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of Education. | of Education. |
BUD thanks CCC and the Alelchior famuly or their role m this work. | BHD thanks GCC and the Melchior family for their role in this work. |
Finally. we thank the referee or an eujovable aud helpful review process. | Finally, we thank the referee for an enjoyable and helpful review process. |
The results for spherically sxauinetric iodels are nuportaut to us both because they can be treated to sole extent with analytic techniques aud because direct conrparisous are possible with ERSAL | The results for spherically symmetric models are important to us both because they can be treated to some extent with analytic techniques and because direct comparisons are possible with ERSM. |
As a check on our computational accuracy. we first show our results for the mean intensity in a homogeneous spherical cloud compared with the analytic solution. | As a check on our computational accuracy, we first show our results for the mean intensity in a homogeneous spherical cloud compared with the analytic solution. |
The solution of the full equation of radiative transfer in a homogeneous sphere can be obtained analytically with the spherical harmonics method (Flanucry. Roberec. Rybicki 1980). | The solution of the full equation of radiative transfer in a homogeneous sphere can be obtained analytically with the spherical harmonics method (Flannery, Roberge, Rybicki 1980). |
In Appendix A we show that for a homogeneous cloud of radius A and ceutre-to-ccdec optica depth την the calculation of J,(r) from eq. (2)) | In Appendix A we show that for a homogeneous cloud of radius $R$ and centre-to-edge optical depth $\tau_{\nu c}$, the calculation of $J_\nu(r)$ from eq. \ref{Jrdef}) ) |
is straightforward aud can be expressed dn ΤΟΙΣ of exponeutialiuteeral functions of the arguimeut Ty(ll zor). where ww=P Approximate expressions for the mean intensity can be obtained in the Πιτ of very large or very σα τν | is straightforward and can be expressed in terms of exponential-integral functions of the argument $x_\pm=\tau_{\nu c}(1\pm x)$ , where $x= r/R$: Approximate expressions for the mean intensity can be obtained in the limit of very large or very small $\tau_{\nu c}$. |
Iu the optically thin limit eq. (3)) | In the optically thin limit eq. \ref{Jexact}) ) |
becomes whereas an approximate formula (diverging for.» 1) for the optically thick case is h(E | becomes, whereas an approximate formula (diverging for $x\rightarrow 1$ ) for the optically thick case is ( ). |
Iu Figure l.. we courpare the exact aud approximate analytical results with that obtained nunuerncallv integrating eq. (2)). | In Figure \ref{homsphere}, we compare the exact and approximate analytical results with that obtained numerically integrating eq. \ref{Jrdef}) ). |
We see that the nunerical code reproduces the analytic result very well | We see that the numerical code reproduces the analytic result very well. |
Thus. m simple ecolctrics. the procedure used for the angle integration is adequate. | Thus, in simple geometries, the procedure used for the angle integration is adequate. |
There are ereat) advantages im developing analytic formmlac which allow an approximate cstimate of the temperature iu deuse preprotostellar cores. | There are great advantages in developing analytic formulae which allow an approximate estimate of the temperature in dense pre–protostellar cores. |
This permits for example ai ciscussion of scaling relationships for different external radiatiou fields and opacities. | This permits for example a discussion of scaling relationships for different external radiation fields and opacities. |
We have Maccordinglv attempted to derive an analytical estimate of the temperature at the ceuter of a spherically sviuucetric loud for comparison with our πιοΊσα. results. | We have accordingly attempted to derive an analytical estimate of the temperature at the center of a spherically symmetric cloud for comparison with our numerical results. |
Tn order to develop au analytic approximation. we must first make some crude approximations concerning the dust opacity aud externa radiation field iu COYCS such as L1511. | In order to develop an analytic approximation, we must first make some crude approximations concerning the dust opacity and external radiation field in cores such as L1544. |
Followi18o BlacXA (1991) and MÀNp.t oeiterstellar radiaion field at wavelengths above loan Kam) be represcuted by the sui of four comvibulons: (1) an opticalNIR component peasing at Ay,= lun. due to the chussion of disk «wart and eiaut stars: (2) the diffuse FIR chussion from dst eralus. peasing at A,100 pan: (3) mudIR radiation from πια] nonthermally beate eralus in the range 51X) παπα 1) the cosmic backerounud radiation. peaking at A,=1 nuu. | Following Black (1994) and MMP, the interstellar radiation field at wavelengths above 0.1 $\mu$ m can be represented by the sum of four contributions: (1) an optical-NIR component peaking at $\lambda_{\rm p}=
1$ $\mu$ m, due to the emission of disk dwarf and giant stars; (2) the diffuse FIR emission from dust grains, peaking at $\lambda_{\rm p} =
100$ $\mu$ m; (3) mid–IR radiation from small non–thermally heated grains in the range 5–100 and (4) the cosmic background radiation, peaking at $\lambda_{\rm p}
=1$ mm. |
ComIOCits 1.2. and l can be represeited by a suele (or a sunm of) modified black-bodies at temperatures 7; of the kine where Ay is the peak swaveleugth aud IY; are dilution factors. | Components 1,2, and 4 can be represented by a single (or a sum of) modified black-bodies at temperatures $T_i$ of the kind )^p where $\lambda_{\rm p}$ is the peak wavelength and $W_i$ are dilution factors. |
The values of the parameters Ay. Wy) aud Ti are listed in Table D1.. | The values of the parameters $\lambda_{\rm p}$, $W_i$ and $T_i$ are listed in Table \ref{isf_param}. |
The third (MIB) component is found to be variable near dense cores such as LISLL aud cau be approximated as a power law. | The third (MIR) component is found to be variable near dense cores such as L1544 and can be approximated as a power law. |
Iu the following discussion however. we at first ucelect it and then make au estimate of its lmaportance. | In the following discussion however, we at first neglect it and then make an estimate of its importance. |
The solar iradiance. or the solar flux received at the op of the Earth’s atmosphere. is known to vary over a laree number of time scales. ranging from έτος o inonthns and decades. | The solar irradiance, or the solar flux received at the top of the Earth's atmosphere, is known to vary over a large number of time scales, ranging from minutes to months and decades. |
The changes iu the total solar output have been measured since 1978 (Willson&IImd-son1988) and different composites of the measurements iive been preseuted by Eróhlich&Lean(1998):Will-son&Mordvinov(2003) απ Dewitteetal.C2001). | The changes in the total solar output have been measured since 1978 \citep{willson88} and different composites of the measurements have been presented by \citet{frohlich98,willson2003}
and \citet{Dewitte2004}. |
. While the short-term: (uiuutes to hour) variability is uainlv due to solar oscillations auc eranulation. the daily to decadal variability is attributed to the changes in the surface maeuetic field combined with the solar rotation that transports solar active regions iuto aud out of view. | While the short-term (minutes to hour) variability is mainly due to solar oscillations and granulation, the daily to decadal variability is attributed to the changes in the surface magnetic field combined with the solar rotation that transports solar active regions into and out of view. |
Tncdeed. IRrivovaetal.(2003) oud hat more than of the solar variability between 1996 aud 2002 couk be explained by chiuges in the solar surface field. | Indeed, \citet{krivova2003_cycle23} found that more than of the solar variability between 1996 and 2002 could be explained by changes in the solar surface field. |
Similar conclusions were reached by Wenuzleretal.(2006) who reconstructed solar irracliance from itt Peal maenetoerams coverine the last 3 solar cvcles. | Similar conclusions were reached by \citet{wenzler2006} who reconstructed solar irradiance from Kitt Peak magnetograms covering the last 3 solar cycles. |
Solar variability is a strong function of wavelength: while solar output is small in the UV. the relative variability is more than oue order of magnitude larecr in he UV than in the visible. | Solar variability is a strong function of wavelength: while solar output is small in the UV, the relative variability is more than one order of magnitude larger in the UV than in the visible. |
Uutil very receutly. the spectral dependence of the solar variability had iainly been determined in the UV. in particular by the measurements aeu by the SUSIM aud SOLSTICE instruments onboard UARS (sec.e.g.àYFlovdetal.2003a). | Until very recently, the spectral dependence of the solar variability had mainly been determined in the UV, in particular by the measurements taken by the SUSIM and SOLSTICE instruments onboard UARS \citep[see, e.g.,][]{floyd-et-al-2003a}. |
. Tuformation iu he visible was restricted to the three colour channels of he SPM instrmucut of SOMO/VIRGO (Frohlichetal. 1995).. though degradation Lampered the use of these data ievoud timescales of the order of a fewmouthls!. | Information in the visible was restricted to the three colour channels of the SPM instrument of SOHO/VIRGO \citep{frohlich95}, though degradation hampered the use of these data beyond timescales of the order of a few. |
. The variability at most other wavelengths had to be inferred usine a variety of approaches. such as ¢.e.. pioueere by Lean(1989) who produced he first estimate of solar-cvcle variability over a large wavelength range. | The variability at most other wavelengths had to be inferred using a variety of approaches, such as e.g., pioneered by \citet{lean89} who produced the first estimate of solar-cycle variability over a large wavelength range. |
Au alternative approach was followed by Unruhetal.(1999) who usec facular aud spot model atinospheres to calculate the dux changes due to maguctic features. | An alternative approach was followed by \citet{unruh99lumi} who used facular and spot model atmospheres to calculate the flux changes due to magnetic features. |
Fligeeetal.(2000) and I&aivovactal.(2003) used solar surface Huaees and maegnetoerams to calculate the variability on time scales ranging from days to vears. | \citet{fligge2000irrad} and \citet{krivova2003_cycle23}
used solar surface images and magnetograms to calculate the variability on time scales ranging from days to years. |
Were we built on this approach and preseut comparisons between modelled and measured spectral imaciauces during three mouths in 9001. | Here we built on this approach and present comparisons between modelled and measured spectral irradiances during three months in 2004. |
Thaulks to missions such as SORCE aud SCIAMACIIY the observational outlook has now become much better and we have. for the first time. variability observations that span from the UV to the near IR (larderetal.2005101Rottimanetal.2005:Skupin 2005). | Thanks to missions such as SORCE and SCIAMACHY the observational outlook has now become much better and we have, for the first time, variability observations that span from the UV to the near IR \citep{harder2005calib,rottman2005sorce,skupin-et-al-2005}. |
. In the following we consider SORCE data oulv. | In the following we consider SORCE data only. |
First comparisous between SORCE measurements and models have been preseuted by. e.g.. Fouteulaetal.(200L) and Leanetal.(2005).. | First comparisons between SORCE measurements and models have been presented by, e.g., \citet{fontenla2004sorce} and \citet{lean2005sorce}. |
All data prescuted here have beeu recorded between 21 April and 1 Aneust 2001. | All data presented here have been recorded between 21 April and 1 August 2004. |
During this time he Stun was in a relativolv quiet phase. especially in May when ouly a very snall spot eroup appeared ou the solar disk. | During this time the Sun was in a relatively quiet phase, especially in May when only a very small spot group appeared on the solar disk. |
A new and Lueer active region eimerged over the next mouth. resulting in a depression of just over 1 permuille iu total solar irradiance (TST) im July. | A new and larger active region emerged over the next month, resulting in a depression of just over 1 permille in total solar irradiance (TSI) in July. |
lu the next section we briefly describe our irradiance modelling approach. | In the next section we briefly describe our irradiance modelling approach. |
We then discuss the data analysis for the differeut instruments (Sec. 3)). | We then discuss the data analysis for the different instruments (Sec. \ref{sec:data}) ). |
Iu Sec. L. | In Sec. \ref{sec:comps}, |
we conpare the relative inradiauce changes derived frou the models with a nuuber of different data sets spanning a waveleugth ranee from 200 to 1600 nm. | we compare the relative irradiance changes derived from the models with a number of different data sets spanning a wavelength range from 200 to 1600 nm. |
In particular. we conipare our inodel to data from SORCE/SIML. UARS/SUSIAL and SoIIO/VIRGO. | In particular, we compare our model to data from SORCE/SIM, UARS/SUSIM, and SoHO/VIRGO. |
. We conclude this section by presenting observed aud modelled. timeserics | We conclude this section by presenting observed and modelled timeseries |
vdWO09 observed a ceiling mass of ~Lott3f. for disk-domunated. quiescent galaxies iu the prescut-day universe. | vdW09 observed a ceiling mass of $\sim 10^{11}\ M_{\sun}$ for disk-dominated, quiescent galaxies in the present-day universe. |
In accordance. Bernardietal.(2010) fouud that early-type galaxies with verv high masses (2xLott AL.) differ iu niv wars from those with lower masses (<104+ AL). | In accordance, \citet{bernardi2010} found that early-type galaxies with very high masses $\sim 2\times10^{11}\ M_{\sun}$ ) differ in many ways from those with lower masses $<10^{11}\ M_{\sun}$ ). |
In this paper we show that a similar transition mass exists at 5—0.7 and that its value has uot shifted by more than 0.05 dex between +~0.7 and the preseut. | In this paper we show that a similar transition mass exists at $z\sim 0.7$ and that its value has not shifted by more than 0.05 dex between $z\sim 0.7$ and the present. |
Thus. at all redshifts. roughly of the stellar mass in early-type systems is contained in these relatively round systems. | Thus, at all redshifts, roughly of the stellar mass in early-type systems is contained in these relatively round systems. |
As expected from Figure LL. round οντοις (raj= 0.6) have a characteristic mass of M*~9«10/9AL. while highly flattened (proj<O.L). passively evolving galaxieshave AZ*~Lx1019AL. | As expected from Figure \ref{rix}, round systems $q_{proj} > 0.6$ ) have a characteristic mass of $M^{\star} \sim 9 \times 10^{10}\
M_{\sun}$ while highly flattened $q_{proj}<0.4$ ), passively evolving galaxieshave $M^{\star} \sim 4 \times 10^{10}\ M_{\sun}$. |
We find no significant evolution iu the value for AM between our two samples. ouly evolution iu the co-moving ΠΠΟΥ clensity. | We find no significant evolution in the value for $M^{\star}$ between our two samples, only evolution in the co-moving number density. |
This mass ceiling has the same mass. or. in other words. AL* does not evolve for more elongated or "diskv carly-type galaxies despite the erowth of the passively evolving population by a factor of ~2.3 in mass between :~1 and today (seeFigure&::2011). | This mass ceiling has the same mass, or, in other words, $M^{\star}$ does not evolve for more elongated or “disky” early-type galaxies despite the growth of the passively evolving population by a factor of $\sim2-3$ in mass between $z\sim1$ and today \citep[see Figure
\ref{mf};. |
This has two müplications. | This has two implications. |
First. the progenitors of today’s population of massive earlv-tvpe population were not more disk-dominated svstenis at 2=0.7 that faded into the passively-evolving population. | First, the progenitors of today's population of massive early-type population were not more disk-dominated systems at $z=0.7$ that faded into the passively-evolving population. |
lustead. these galaxies must already be almost round. roughly 2:3 in intrinsic axis-ratio. galaxies before the truucation of star-formation (seeI&ocevskietal.2010.forcandidate progenitors).. | Instead, these galaxies must already be almost round, roughly 2:3 in intrinsic axis-ratio, galaxies before the truncation of star-formation \citep[see][for candidate
progenitors]{kocevski2011}. |
Second. if ποοΊο builds up the population of galaxies above ~LOYAY... that inergiug lost cause them to become rounder svstenis. | Second, if merging builds up the population of galaxies above $\sim 10^{11}\ M_{\sun}$, that merging most cause them to become rounder systems. |
At the highest masses. we find that. not only is here a lack of fattened or "diskv galaxies. but that he distribution is consistent with a larecly triaxial sopulation. | At the highest masses, we find that, not only is there a lack of flattened or “disky” galaxies, but that the distribution is consistent with a largely triaxial population. |
This can be seen by the lack of galaxies hat are round in projection galaxies at high masses in Figure L. | This can be seen by the lack of galaxies that are round in projection galaxies at high masses in Figure \ref{rix}. |
These apparently round galaxies are seen at ower luasses. so we do know that the lower fraction of high mass. round galaxies not just a systematic neasurelent error. | These apparently round galaxies are seen at lower masses, so we do know that the lower fraction of high mass, round galaxies not just a systematic measurement error. |
Padilla&Strauss(2008) found a simular result. but the lack of evolution we find means hat this triaxialitv is set in the formation of these systems out to z~1. | \citet{padilla2008} found a similar result, but the lack of evolution we find means that this triaxiality is set in the formation of these systems out to $z\sim1$. |
This result. when combined with he observed evolution in the normalization of the mass Muction. and hints about the shape (see 1L2.0)). this axial population is built up over the redshift range we observe. but is done so in such a way to produce similar shaped systems over that time. | This result, when combined with the observed evolution in the normalization of the mass function, and hints about the shape (see \ref{roundmf}) ), this triaxial population is built up over the redshift range we observe, but is done so in such a way to produce similar shaped systems over that time. |
Massive ellipticals are assumed to form: out of multiple merecrs of near equal mass «πίστας aud the merecr rate is expected to be high even at redshifts of 7~O.7 (c.c.DeLucia&Dlaizot2007). | Massive ellipticals are assumed to form out of multiple mergers of near equal mass systems and the merger rate is expected to be high even at redshifts of $z\sim0.7$ \citep[e.g.][]{delucia2007}. |
. Detailed simulations with cosinological initial conditions show that additional niechanisnis are required to reproduce the observed shapes and kinematic profiles of massive ellipticals (6.8.Durkertetal.2008:Novak2008). | Detailed simulations with cosmological initial conditions show that additional mechanisms are required to reproduce the observed shapes and kinematic profiles of massive ellipticals \citep[e.g.][]{burkert2008,novak2008}. |
. Minor mergers aud tidal encouuters also provide a 11echauisui for making the ost luassive quiescent galaxies appear round. | Minor mergers and tidal encounters also provide a mechanism for making the most massive quiescent galaxies appear round. |
Vulcanietal.(2011) finds that the most massive cluster galaxies. objects too massive to be iu our saunple. are less round at high redshift. | \citet{vulcani2011b} finds that the most massive cluster galaxies, objects too massive to be in our sample, are less round at high redshift. |
This poiuts to observational evidence of the process of galaxies beconmiug rounder with tine. possibly because of the mechanisius suggested in Burkertetal.(2008).. but ouly for the rarest and most extreme of svstenis. | This points to observational evidence of the process of galaxies becoming rounder with time, possibly because of the mechanisms suggested in \citet{burkert2008}, , but only for the rarest and most extreme of systems. |
At masses <101AJ. the eurh-tvpe population becomes more aud more "diskv. | At masses $<10^{11}\ M_{\sun}$, the early-type population becomes more and more “disky”. |
This ean be seeu iu two wars. first. we fud more ronud galaxies. νο20.9. | This can be seen in two ways, first, we find more round galaxies, $q_{proj}>0.9$. |
Secoud. we fid more flattened systems. 4,,,;«0.1. | Second, we find more flattened systems, $q_{proj}<0.4$. |
This can be seen in both the minima axis-ratio we find iu Figure L. and. the distribution of ¢/a values we inter from our parametric modeling iu Figure 5.. | This can be seen in both the minimum axis-ratio we find in Figure \ref{rix}, and, the distribution of $c/a$ values we infer from our parametric modeling in Figure \ref{ctoa}. |
Quiescent galaxies with masses that dominate the cosmic stellar mass budget (3«100AZ,cM,101.AZ.) show a broad but nou-evolviug range in axis at both 2~0.7 aud +~0.06. | Quiescent galaxies with masses that dominate the cosmic stellar mass budget $3 \times10^{10}\ M_{\sun}<M_{\star}<10^{11}\ M_{\sun}$ ) show a broad but non-evolving range in axis-ratios, at both $z\sim 0.7$ and $z\sim 0.06$. |
The broad rauge m axis-ratios implies that the population cau form through a number of chanucls. | The broad range in axis-ratios implies that the population can form through a number of channels. |
Because we fiud so little evolution in the axi-ratios. however. means that. whatever the mechanisnis that formi carly-type galaxies iu this mass range. they mst have worked at similar rates across the last 7-8 Cars of look-back time. | Because we find so little evolution in the axis-ratios, however, means that, whatever the mechanisms that form early-type galaxies in this mass range, they must have worked at similar rates across the last 7-8 Gyrs of look-back time. |
This evolution caunot explained cutirely by the increase in the uunuber of bulec-dominated galaxies (sav. niergeer products). nor can it be explained eutirelv by the cessation of star formation in disk-donuünated galaxies without structural changes. | This evolution cannot explained entirely by the increase in the number of bulge-dominated galaxies (say, merger products), nor can it be explained entirely by the cessation of star formation in disk-dominated galaxies without structural changes. |
Several evolutionary processes that cause the formation of quiescent galaxies must contribute in order to explain the uuchangiug fractions of bulec- aud cdisk-domunatec quiescent ealaxies; | Several evolutionary processes that cause the formation of quiescent galaxies must contribute in order to explain the unchanging fractions of bulge- and disk-dominated quiescent galaxies. |
Moreover. the relative importance of the various evolutionary processes has not stronely changed over the past 7-8 Cors. | Moreover, the relative importance of the various evolutionary processes has not strongly changed over the past 7-8 Gyrs. |
This is reminiscent of the general result that the morphological mix of galaxies of these masses docs not siguificanthy change over the sale tine (vdWO.IT09.Duudyetal.2010). | This is reminiscent of the general result that the morphological mix of galaxies of these masses does not significantly change over the same time \citep[vdW07,
H09,][]{bundy2010}. |
. Our work finds consistent evolution in the nunboer density of passively evolving galaxies with redshift. | Our work finds consistent evolution in the number density of passively evolving galaxies with redshift. |
Most work finds significant evolution. factors of 2 or 3. in the nuuber density of galaxies in the mass range of our saniple(e.s.bertetal.2010:BundyBrauneretal. 2011). | Most work finds significant evolution, factors of 2 or 3, in the number density of galaxies in the mass range of our \citep[e.g.][]{ilbert2010,bundy2010,brammer2011}. |
. The combination of a flattened population at low masses with the increased nuuber density of galaxies with redshift says that. at lower masses. the builkiup of the mass function of passively evolving galaxies. or earlv-tvpes. is the build up of passive disk-like galaxies. such as SOs or c"diskw ellipticals (Duudyetal. 2010).. | The combination of a flattened population at low masses with the increased number density of galaxies with redshift says that, at lower masses, the buildup of the mass function of passively evolving galaxies, or early-types, is the build up of passive disk-like galaxies, such as S0s or “disky” ellipticals \citep{bundy2010}. . |
Tow can we explain the existence and continued erowtlof a population of quiesceut. flatteued. galaxies? | How can we explain the existence and continued growthof a population of quiescent, flattened galaxies? |
Cras stripping in eroup aud cluster euvironments has long been argued to play a role (Spitzer&Dade 1951).. and was recently shown to explain the existence of the morphologv-deusitv relation | Gas stripping in group and cluster environments has long been argued to play a role \citep{spitzer1951}, , and was recently shown to explain the existence of the morphology-density relation |
until 8.8 hours after the burst (?).. | until 8.8 hours after the burst \citep{rep_75}. |
The X-ray Telescope (XRT) observed the GRB location and detected a source at R.A = 17h 50m 58.495. Dec = -68° 557 277 with an uncertainty of 5.47 at the confidence level. within the error region. | The X-ray Telescope (XRT) observed the GRB location and detected a source at R.A = 17h 50m 58.49s, Dec = $-$ $^{\circ}$ 55' 27" with an uncertainty of 5.4" at the confidence level, within the error region. |
XRT carried out a further observation from 325—517 kks after the trigger. but the source had faded to below the XRT detection limit (the 3 sigma upper limit to the observed count rate was 0.011 counts/s. 2)). | XRT carried out a further observation from $325-517$ ks after the trigger, but the source had faded to below the XRT detection limit (the 3 sigma upper limit to the observed count rate was 0.011 counts/s, \citet{gcn6626}) ). |
Therefore it was not possible to estimate a break time from the lighteurve. | Therefore it was not possible to estimate a break time from the lightcurve. |
The X-ray spectrum over the interval Ty + 31.8 ks to To + 54.] ks was fit byan absorbed power-law with a photon index ϱ = -2.7 + 0.6. a fixed Galactic column density of 6 x107 cm and an average observed 0.3-10 keV flux of 245%)x 107 erg em s! (9). | The X-ray spectrum over the interval $_0$ + 31.8 ks to $_0$ + 54.1 ks was fit byan absorbed power-law with a photon index $\beta$ = $-$ 2.7 $\pm$ 0.6, a fixed Galactic column density of 6 $\times 10^{20}$ $^{-2}$ and an average observed 0.3–10 keV flux of $^{+2.0}_{-1.4})~ \times$ $^{-13}$ erg $^{-2}$ $^{-1}$ \citep{gcn6626}. |
The VLT observed the error region of 0070707 approximately 11 hours after the burst occurred (?) and found a single source with magnitude R~23.0 within the XRT error circle. | The VLT observed the error region of 070707 approximately 11 hours after the burst occurred \citep{gcn6612} and found a single source with magnitude $R \sim 23.0$ within the XRT error circle. |
This source faded by the second observation at ~ 34 hours post-trigger and was determined to be the optical afterglow (?).. | This source faded by the second observation at $\sim$ 34 hours post-trigger and was determined to be the optical afterglow \citep{gcn6613}. |
An upper limit was placed on the redshift of =<3.6 (S. Piranomonte. priv. | An upper limit was placed on the redshift of $z < 3.6$ (S. Piranomonte, priv. |
comm.) | comm.) |
due to the absence of a Lyman alpha limit. | due to the absence of a Lyman alpha limit. |
The host galaxy has an apparent magnitude of R=27.3 and is the faintest host detected so far for a short burst (e.g. the host galaxy of 0050709 had R~21 (?).. the host galaxy of 0050724 had K~15 (?) and the host galaxy of 0060121 had R=26.6 (2))). | The host galaxy has an apparent magnitude of $R = 27.3$ and is the faintest host detected so far for a short burst (e.g. the host galaxy of 050709 had $R\sim21$ \citep{hjorth2005}, the host galaxy of 050724 had $K\sim15$ \citep{berger05} and the host galaxy of 060121 had $R = 26.6$ \citep{levan2006}) ). |
The time history and spectral properties (e.g. Too. spectrallag and spectral shape) of the short burst 0070707 are similar to those of several short BATSE bursts (?) and bursts from the KONUS-Wind short GRB catalogue (?).. | The time history and spectral properties (e.g. $_{90}$, spectrallag and spectral shape) of the short burst 070707 are similar to those of several short BATSE bursts \citep{yuki2006} and bursts from the KONUS–Wind short GRB catalogue \citep{konus_cat}. |
The ratio of y-ray fluence to X-ray fluence (yo 400/X5-30) IS ~ 8. consistent with values found for BATSE short GRBs (?).. | The ratio of $\gamma$ -ray fluence to X-ray fluence $\gamma_{30-400}$ $_{2-30}$ ) is $\sim$ 8, consistent with values found for BATSE short GRBs \citep{ghir2004}. |
There are no direct measurements of the redshift of 0070707. | There are no direct measurements of the redshift of 070707. |
The isotropic peak luminosity L4:ice, can be calculated using the 50-300kkeV peak flux of the joint data from 3.2 (Puis=2xI07 ergcem ss?!) and assuming that 0070707 1s at the average redshift for short bursts. 1.8. z=0.35. | The isotropic peak luminosity $_{peak,iso}$ can be calculated using the keV peak flux of the joint data from \ref{s} $_{50-300} = 2 \times 10^{-7}$ $^{-2}$ $^{-1}$ ) and assuming that 070707 is at the average redshift for short bursts, i.e. $z = 0.35$. |
This yields a value for Ly,44, of LL x 107? ergs s. | This yields a value for $_{peak,iso}$ of 1.1 $\times$ $^{50}$ ergs $^{-1}$. |
The observed bolometric fluence kkeV) extrapolated from the power-law fit to the jointdata yields an isotropic equivalent energy Ej; of 1.8 x 10°! ergs. | The observed bolometric fluence keV) extrapolated from the power–law fit to the jointdata yields an isotropic equivalent energy $_{iso}$ of 1.8 $\times$ $^{51}$ ergs. |
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