hints at a nioderate decline to Ly/£Ly.=δια40.26. deviating lo below the mean saturation value.," As parameterized in Figure \ref{supersaturation} the bin containing the fastest rotators hints at a moderate decline to $L_X / L_{bol} = -3.41 \pm 0.26$, deviating $\sigma$ below the mean saturation value."
+ Ao two-sample Iohuogorov-Smürnov (ISS) test comparing the N-rav huninosity ratio of stars above and below the proposed corotation supersaturation threshold of zy:Rs~3 gives a probability of <10? that the two samples are drawn from the same population., A two-sample Kolmogorov-Smirnov (KS) test comparing the X-ray luminosity ratio of stars above and below the proposed corotation supersaturation threshold of $R_{Kepler} / R_\star \sim 3$ gives a probability of $< 10^{-3}$ that the two samples are drawn from the same population.
+ 7) offered an alternative explanation that the reduction ii N-ray huuinositv is caused by a decrease in the filling factor of coronal active regions on the stellar surface., \citet{step01} offered an alternative explanation that the reduction in X-ray luminosity is caused by a decrease in the filling factor of coronal active regions on the stellar surface.
+ They propose that a poleward migration of active regions would be induced by nouuniform heating of the convective envelope in accordance with the vou Zeipel theorem., They propose that a poleward migration of active regions would be induced by nonuniform heating of the convective envelope in accordance with the von Zeipel theorem.
+ ?7)— argued that the effect of rotation ou the internal radiative transfer within a star causes a lieher fraction of dux at the poles than at the equator. an effect known as eravitv-«darkeniug.," \citet{vonz24} argued that the effect of rotation on the internal radiative transfer within a star causes a higher fraction of flux at the poles than at the equator, an effect known as gravity-darkening."
+ For sufiicicutly fast rotating stars this flux iuibalance can also induce strong convective updrafts iu the outer convective envelope (?).. capable of sweeping magnetic flux tubes iu the interior to the poles before they rise to the surface.," For sufficiently fast rotating stars this flux imbalance can also induce strong convective updrafts in the outer convective envelope \citep{step98}, capable of sweeping magnetic flux tubes in the interior to the poles before they rise to the surface."
+ According to the vou Zeipel theorem the ciuerecut surface flux for a rotating solid body is proportional to the local effective gravity. ο. given by where wis the aneular rotation rate. & is the radius. and 0 is the colatitucle (differencesiuthetheoremforadiffeveutiallvrotatingstarareminor.," According to the von Zeipel theorem the emergent surface flux for a rotating solid body is proportional to the local effective gravity, $g_{eff}$, given by where $\omega$ is the angular rotation rate, $R$ is the radius, and $\theta$ is the colatitude \citep[differences in the theorem for a differentially rotating star are minor,][]{maed99}."
+" 7). 7) have argued that the resulting chauge iu XN-aav huninositv is proportional to the relative streneth of this ""polar updraft effect at the boundary between the radiative interior aud thecouvective envelope. which we quautifv as where R, is the radius of the radiativecouvective boundary."," \citet{step01} have argued that the resulting change in X-ray luminosity is proportional to the relative strength of this “polar updraft” effect at the boundary between the radiative interior and theconvective envelope, which we quantify as where $R_c$ is the radius of the radiative–convective boundary."
+ The efficicney of this mechanisin therefore scales as the rotation rate to the second power and the radius of the radiativeconvective boundary to the third, The efficiency of this mechanism therefore scales as the rotation rate to the second power and the radius of the radiative–convective boundary to the third
+Figure 14. shows the annotated (33° roll angle) for the star HD120780.,Figure \ref{annotate_hd120780} shows the annotated $^{\circ}$ roll angle) for the star HD120780.
+ The angle is projected along the solid lines radiating from the central pixel., The angle is projected along the solid lines radiating from the central pixel.
+ The brightest object in. this image is found between the projected solid lines. however since no counterpart was found this is à residual super speckle and shows that such artifacts still remain. after SDI reduction.," The brightest object in this image is found between the projected solid lines, however since no counterpart was found this is a residual super speckle and shows that such artifacts still remain after SDI reduction."
+ The potential candidate around this star is found to reside at a separation of 0.30+0.01 from the central pixel with a API contrast of ~8.8 magnitudes. and is again well below the strong methane boundary.," The potential candidate around this star is found to reside at a separation of $\pm$ $''$ from the central pixel with a $\Delta$ F1 contrast of $\sim$ 8.8 magnitudes, and is again well below the strong methane boundary."
+ Significantly though. the separation 1s in agreement with the object found around the star HD25874.," Significantly though, the separation is in agreement with the object found around the star HD25874."
+ Also the position angle of the object is 240° which again is in agreement with the find around HD25874., Also the position angle of the object is $^{\circ}$ which again is in agreement with the find around HD25874.
+ The fact that both the separation and position angle agree for the possible candidate objects around both HD25874 and HD120780 strongly indicate that these detections are actually artifacts inherent in the reduction procedure rather than an actual detections of the companion objects to both these stars. possibly spider arm residuals.," The fact that both the separation and position angle agree for the possible candidate objects around both HD25874 and HD120780 strongly indicate that these detections are actually artifacts inherent in the reduction procedure rather than an actual detections of the companion objects to both these stars, possibly spider arm residuals."
+ The contrast curves created for this system (Fig., The contrast curves created for this system (Fig.
+ 15 upper) are extrapolated from the unsaturated images of the other systems., \ref{hd120780_contrast} upper) are extrapolated from the unsaturated images of the other systems.
+ The detection here is marked by the filled circle with associated uncertainties to highlight the contrasts achieved at such low angular separations., The detection here is marked by the filled circle with associated uncertainties to highlight the contrasts achieved at such low angular separations.
+ The detectability probability in this region is found to be50-70%.. again showing that more data would be useful to better constrain the companion parameters. even at such low separations.," The detectability probability in this region is found to be, again showing that more data would be useful to better constrain the companion parameters, even at such low separations."
+ The detectability region Is better seen in Fig., The detectability region is better seen in Fig.
+ 15. (lower) where again we have over confidence reaching out beyond 0.1” from the star and with fairly high brown dwarf masses. across the stellar mass regime.," \ref{hd120780_contrast} (lower) where again we have over confidence reaching out beyond $''$ from the star and with fairly high brown dwarf masses, across the stellar mass regime."
+ The gray scale in this region shows no real dark unconstrained regions. except for widely separated and low mass companions.," The gray scale in this region shows no real dark unconstrained regions, except for widely separated and low mass companions."
+ Below the SDI limits we see that we can rule out lower mass companions with confidence only out to separations of «0.14 (~2.38 AU) at best., Below the SDI limits we see that we can rule out lower mass companions with confidence only out to separations of $\sim$ $''$ $\sim$ 2.38 AU) at best.
+" Yet above the lc confidence level (7040) we can rule out a range of lower mass companions reaching às far out as 0.25"" (4.25 AU).", Yet above the $\sigma$ confidence level ) we can rule out a range of lower mass companions reaching as far out as $''$ (4.25 AU).
+ Finally no other longer period methane rich companions were detected out to a separation of ~34AU., Finally no other longer period methane rich companions were detected out to a separation of $\sim$ 34AU.
+ We have performed a targeted direct imaging program to detect cool companions orbiting within 2 of their parent stars., We have performed a targeted direct imaging program to detect cool companions orbiting within $''$ of their parent stars.
+ The stars were drawn from the AAPS and Keck planet search projects and consist of objects that exhibit large radial-velocity variation over several years., The stars were drawn from the AAPS and Keck planet search projects and consist of objects that exhibit large radial-velocity variation over several years.
+ These radial-velocities indicate the presence of à massive companion on a long period orbit., These radial-velocities indicate the presence of a massive companion on a long period orbit.
+ Five stars were examined with the NACO-SDI system on the VLT in Paranal. Chile.," Five stars were examined with the NACO-SDI system on the VLT in Paranal, Chile."
+ From the five. two possible detections were found around the stars HD25874 and HD120780.," From the five, two possible detections were found around the stars HD25874 and HD120780."
+ However. further analysis of these detections indicate they are probable residual artifacts since they are found to be located at the same distance from the central pixel and the same position angle in each of the images for each star.," However, further analysis of these detections indicate they are probable residual artifacts since they are found to be located at the same distance from the central pixel and the same position angle in each of the images for each star."
+ In addition we also present detectabilities for each system by analysing the radial-velocity information we have acquired., In addition we also present detectabilities for each system by analysing the radial-velocity information we have acquired.
+ Each of these detections lie within sensitivity boundaries for these stars of between70%.. meaning they could not be ruled out with any high degree of certainty from the radial-velocity data.," Each of these detections lie within sensitivity boundaries for these stars of between, meaning they could not be ruled out with any high degree of certainty from the radial-velocity data."
+" We have summarised the results of this work in Table 7.. which shows both the broad-band AO and narrow-band SDI reduced contrasts and mass limits at separations of 0.5"" and 1.0""."," We have summarised the results of this work in Table \ref{tab:results}, , which shows both the broad-band AO and narrow-band SDI reduced contrasts and mass limits at separations of $''$ and $''$."
+ Also the sensitivity limits from the radial-velocity data have been summarised at confidence limits of 70. 90 and for each system. highlighting the minimum mass. possibly detectable at each confidence level.," Also the sensitivity limits from the radial-velocity data have been summarised at confidence limits of 70, 90 and for each system, highlighting the minimum mass possibly detectable at each confidence level."
+" The table shows the contrasts in magnitudes and mass limits in Jupiter-masses. at 0.5"" and 1.0” angular separations for AO and SDI images and shows the AO performs almost as well as the SDI at separations of 21.0"".", The table shows the contrasts in magnitudes and mass limits in Jupiter-masses at $''$ and $''$ angular separations for AO and SDI images and shows the AO performs almost as well as the SDI at separations of $\ge$ $''$.
+ Also since the stars are fairly old the AO reduction tends to reach only a few Jupiter-masses above the SDI reduction for most of the sample. however when the star is fairly young (HD145825) the SDI reaches far deeper than AO i.e. 13M; in this case.," Also since the stars are fairly old the AO reduction tends to reach only a few Jupiter-masses above the SDI reduction for most of the sample, however when the star is fairly young (HD145825) the SDI reaches far deeper than AO i.e. $_{\rm{J}}$ in this case."
+ Also from this table we see that the radial-velocity confidence limits are all generally the same. apart from HD32778 due its limited number of data points and temporal coverage.," Also from this table we see that the radial-velocity confidence limits are all generally the same, apart from HD32778 due its limited number of data points and temporal coverage."
+ As the stars chosen are fairly old. à consequence of the radial-velocity selection method. the companions must. be sufficiently massive to lie within the detection threshold of the instrument.," As the stars chosen are fairly old, a consequence of the radial-velocity selection method, the companions must be sufficiently massive to lie within the detection threshold of the instrument."
+" Also since three of the five stars were found to have ""liner"" trends. and we expect these to be sufficiently massive. then we suspect either the uncertainty in age means we are underestimating our mass thresholds. the models are over estimating the magnitudes of the companions. the companions are aligned such that theyare found behind. or very close to the central star (x0.1) "," Also since three of the five stars were found to have 'liner' trends, and we expect these to be sufficiently massive, then we suspect either the uncertainty in age means we are underestimating our mass thresholds, the models are over estimating the magnitudes of the companions, the companions are aligned such that theyare found behind, or very close to the central star $\le$ "
+"where the factor C now stands for We take this to hold for the output of real RILD simulations as well. so that we have an approximate parametrisation for the particle distribution in any. grid cell in terms of 27, and 55; alone.","where the factor $C$ now stands for We take this to hold for the output of real RHD simulations as well, so that we have an approximate parametrisation for the particle distribution in any grid cell in terms of $\gamma'_m$ and $\gamma'_M$ alone."
+ A more complete treatment of the particle distribution (e.g. Peer&Waxman (2005))) would cllectively introduce an additional dimension to the RIID simulation and slow clown the caleulations accorcinely., A more complete treatment of the particle distribution (e.g. \cite{Peer2005}) ) would effectively introduce an additional dimension to the RHD simulation and slow down the calculations accordingly.
+ Via reasoning completely analogous to the non-cooling case (where we use eq. (AL)), Via reasoning completely analogous to the non-cooling case (where we use eq. \ref{nume}) )
+ with « instead of 57 to obtain voy) we arrive at an auxiliary function Q given by ocurring in Since this is a function of two variables instead of one. its limiting behaviour is more complicated.," with $\gamma'_M$ instead of $\gamma'_e$ to obtain $\nu'_{cr,M}$ ) we arrive at an auxiliary function $\mathcal{Q}$ given by ocurring in Since this is a function of two variables instead of one, its limiting behaviour is more complicated."
+ Woy «1. Qu.o) approximately reduces to which can be obtained by approximating yay by zero in the integration limits and integrand of equation (B5)).," If $y_M \ll 1$ , $\mathcal{Q}(y_M, y_m )$ approximately reduces to which can be obtained by approximating $y_M$ by zero in the integration limits and integrand of equation \ref{Q_cooling_equation}) )."
+" M l. the approximate result is which follows from approximating the integral from (B5)) by its value at jJ, times the integration domain."," If $y_m / y_M \to 1$ , the approximate result is which follows from approximating the integral from \ref{Q_cooling_equation}) ) by its value at $y_m$ times the integration domain."
+" Ly, «1 as well. we can use the first term of the lower limit series expansion for P in the integral and solve it to refine the approximate result to On the other hand. ΕνD>1. we can approximate the result in terms of Qj.gu) for a smaller value yi, (ec. the last tabulatecl value): recluces to lon j values of 55, and yi"," If $y_m \ll 1$ as well, we can use the first term of the lower limit series expansion for $\mathcal{P}$ in the integral and solve it to refine the approximate result to On the other hand, if $y_m / y_M \gg 1$, we can approximate the result in terms of $\mathcal{Q}( y_M, y'_m )$ for a smaller value $y'_m$ (i.e. the last tabulated value): which further reduces to for sufficiently high values of $y'_m$ and $y_m$."
+" Finally. for yay 3» 1 wefindfrom fitting to tabulated values that Q(gaj. 75"") is bestdescribed by Vm T UA In practice the code uses a two-dimensional table with numerically calculated: values in addition to the analytical expressions above."," Finally, for $y_M \gg 1$ we find from fitting to tabulated values that $\mathcal{Q}( y_M, y_m )$ is best described by In practice the code uses a two-dimensional table with numerically calculated values in addition to the analytical expressions above."
+" Phe contribution from the region where yayX1 is elfectively zero due to the exponential term e.M,", The contribution from the region where $y_M \gg 1$ is effectively zero due to the exponential term $e^{-y_M}$.
+ We summarize the equations for the scaling coefficients in table (C1))., We summarize the equations for the scaling coefficients in table \ref{coefficients_table}) ).
+ Aside from some explicit dependencies on p and & these equations also contain truly numerical constants with values that have been determined by fitting to output of our code., Aside from some explicit dependencies on $p$ and $k$ these equations also contain truly numerical constants with values that have been determined by fitting to output of our code.
+" For example C'p(p.k) contains the constants C'po. Cp, and Cop, (with Ch, denoting Cp, to the power & ete)."," For example $C_D(p,k)$ contains the constants $C_{D0}$, $C_{Dk}$ and $C_{Dkk}$ (with $C_{Dk}^k$ denoting $C_{Dk}$ to the power $k$ etc.)."
+ Fheir numerical values are listed in table (€2))., Their numerical values are listed in table \ref{coefficient_values_table}) ).
+ Instead of incorporating these numerical constants in the total flux formula as we have done here we could also have used a fitting polvnomial. but this approachmore closely rellects the & and p dependencies.," Instead of incorporating these numerical constants in the total flux formula as we have done here we could also have used a fitting polynomial, but this approachmore closely reflects the $k$ and $p$ dependencies."
+ The first term (p1) in these equations is also the first term in eq. (8))., The first term $(p-1)$ in these equations is also the first term in eq. \ref{F_scaling_equation}) ).
+ From the contribution of Αι we obtain a contribution1/(3.—&) via, From the contribution of $N_{tot}$ we obtain a contribution $1/(3-k)$ via
+gives greater weight to the circular polarisation curves in he fitting process.,gives greater weight to the circular polarisation curves in the fitting process.
+ We show in Fig., We show in Fig.
+ 4 the position of the accretion regions »ecdieted by our model superimposed on globes representing a complete rotation of the white ciwarf for the same moclel its as shown in Fig. 3.., \ref{globes} the position of the accretion regions predicted by our model superimposed on globes representing a complete rotation of the white dwarf for the same model fits as shown in Fig. \ref{polfit}.
+ Xn extended: evelotron accretion region. which is offset by a [arge distance from the visible magnetic pole. is first seen at ó «0.3 (and may account [or he increase in the linear. polarisation seen at this phase) and disappears at ὁ ~0.9 the primary acceretion. region.," An extended cyclotron accretion region, which is offset by a large distance from the visible magnetic pole, is first seen at $\phi\sim$ 0.3 (and may account for the increase in the linear polarisation seen at this phase) and disappears at $\phi\sim$ 0.9 – the primary accretion region."
+ The ‘tail’ of this region is optically thin evclotron radiation., The `tail' of this region is optically thin cyclotron radiation.
+ Fig., Fig.
+ 4 shows that the Leading edge of the primary region (optically thick evelotron radiation) quickly. appears over the limb of the white dwarf ancl accounts for the rapic increase in circularly polarised flux at ὁ ~0.3 (Fig. 3))., \ref{globes} shows that the leading edge of the primary region (optically thick cyclotron radiation) quickly appears over the limb of the white dwarf and accounts for the rapid increase in circularly polarised flux at $\phi\sim$ 0.3 (Fig. \ref{polfit}) ).
+ On the other hand the ‘tail’ of the accretion region takes a longer time to disappear over the limb of the white dwarf resulting in a less rapid «ecrease in the cireularly. polarisec Hux., On the other hand the `tail' of the accretion region takes a longer time to disappear over the limb of the white dwarf resulting in a less rapid decrease in the circularly polarised flux.
+ A much. smaller. accretion region is precieted very close to the spin pole an is alwavs visible — the secondary accretion region., A much smaller accretion region is predicted very close to the spin pole and is always visible – the secondary accretion region.
+ For higver values of 2 (€ 507). we finc other moclels which fit the data equally well.," For higher values of $\beta$ $>50^{\circ}$ ), we find other models which fit the data equally well."
+ In these cases the primary accretion region is less extended in. magnetic longitude and the secondary accretion is brighter ancl for 3«507., In these cases the primary accretion region is less extended in magnetic longitude and the secondary accretion is brighter and for $\beta <50^{\circ}$.
+ We consider the 3«50° set of moclels to be the most. likely solutions since in these cases the secondary region is much less significant than the primary and hence the most sipple scenario (as is the rationale of our optimisation technique)., We consider the $\beta <50^{\circ}$ set of models to be the most likely solutions since in these cases the secondary region is much less significant than the primary and hence the most simple scenario (as is the rationale of our optimisation technique).
+ While we are confident that all our best fitting mocdols predict the same general location. extent ancl phasing of the primary accretion region. we cannot be certain of the finer details of its structure as the resolution of the image is limited due to the finite sampling of our data ancl the fact that the optimisation technique tends to smooth out fine structure.," While we are confident that all our best fitting models predict the same general location, extent and phasing of the primary accretion region, we cannot be certain of the finer details of its structure as the resolution of the image is limited due to the finite sampling of our data and the fact that the optimisation technique tends to smooth out fine structure."
+ We now go on to compare our model results with the scenario of Alisadi et al (1996) who suggest that the deep X-ray dip is due to the accretion stream far from the white ανα obscuring the emission [rom the bright accretion reeion., We now go on to compare our model results with the scenario of Misaki et al (1996) who suggest that the deep X-ray dip is due to the accretion stream far from the white dwarf obscuring the emission from the bright accretion region.
+ The optical ancl X-ray data reported by Thomas Iteinsch (1996) and. Misaki ct al (1996) respectively are consistent with CP Tuc being a polar although their data could not rule out an Intermediate Polar interpretation (LPs - the non- magnetic CVs)., The optical and X-ray data reported by Thomas Reinsch (1996) and Misaki et al (1996) respectively are consistent with CP Tuc being a polar although their data could not rule out an Intermediate Polar interpretation (IPs - the non-synchronous magnetic CVs).
+ One of the main characteristics, One of the main characteristics
+The partial differcutial equations describing the disk can now be recast as ordinary differential equations aud algebraic relations in the normalized quautities.,The partial differential equations describing the disk can now be recast as ordinary differential equations and algebraic relations in the normalized quantities.
+ The equation of mass continuity becomes and the couservation of momentum (Equation (5))) becomes in the » direction. iu the @ direction. aud in the z direction. where and The quantity e9/v% is denoted by e in L9G.," The equation of mass continuity becomes and the conservation of momentum (Equation \ref{eq:moment}) )) becomes in the $r$ direction, in the $\phi$ direction, and in the $z$ direction, where and The quantity $v_{\rm A,0}^{2}/v_{\rm K}^{2}$ is denoted by $\sigma$ in L96."
+ By definition YW=δι.," By definition, $\Psi = \Psi_{*}\varphi$."
+ The relation (oe). therefore gives ο(rior)=rapeVueVQa .," The relation $B_{r}r =
+-\left(\frac{\partial \Psi}{\partial z}\right)_{r}$ therefore gives $\left(B_{z*} \varphi^{1-2\zeta} b_{r}\right)
+\left(r_{*}\varphi^{\zeta}x\right) =
+\frac{\Psi_{*}\varphi}{r_{*}\varphi^{\zeta}}\left(\frac{x^{\prime}}{\zeta
+x^{2}}\right)$ ."
+" From the definition of the flux function. E=B. so we have At a poiut (rn,z)=rags(α(8).sa(s)) on a field line. the uuit tangent vector along the feld line is This relation. combined with Equation (31)). gives Equation (32)). which expresses b, iu terum of b,. x. and s. automatically satisfies the solenoidal couclition (Equation (11)))."," From the definition of the flux function, $\frac{\Psi_{*}}{r_{*}^{2}} = B_{z*}$, so we have At a point $\left(r,z\right) = r_{*}\varphi^{\zeta}\left(x(s),sx(s)\right)$ on a field line, the unit tangent vector along the field line is This relation, combined with Equation \ref{eq:xprime}) ), gives Equation \ref{eq:brbz}) ), which expresses $b_{z}$ in terms of $b_{r}$, $x$, and $s$, automatically satisfies the solenoidal condition (Equation \ref{eq:solen}) ))."
+ Auperre’s law (Equation (9))) becomes in the r direction. in the @ direction. and in the z direction.," Ampèrre's law (Equation \ref{eq:amp}) )) becomes in the $r$ direction, in the $\phi$ direction, and in the $z$ direction."
+" Usine the algebraic relation for bz. the $ component simplifies further to The poloidal comrponeuts of Amperre’s law (Equations(33)) and (35)}) together imply Oluu’s law (Equation (1))) becomes where the comoving electric field is given by The three components of Equation (38)) are where b=v+b,b2 . and the magnetic ficld- aud -perpeudicular componcuts of the comoving electric field satisty aud"," Using the algebraic relation for $b_{z}$, the $\phi$ component simplifies further to The poloidal components of Ampèrre's law (Equations\ref{eq:ampr2}) ) and \ref{eq:ampz2}) )) together imply Ohm's law (Equation \ref{eq:ohm}) )) becomes where the comoving electric field is given by The three components of Equation \ref{eq:ohm2}) ) are where $b = \sqrt{b_{r}^{2}+b_{\phi}^{2}+b_{z}^{2}}\,$ , and the magnetic field-parallel and -perpendicular components of the comoving electric field satisfy and"
+a MCMC algorithm (e.g. 222».,"a MCMC algorithm (e.g. \citealt{Tegmark2004,Cameron2007,Gibson2008}) )."
+ We begin by calculating the 7 statistic of a set of proposed parameters. where /; is the flux observed at time 7. 7; is the model flux and 7; is the uncertainty of each /; as described in Section 3.2..," We begin by calculating the $\chi^2$ statistic of a set of proposed parameters, where $f_j$ is the flux observed at time $j$, $ m_j$ is the model flux and $\sigma_j$ is the uncertainty of each $f_j$ as described in Section \ref{errors}."
+" At each step in the MCMC chain. each proposed parameter is perturbed by a random amount which we call a ""jump function”."," At each step in the MCMC chain, each proposed parameter is perturbed by a random amount which we call a “jump function”."
+ Each jump function is proportional to the uncertainty of each parameter multiplied by a random Gaussian number with mean zero and unit standard deviation., Each jump function is proportional to the uncertainty of each parameter multiplied by a random Gaussian number with mean zero and unit standard deviation.
+ The new parameter set is acceptedwith probability. where A\> is the difference in the X7 of subsequent parameters sets.," The new parameter set is acceptedwith probability, where $\Delta\chi^2$ is the difference in the $\chi^2$ of subsequent parameters sets."
+ Note. the new parameter set is always accepted if its X7 is lower than the previous parameter set 2?= 1.," Note, the new parameter set is always accepted if its $\chi^2$ is lower than the previous parameter set $P=1$ )."
+ The jump functions are sealed by a common factor in order to ensure that 25% of the steps are accepted. as suggested by ?..," The jump functions are scaled by a common factor in order to ensure that $25\%$ of the steps are accepted, as suggested by \citet{Tegmark2004}."
+ To estimate the uncertainty of each parameter and calculate the jump functions. an initial MCMC fit was performed.," To estimate the uncertainty of each parameter and calculate the jump functions, an initial MCMC fit was performed."
+ With these jump functions. we computed seven MCMC chains each of 150 O00 points and different initial parameters.," With these jump functions, we computed seven MCMC chains each of 150 000 points and different initial parameters."
+ The initial of each chain that corresponded to the burn in phase were discarded and the remaining parts merged into a master chain., The initial of each chain that corresponded to the burn in phase were discarded and the remaining parts merged into a master chain.
+ We estimated the best fit parameter as the mode of its probability distribution and the | c limits as the value at which the integral of the distribution equals from both sides of the mode., We estimated the best fit parameter as the mode of its probability distribution and the 1 $\sigma$ limits as the value at which the integral of the distribution equals from both sides of the mode.
+ We computed the ? statistic for each titted parameter and concluded that chain convergence was good.," We computed the \citet{Gelman92}
+ statistic for each fitted parameter and concluded that chain convergence was good."
+ To test how the limb darkening coefficients affect the derived system parameters. we repeated the MCMC procedure also fitting for the linear LD “a” which is the most sensitive to the observing filter.," To test how the limb darkening coefficients affect the derived system parameters, we repeated the MCMC procedure also fitting for the linear LD $a$ ” which is the most sensitive to the observing filter."
+ The quadratic LD coefficient. b—0.210172. was kept fixed. because as reported by ?.. the high precision of the RISE light curves is not enough to fully constrain the LD coefficients (i.e.. the MCMC does not converge when fitting both coefficients).," The quadratic LD coefficient, $b=0.210172$, was kept fixed, because as reported by \citet{Gibson2008}, the high precision of the RISE light curves is not enough to fully constrain the LD coefficients (i.e., the MCMC does not converge when fitting both coefficients)."
+ We restricted the linear LD coefficient to be the same for all the light curves since they were all taken with the same filter., We restricted the linear LD coefficient to be the same for all the light curves since they were all taken with the same filter.
+ Therefore. in the second MCMC procedure we titted 13 parameters.," Therefore, in the second MCMC procedure we fitted 13 parameters."
+ We estimated @=0.337+0.034.," We estimated $a = 0.337
+\pm 0.034$."
+ Comparing the fitted and fixed LD solutions. we concluded that although the fitted linear LD coefficient is statistically significantly different from the theoretical value. this does not affect the derived system parameters to any extent.," Comparing the fitted and fixed LD solutions, we concluded that although the fitted linear LD coefficient is statistically significantly different from the theoretical value, this does not affect the derived system parameters to any extent."
+ The two solutions are within 1.5c of each other., The two solutions are within $1.5\sigma$ of each other.
+ Contrary to what was found by other authors (e.g. ?2)). the uncertainties of the fitted LD solution are slightly smaller than those of the fixed solution.," Contrary to what was found by other authors (e.g., \citealt{Gibson2008,Southworth2008}) ), the uncertainties of the fitted LD solution are slightly smaller than those of the fixed solution."
+ The 4? of the fitted LD solution is similar to the fixed LD solution which does not justify the addition of an extra free parameter in the fit., The $\chi^2$ of the fitted LD solution is similar to the fixed LD solution which does not justify the addition of an extra free parameter in the fit.
+ Consequently. we conclude that our light curve is of insufficient quality to better constrain the linear LD relativeto that achieved by theoretical models and we choose to present the fixed LD solution.," Consequently, we conclude that our light curve is of insufficient quality to better constrain the linear LD relativeto that achieved by theoretical models and we choose to present the fixed LD solution."
+ The estimated transit times. combined the original ephemeris (?) were used to update the linear ephemeris. We found [P-—4.3225060+0.0000031 and 7532455084.51974+0.00020 which was set to the mid transit time of the 2009 light curve.," The estimated transit times, combined the original ephemeris \citep{Bouchy2010} were used to update the linear ephemeris, We found $P=4.3225060 \pm 0.0000031$ and $T_0 =2455084.51974 \pm
+0.00020$ which was set to the mid transit time of the 2009 light curve."
+ This ephemeris was used in the tinal MCMC procedures., This ephemeris was used in the final MCMC procedures.
+ For future reference. the time residuals from the linear ephemeris are given in Table |..," For future reference, the time residuals from the linear ephemeris are given in Table \ref{ttv}."
+ We conclude that the time residuals of WASP-2Ib are consistent with a linear ephemeris., We conclude that the time residuals of WASP-21b are consistent with a linear ephemeris.
+ The geometric system parameters of WASP-2] and the le uncertainties derived from the MCMC analysis with fixed limb darkening coefficients are given in Table 2.., The geometric system parameters of WASP-21 and the $\sigma$ uncertainties derived from the MCMC analysis with fixed limb darkening coefficients are given in Table \ref{mcmc}.
+ These parameters are directly measured from the transit light curve and are only weakly dependent of stellar properties through the limb darkening coefficients., These parameters are directly measured from the transit light curve and are only weakly dependent of stellar properties through the limb darkening coefficients.
+ Note. all the derived parameters presented in Table were calculated at each point of the chain.," Note, all the derived parameters presented in Table \ref{mcmc} were calculated at each point of the chain."
+ Therefore. the final derived values and errors were determined from their probability distribution as done for the fitted values.," Therefore, the final derived values and errors were determined from their probability distribution as done for the fitted values."
+ We obtain a significantly lower density than was previous reported in the discovery paper ps=0844O.09p. )., We obtain a significantly lower density than was previous reported in the discovery paper $ \rho_* = 0.84 \pm 0.09 \rho_{\odot} $ ).
+ To obtain the stellar and planetary physical properties. the geometric parameters have to be scaled with the stellar mass.," To obtain the stellar and planetary physical properties, the geometric parameters have to be scaled with the stellar mass."
+ The new high quality transit light curves give a direct estimate of the stellar density., The new high quality transit light curves give a direct estimate of the stellar density.
+ This allows a more accurate estimation of stellar mass than log g derived from spectral analysis (2).., This allows a more accurate estimation of stellar mass than log g derived from spectral analysis \citep{Sozzetti2007}.
+ Currently there are two main methods to derive the stellar mass from the stellar density., Currently there are two main methods to derive the stellar mass from the stellar density.
+ The first uses isochrones and mass tracks from stellar models (2). and the second uses an empirical calibration derived from stellar eclipsing binaries (22).," The first uses isochrones and mass tracks from stellar models \citep{Sozzetti2007} and the second uses an empirical calibration derived from stellar eclipsing binaries \citep{Torres2010,Enoch2010}."
+ ? derived the stellar mass through the empirical calibration between Z;.rr. p. and [Fe/H] (?). with the parametrisation of ?..," \citet{Bouchy2010} derived the stellar mass through the empirical calibration between $T_{eff}$, $ \rho_* $ and [Fe/H] \citep{Torres2010} with the parametrisation of \citet{Enoch2010}."
+ Following the same procedure. with the improved p.. we derive a stellar mass of 1.02+0.05M .," Following the same procedure, with the improved $\rho_*$, we derive a stellar mass of $1.02 \pm 0.05 \Msun$ ."
+. In Table 3. we present the mass and radius of WASP-21 and WASP-21b and the Io uncertainties derived from the MCMC for a stellar mass of 1.02+0.05M .," In Table \ref{mcmc2} we present the mass and radius of WASP-21 and WASP-21b and the $\sigma$ uncertainties derived from the MCMC for a stellar mass of $1.02 \pm
+0.05 \Msun$ ."
+. We obtain a signiticantly larger stellar and planetary radius than, We obtain a significantly larger stellar and planetary radius than
+fully consistent with our Hix equations.,fully consistent with our flux equations.
+" Using the cosmology OQ,= 027. Oy=0.73 and Hubble parameter fy=71 km s+ AlpeJ|. they have ροκ=1.635 and +=0.835 (Metzgeretal. 1997)... leacing. at fang=23.3 days. to voy=921:107 Hz. 5,4=426-101 Hz. |,=0.756 niJv. p—2.22 and &=0.0307."," Using the cosmology $\Omega_M = 0.27$ , $\Omega_\Lambda=0.73$ and Hubble parameter $H_0=71$ km $^{-1}$ $^{-1}$, they have $r_{obs,28} = 1.635$ and $z =0.835$ \citep{Metzger1997}, , leading, at $t_{obs,d} = 23.3$ days, to $\nu_{c,1} = 9.21\cdot10^{13}$ Hz, $\nu_{m,1} = 4.26\cdot10^{10}$ Hz, $F_{m,1} = 0.756$ mJy, $p = 2.22$ and $k = 0.0307$."
+ Both VanderLlorstetal.(2008) and Galamaοἱal.(1999) take for the hydrogen mass fraction of the circumburst medium Vo=0.7. which in our Ilux equations is mathematically equivalent (though conceptually different) to setting £y=(1|X)/20.85.," Both \citet{vanderHorst2008} and \citet{Galama1999} take for the hydrogen mass fraction of the circumburst medium $X=0.7$, which in our flux equations is mathematically equivalent (though conceptually different) to setting $\xi_N = (1+X)/2 = 0.85$."
+ Unfortunately this still leaves us with four variables to determine (cp. CE. I. no) and only three constraints (peak flux. cooling ancl peak frequeney).," Unfortunately this still leaves us with four variables to determine $\epsilon_B$, $\epsilon_E$, $E_{52}$ , $n_0$ ) and only three constraints (peak flux, cooling and peak frequency)."
+ From a theoretical studs of the microstructure of collisionless shocks \ledvedey(2006) arrives at the following Constraint: We include this constraint to have a closed: set. οἱ equatIons., From a theoretical study of the microstructure of collisionless shocks \citet{Medvedev2006} arrives at the following constraint: We include this constraint to have a closed set of equations.
+ For the values quoted above we obtain: £5»=0.155. Do 19δ. ep=0.1057. cep=0.325.," For the values quoted above we obtain: $E_{52} = 0.155$, $n_0 = 1.28$, $\epsilon_B = 0.1057$, $\epsilon_E = 0.325$."
+ In figures 2. and 53. we have plotted a comparison between the spectrum generated by using these values as input parameters. [or the DM solution and the spectrum as it is represented. by applying theresults of the broadband fit of for the values ofthe critical frequencies and the peak flux to equation (11))., In figures \ref{GRB970508_spectrum_figure} and \ref{GRB970508_lightcurves_figure} we have plotted a comparison between the spectrum generated by using these values as input parameters for the BM solution and the spectrum as it is represented by applying theresults of the broadband fit of for the values ofthe critical frequencies and the peak flux to equation \ref{F1smooth_cooling_equation}) ).
+ Our scaling cocllicicnts were fixed for arbitrary& and, Our scaling coefficients were fixed for arbitrary$k$ and
+ (Shu.Adams., \citep{sal87}.
+&Lizauo1987).. Bodenheiuer1995 (v, \citealp{b95} \citealp{j07}) \citep{dj08}.
+anJorgensen2008).. Enochetal.(2009) 0141. Ulrich (1976)) al.(2009) Πασά diniensiouless miass-to-flux ratio A a few: Trolaud&Crutcher 2008)) is coutrolled. by maguetie braking. which has been difficult to quantify uutil recently.," \citet{ecdd09} $0.1\msun$ \citeauthor{u76}' \citeyear{u76}) \citet{ecdd09} mean dimensionless mass-to-flux ratio $\lambda \sim$ a few; \citealp{tc08}) ) is controlled by magnetic braking, which has been difficult to quantify until recently."
+ Detailed calculations have now shown that the magnetic braking is apparently so efücieut as to iulibit the formation of rotationally supported disks for a moderate level of core naenetization. as long as the maeuetic field is frozen in the matter (i.c.. the ideal MIID limit: Allen. 2003: Callietal. 2006:: Price&Bate 2007:: ITennebe&Fromaneg 2008: Mellon&Li 2008: see. howew IIleuuebelle&Ciardi2009. for a different view. aud al Machidaetal.2005 ancl Banerjee&Pudritz2006 for t evolution of rotating. maenetized cores prior to the main accretion phase).," Detailed calculations have now shown that the magnetic braking is apparently so efficient as to inhibit the formation of rotationally supported disks for a moderate level of core magnetization, as long as the magnetic field is frozen in the matter (i.e., the ideal MHD limit; \citealp{als03}; \citealp{glsa06}; \citealp{pb07}; \citealp{hf08}; \citealp{ml08}; see, however, \citealp{hc09} for a different view, and also \citealp{mmth05} and \citealp{bp06} for the evolution of rotating, magnetized cores prior to the main accretion phase)."
+ In order for a rotationally supported disk to exist around a rapidly accreting protostar. the magnetic braking nust be weakened one way or another.," In order for a rotationally supported disk to exist around a rapidly accreting protostar, the magnetic braking must be weakened one way or another."
+ The most obvious possibility is through non-deal MITD effects which are expected to be important in lehthy ionized star-forming cores.," The most obvious possibility is through non-ideal MHD effects, which are expected to be important in lightly ionized star-forming cores."
+ However. we have recently shown that ambipolar diffusion. the best studied uou-ideal MIID effect iu star formation. docs not sufficiently weaken the braking to allow rotationally supported disks to form for realistic levels of core maguctization aud ionization: iu some cases. the magnetic braking is even ibanced (Mellon&Li2009.. Krasnopolsky&EKóuiel 2002:: see Dasu&Mouschovias1995.. IToskiug&Whit-worth2001. and Duffiu&Pudritz2009.. who cousicered the effects of ambipolar diffusion on magnetic brakiug and core fragmentation during the eulier phase of prestellar core evolution).," However, we have recently shown that ambipolar diffusion, the best studied non-ideal MHD effect in star formation, does not sufficiently weaken the braking to allow rotationally supported disks to form for realistic levels of core magnetization and ionization; in some cases, the magnetic braking is even enhanced \citealp{ml09}, \citealp{kk02}; see \citealp{bm95}, \citealp{hw04}
+ and \citealp{dp09}, who considered the effects of ambipolar diffusion on magnetic braking and core fragmentation during the earlier phase of prestellar core evolution)."
+ This motivates us to examine whether two of the remaining non-ddeal MIID effects. olunic dissipation aud the Tall effect. can weaken the couplingbetween the magnetic fieldauc the bulk neutral matter (and thus the streneth of maeguetic braking) enough to enable the formation of rotationallysupported disks.," This motivates us to examine whether two of the remaining non-ideal MHD effects, ohmic dissipation and the Hall effect, can weaken the couplingbetween the magnetic fieldand the bulk neutral matter (and thus the strength of magnetic braking) enough to enable the formation of rotationallysupported disks."
+ This paper concentrates ou the olunic dissipation., This paper concentrates on the ohmic dissipation.
+"this is actually true in our simulations. we examined the central densities and masses of our halos at +=25: we find that all halos except for one have masses below 10? AZ. and central densities below — 2  tone halo has AZ=10A7. and p,=Y107 H "")","this is actually true in our simulations, we examined the central densities and masses of our halos at $z=25$; we find that all halos except for one have masses below $10^6$ $\Msun$ and central densities below $\sim$ 2 $^{-3}$ (one halo has $M > 10^6 \Msun$ and $n_c = 7 \times 10^3$ $^{-3}$ )."
+ Our halo abundance is consistent with predictions from the conditional Press-Schechter formalism (MBHO6)., Our halo abundance is consistent with predictions from the conditional Press-Schechter formalism (MBH06).
+ We chose this approach because we are interested in investigating the impact of relic HII regions on the formation of structure which forms significantly later. rather than studying the impact of a Pop III star on a nearby. highly collapsed halo.," We chose this approach because we are interested in investigating the impact of relic HII regions on the formation of structure which forms significantly later, rather than studying the impact of a Pop III star on a nearby, highly collapsed halo."
+ In particular. as noted in the introduction. our motivation is to follow up on the suggestion of ? that relic HII regions will imprint an excess entropy on the low-density gas at early times. which will then suppress later star formation.," In particular, as noted in the introduction, our motivation is to follow up on the suggestion of \citet{OH03} that relic HII regions will imprint an excess entropy on the low-density gas at early times, which will then suppress later star formation."
+ Of course. we stress that it is still important to examine our results with a more realistic RHD treatment.," Of course, we stress that it is still important to examine our results with a more realistic RHD treatment."
+ The results discussed in the section on positive feedback (section 4) are even more robust because in that case the halos form out of gas that was barely overdense when the ionization occurred., The results discussed in the section on positive feedback (section \ref{sec:pos}) ) are even more robust because in that case the halos form out of gas that was barely overdense when the ionization occurred.
+ As can be seen in Fig. 7..," As can be seen in Fig. \ref{fig:NoUVB_vs_Heat},"
+ even at >=21 (long after the transient ionization). the central density of the halo Gin the noUVB case) is only 0.2 ," even at $z=21$ (long after the transient ionization), the central density of the halo (in the noUVB case) is only 0.2 $^{-3}$."
+ Another useful comparison of our results can be made with ?.. who used a full cosmological simulation of Pop III star formation including radiative transfer (albeit in a region considerably smaller han considered there).," Another useful comparison of our results can be made with \citet{WA08b}, who used a full cosmological simulation of Pop III star formation including radiative transfer (albeit in a region considerably smaller than considered there)."
+ While a direct comparison is difficult because ? do not explicitly measure the suppression due to photo-ionization. a number of conclusions can be drawn.," While a direct comparison is difficult because \citet{WA08b} do not explicitly measure the suppression due to photo-ionization, a number of conclusions can be drawn."
+ First. it is clear rom Fig.," First, it is clear from Fig."
+ 7 of that paper that they also see the evaporation of gas Tom low-mass halos and filaments. and the resulting density field in the two simulations looks remarkably similar.," 7 of that paper that they also see the evaporation of gas from low-mass halos and filaments, and the resulting density field in the two simulations looks remarkably similar."
+ In addition. the ime difference between the first Pop TT object and the formation of he second is about 80 Myr. consistent with the time-delay derived in this paper.," In addition, the time difference between the first Pop III object and the formation of the second is about 80 Myr, consistent with the time-delay derived in this paper."
+ On the other hand. one important point that comes out of that paper is that. because Pop III star formation tends to be clustered. relic HII regions do not sit isolated for longer periods of time and so a region of space may undergo multiple ionization episodes. something which is not considered in this work.," On the other hand, one important point that comes out of that paper is that, because Pop III star formation tends to be clustered, relic HII regions do not sit isolated for longer periods of time and so a region of space may undergo multiple ionization episodes, something which is not considered in this work."
+ Next. we turn to LW photons.," Next, we turn to LW photons."
+ First. we remind the reader that we are discussing a persistent LW background. and that although the LW flux is produced by the same of objects that generate the transient UV flux. it is not produced by the same physical objects.," First, we remind the reader that we are discussing a persistent LW background, and that although the LW flux is produced by the same of objects that generate the transient UV flux, it is not produced by the same physical objects."
+ As discussed in more detail in 2.. the LW background is accumulated from a large number of distant sources.," As discussed in more detail in \citet{HRL96}, the LW background is accumulated from a large number of distant sources."
+ The effective optical depth in the LW lines in the background IGM. over a Hubble distance. is at most 7zc1—2 (??). so the Olbers” integral for the LWB is dominated by the large number of sources at a good fraction of the Hubble distance.," The effective optical depth in the LW lines in the background IGM, over a Hubble distance, is at most $\tau\approx 1-2$ \citep{HAR00,RGS01}, so the Olbers' integral for the LWB is dominated by the large number of sources at a good fraction of the Hubble distance."
+ In comparison. the ionizing photons come exclusively from the much fewer source(s) nearby. within the local HII region.," In comparison, the ionizing photons come exclusively from the much fewer source(s) nearby, within the local HII region."
+ LW photons can be self-shielded by a sufficient column of H»., LW photons can be self-shielded by a sufficient column of $_2$.
+ Here again. there are two regimes.," Here again, there are two regimes."
+ One is the large quantities of low-density gas in the relic HII regions that form substantial amounts of molecular hydrogen immediately after the ionizing flux is turned off (e.g. 2)., One is the large quantities of low-density gas in the relic HII regions that form substantial amounts of molecular hydrogen immediately after the ionizing flux is turned off (e.g. \citealt{KM05}) ).
+ Although H» formed in this manner may slow the build-up of a persistent LW background. it does not — for a given LW background — provide a large amount of shielding.," Although $_2$ formed in this manner may slow the build-up of a persistent LW background, it does not – for a given LW background – provide a large amount of shielding."
+ For example. ? find typical values of LW optical depths of order unity for a single relic HII region.," For example, \citet{JGB07} find typical values of LW optical depths of order unity for a single relic HII region."
+ This means that the LW flux inside the relic HII region can be reduced by a factor of a few. but it also means that it doesn’t take long for the persistent LW background to “burn off” the H» produced in relic HII regions.," This means that the LW flux inside the relic HII region can be reduced by a factor of a few, but it also means that it doesn't take long for the persistent LW background to ""burn off"" the $_2$ produced in relic HII regions."
+ The second regime occurs after the ionized gas falls back into the collapsed halos. where H» is in equilibrium between formation and photo-destruction.," The second regime occurs after the ionized gas falls back into the collapsed halos, where $_2$ is in equilibrium between formation and photo-destruction."
+ In this case. when a sufficient column density of H» is formed. it can effectively attenuate the LW flux.," In this case, when a sufficient column density of $_2$ is formed, it can effectively attenuate the LW flux."
+ Based on our simulations. we find that when CD gas forms. the LW optical depth can be of order unity or larger. indicating that self-shielding can play a role during the recollapse (although probably not during the photoionization process itself).," Based on our simulations, we find that when CD gas forms, the LW optical depth can be of order unity or larger, indicating that self-shielding can play a role during the recollapse (although probably not during the photoionization process itself)."
+ Therefore. we caution readers that the critical LW values we give above may need to be increased when self-shielding is properly accounted for.," Therefore, we caution readers that the critical LW values we give above may need to be increased when self-shielding is properly accounted for."
+ Finally. we note that in simulations with full radiative transfer. the photoionizing flux can also vary across a given halo trelic H II regions have radii of 2-3 kpe. while the Lagrangian radius of a typical halo in our simulations is of order 0.5 kpe).," Finally, we note that in simulations with full radiative transfer, the photoionizing flux can also vary across a given halo (relic H II regions have radii of 2-3 kpc, while the Lagrangian radius of a typical halo in our simulations is of order 0.5 kpc)."
+ This also means that. while most forming halos lie either completely inside or outside a relic H II region. some proto-halos could be partially ionized.," This also means that, while most forming halos lie either completely inside or outside a relic H II region, some proto-halos could be partially ionized."
+ UV radiation from early astrophysical sources could have a large impact on subsequent star formation in nearby protogalaxies. and in general on the progress of cosmological reionization.," UV radiation from early astrophysical sources could have a large impact on subsequent star formation in nearby protogalaxies, and in general on the progress of cosmological reionization."
+ Theoretical arguments based on the absence of metals in the early Universe suggest that the first stars were likely massive. bright. yet short-ived. with lifetimes of a few million years.," Theoretical arguments based on the absence of metals in the early Universe suggest that the first stars were likely massive, bright, yet short-lived, with lifetimes of a few million years."
+ Here we study he impact of the transient radiation arising from such stars on early protogalaxies., Here we study the impact of the transient radiation arising from such stars on early protogalaxies.
+ We apply the same statistical approach as in IBHO6. studying various combinations of transient UVBs and xersistent LWBs. using the hydrodynamical simulation code. Enzo.," We apply the same statistical approach as in MBH06, studying various combinations of transient UVBs and persistent LWBs, using the hydrodynamical simulation code, Enzo."
+ We also study a more typical and relevant region whose proto-galaxies form at lower redshifts. +~ 13-20.," We also study a more typical and relevant region whose proto-galaxies form at lower redshifts, $z\sim$ 13–20."
+ This allows us to trace eedback effects on longer time-scales and lower redshifts., This allows us to trace feedback effects on longer time-scales and lower redshifts.
+ We confirm the results of MBHO6 that feedback in the relic HII regions resulting from such transient radiation. is itself ransient.," We confirm the results of MBH06 that feedback in the relic HII regions resulting from such transient radiation, is itself transient."
+ Feedback effects dwindle away after ~30% of the Hubble time. and the same critical specific intensity of Jui~Kl10“oresstem=Uzlsr| separates. positive and negative feedback regimes.," Feedback effects dwindle away after $\sim30\%$ of the Hubble time, and the same critical specific intensity of $J_{\rm UV} \sim 0.1 \times 10^{-21}{\rm ergs~s^{-1}~cm^{-2}~Hz^{-1}~sr^{-1}}$ separates positive and negative feedback regimes."
+ A weaker UVB stimulates subsequent star formation inside the relic HIT regions. by enhancing the molecular hydrogen abundance.," A weaker UVB stimulates subsequent star formation inside the relic HII regions, by enhancing the molecular hydrogen abundance."
+ A stronger UVB delays. star-formation by reducing the gas density at the centers of collapsing halos., A stronger UVB delays star-formation by reducing the gas density at the centers of collapsing halos.
+ The fact that we have contirmed the results in MBHO06 (which were mostly inferences obtained by extrapolating from higher redshifts since we were unable to continue those simulations to low enough redshifts) suggests that overall feedback is fairly insensitive to the large-scale environment. overdensity. and redshift-dependent halo parameters. and can accurately be modeled in this regime with just the intensity of the impinging UVB. 5..," The fact that we have confirmed the results in MBH06 (which were mostly inferences obtained by extrapolating from higher redshifts since we were unable to continue those simulations to low enough redshifts) suggests that overall feedback is fairly insensitive to the large-scale environment, overdensity, and redshift-dependent halo parameters, and can accurately be modeled in this regime with just the intensity of the impinging UVB. \ref{sec:ss},"
+As anticipated briefly in Sect. 2.1..,"As anticipated briefly in Sect. \ref{empiric},"
+ the empirical correlation (4)) found between the two relevant scale lengths /r and Ro allows us to make a preliminary. test by focusing on a single. representative idealized case.," the empirical correlation \ref{corr}) ) found between the two relevant scale lengths $h$ and $R_{\Omega}$ allows us to make a preliminary test by focusing on a single, representative idealized case."
+ Therefore. a discussion of specific observed objects. such as NGC 3198. is postponed and will be carried out separately (see following Sect. ??))," Therefore, a discussion of specific observed objects, such as NGC 3198, is postponed and will be carried out separately (see following Sect. \ref{3198}) )"
+ As a typical reference case. we consider à galaxy characterized by a purely exponential stellar disk. exp(—£). and by the following rotation curve If we take the correlation found in Eq. (1) ," As a typical reference case, we consider a galaxy characterized by a purely exponential stellar disk, $\hat{\Sigma}(R) \equiv
+\exp(-\xi)$ , and by the following rotation curve If we take the correlation found in Eq. \ref{corr}) )"
+is reproduced., is reproduced.