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Fig. 3.. | Fig. \ref{Ly}, |
Table 1)) we derived a total cohunu density of log (ΠΤΙ) =20.6340.09 in close aereeciment with the Pettiui ct al. ( | Table \ref{HI-DI}) ) we derived a total column density of $\log N$ ) $= 20.63\pm 0.09$ in close agreement with the Pettini et al. ( |
1991) value of 20.720.1. | 1994) value of $20.7\pm 0.1$. |
Adding other components to the hydrogen fit with ow column densitics as the ones observed in the stronger uetal lines. componcut Lat 30 kin laud colmpoucuts 1. 5 and 6 at slightly higher redshifts than the two main coniponeuts. does not change siguificauth neither the otal column deusitv of the two main components nor he fit ou the blue aud red wines of the Lyman lines. | Adding other components to the hydrogen fit with low column densities as the ones observed in the stronger metal lines, component 1 at $-30$ km $^{-1}$ and components 4, 5 and 6 at slightly higher redshifts than the two main components, does not change significantly neither the total column density of the two main components nor the fit on the blue and red wings of the Lyman lines. |
Fie. | Fig. |
3 however clearly shows that the fit with three bydrogen coniponeuts systematically fails to reproduce the edge of he blue wing iu the higher members of the lydrogeu Lyman series. Lye. Lye. LylO iux Lv12. at about 55 aus I from the two ID main components (2 aud 3) which is the expected displacement of the corresponding D lines. | \ref{Ly} however clearly shows that the fit with three hydrogen components systematically fails to reproduce the edge of the blue wing in the higher members of the hydrogen Lyman series, $\epsilon$, Ly8, Ly10 and Ly12, at about $-82$ km $^{-1}$ from the two H main components (2 and 3) which is the expected displacement of the corresponding D lines. |
We then add to the model the deuteruu by assuming its contribution only to the two main components (2 and 3). | We then add to the model the deuterium by assuming its contribution only to the two main components (2 and 3). |
In the fitting procedure we assumed the same redshift | In the fitting procedure we assumed the same redshift |
into line radiation as discussed in the Appendix. with about of the continuum energy. or ~10 eres. coming out in Si Ik, photons. | into line radiation as discussed in the Appendix, with about of the continuum energy, or $\sim 10^{48}$ ergs, coming out in Si $_\alpha$ photons. |
We note that. because of projection effects. an observer who is located in the direction of the jet will infer a [actor of 2 higher line luminosity than the (rue isotropically averaged luminosity. | We note that, because of projection effects, an observer who is located in the direction of the jet will infer a factor of 2 higher line luminosity than the true isotropically averaged luminosity. |
This gives à safety [actor of 2 in the above estimates. | This gives a safety factor of $\sim2$ in the above estimates. |
Because the pair cloud is moving outward with a Lorentz factor ~2. its optical depth is likely (ο reduce substantially bv the time the line photons arrive at the cloud. | Because the pair cloud is moving outward with a Lorentz factor $\sim2$, its optical depth is likely to reduce substantially by the time the line photons arrive at the cloud. |
Therefore. the x-ray lines will not be significantly. attenuated by scattering off pairs. | Therefore, the x-ray lines will not be significantly attenuated by scattering off pairs. |
Note in addition that GRD radiation scattered. olf the pair cloud. associated with the counter-jet is likely to be blocked by the expanding supernova ejecta (for the parameters we have considered. viz.. 0;=0.2. | Note in addition that GRB radiation scattered off the pair cloud associated with the counter-jet is likely to be blocked by the expanding supernova ejecta (for the parameters we have considered, viz., $\theta_j=0.2$, |
transverse v/620.3). | transverse $v/c\gta0.3$ ). |
Thus. that radiation will not contribute to the observed continuum Πακ. | Thus, that radiation will not contribute to the observed continuum flux. |
The discovery of x-ray. lines in the allerelow spectra of GRBs. hours to davs alter the iniüial eamuna-rayw burst. poses severe problem [or theoretical models. | The discovery of x-ray lines in the afterglow spectra of GRBs, hours to days after the initial gamma-ray burst, poses severe problem for theoretical models. |
We have focused in (his paper on the strong Si line observed by Reeves et al. ( | We have focused in this paper on the strong Si line observed by Reeves et al. ( |
2002) in GRB 011211. but our argumentis are valid more generally ancl apply. lor instance. to the Fe line observed in GRD 991216 by Piro et al. ( | 2002) in GRB 011211, but our arguments are valid more generally and apply, for instance, to the Fe line observed in GRB 991216 by Piro et al. ( |
2000). | 2000). |
In the case of GRB 011211. we have the following constraints: 1. | In the case of GRB 011211, we have the following constraints: 1. |
The line-emittinge 0eas must be movinge towards the observer at an averagee speed of οO.1e relative to the rest [rame of the GRB. | The line-emitting gas must be moving towards the observer at an average speed of $v\sim0.1c$ relative to the rest frame of the GRB. |
This suggests that the gas is associated with ejecla from a supernova explosion. | This suggests that the gas is associated with ejecta from a supernova explosion. |
2. | 2. |
The time delay of ~107 s (measured in the source [rame) between the initial GRD and the later line emission constrains the geometry of the source. | The time delay of $\sim10^4$ s (measured in the source frame) between the initial GRB and the later line emission constrains the geometry of the source. |
In models in which the lines are produced by direct irradiation by GRB photons. the distance of the line-emitting region [rom (he central engine (2) and the lateral extent of the irradiated region (42) are constrained by equation (4). | In models in which the lines are produced by direct irradiation by GRB photons, the distance of the line-emitting region from the central engine $D$ ) and the lateral extent of the irradiated region $R$ ) are constrained by equation (4). |
In indirect irradiation models. in which the GRB radiation is first scattered off a nearby screen. the distance to the cloud is constrained to be ~10112 em and D is smaller ~1057? cem. | In indirect irradiation models, in which the GRB radiation is first scattered off a nearby screen, the distance to the cloud is constrained to be $\sim10^{14.5}$ cm and $D$ is smaller $\sim10^{13.5}$ cm. |
3. | 3. |
According to Reeves et al. ( | According to Reeves et al. ( |
2002). GRB OLI211 emitted 107 eres or more in each of five IX, lines. | 2002), GRB 011211 emitted $10^{48}$ ergs or more in each of five $_\alpha$ lines. |
This estimate of the energv is conservative since there might well have been additional line emission during the 11 hours between the time of the GRB and when the x-ray observations began. | This estimate of the energy is conservative since there might well have been additional line emission during the 11 hours between the time of the GRB and when the x-ray observations began. |
Regardless. the energv in the lines is a substantial fraction of the | Regardless, the energy in the lines is a substantial fraction of the |
least 45° (although the major mergers themselves represent only 0.396 of the total event distribution). | least $45\degr$ (although the major mergers themselves represent only $0.3\%$ of the total event distribution). |
We can also consider the CDF of Ay (Fig. | We can also consider the CDF of $\DM$ (Fig. \ref{f:dmfraccdf}) ). |
In this case, if we select just large flips, we find that the vast majorityB). (9396 of those with 02 45°) coincide with minor mergers (Ay< 0.3). labels | In this case, if we select just large flips, we find that the vast majority $93\%$ of those with $\theta \geq 45\degr$ ) coincide with minor mergers $\DM \leq 0.3$ ). |
:The distribution of flips of the inner halo angular momentum is more directly relevant when considering the stability of galaxies that might form within. | The distribution of flips of the inner halo angular momentum is more directly relevant when considering the stability of galaxies that might form within. |
The distribution of events as a function of the inner halo spin direction change the mass change of the whole halo is shown in Fig.D]. | The distribution of events as a function of the inner halo spin direction change the mass change of the whole halo is shown in Fig. \ref{f:MWinnerplots}. |
In comparison to the distribution for total halo spin flips, the inner halo exhibits a far greater spread to low-cos6. | In comparison to the distribution for total halo spin flips, the inner halo exhibits a far greater spread to $\cos\theta$. |
Cumulative distributions are shown in the middle and right panels of Fig.B]. | Cumulative distributions are shown in the middle and right panels of Fig. \ref{f:MWinnerplots}. |
We find that the frequency of minor merger events (the blue line in the middle panel) that have a large inner spin flip is about 6.7%, which is a significant increase on that for total halo flips shown in Fig. | We find that the frequency of minor merger events (the blue line in the middle panel) that have a large inner spin flip is about $6.7\%$, which is a significant increase on that for total halo flips shown in Fig. |
B] (0.7%). | \ref{f:coscdf} $0.7\%$ ). |
The fraction of major merger events that also have significant inner flips is slightly increased, to 26.6%. | The fraction of major merger events that also have significant inner flips is slightly increased, to $26.6\%$. |
Selecting just large flips (right panel), the frequencies are similarly increased compared to the total halo flip distribution: 98.9% of flips of at least 45° coincide with minor mergers, dropping slightly to 97.3% for flips of at least 90°. | Selecting just large flips (right panel), the frequencies are similarly increased compared to the total halo flip distribution: $98.9\%$ of flips of at least $45\degr$ coincide with minor mergers, dropping slightly to $97.3\%$ for flips of at least $90\degr$. |
While it is important to understand the overall frequency of flip events, and their tendency to correlate with mergers, we are also concerned with the frequency of spin orientation changes over the course of halo lifetimes. | While it is important to understand the overall frequency of flip events, and their tendency to correlate with mergers, we are also concerned with the frequency of spin orientation changes over the course of halo lifetimes. |
labels:An important MWhflipfracDtquestion to answer is what is the likelihood of a halo exhibiting a spin flip (of a given magnitude 6) and measured over a timescale τ) at some point during its lifetime? | An important question to answer is what is the likelihood of a halo exhibiting a spin flip (of a given magnitude $\theta_0$ and measured over a timescale $\tau$ ) at some point during its lifetime? |
This can be | This can be |
derive accurate SERs and examine the relationship between galaxy properties and the dust. content in a galaxy. | derive accurate SFRs and examine the relationship between galaxy properties and the dust content in a galaxy. |
SET are related to the luminosity through a linear scale factor. determined with the assumption of a constant stellar initial mass function (IME). | SFRs are related to the luminosity through a linear scale factor, determined with the assumption of a constant stellar initial mass function (IMF). |
These conversion factors will change if the EME is varied. | These conversion factors will change if the IMF is varied. |
One recently. suggested: possibility is an evolving IAIF (Wilkinsetal.20082.b:Ciunawardhbanaetal.inprep.) which would lead. to an evolving SER conversion factor. | One recently suggested possibility is an evolving IMF \citep{Wlk:08a, Wlk:08b, Gun:10} which would lead to an evolving SFR conversion factor. |
Pflamm-Altenburge..Weidner&Ixroupa(20090) proposed an Integrated Galaxy LAL (GLME) that combines the LIMES of all voung-star clusters to form a ealaxv-wide ΕλΠο | \citet{Pfl:09} proposed an Integrated Galaxy IMF (IGIMF) that combines the IMFs of all young-star clusters to form a galaxy-wide IMF. |
, "This was developed to account. for the inconsistencies in current IMES which are based on isolated stellar clusters and then applied on galaxy wide scales. | This was developed to account for the inconsistencies in current IMFs which are based on isolated stellar clusters and then applied on galaxy wide scales. |
One of the inconsistencies that the IGIME accounts for is the discrepaney between FUN. and lla derived. SEIs at [ow SEDBs (Sullivanetal.2000:Lee2009). | One of the inconsistencies that the IGIMF accounts for is the discrepancy between FUV and $\alpha$ derived SFRs at low SFRs \citep{Sul:00, Lee:09}. |
. We compare and contrast SERs corrected. for. dust using dillerent. obscuration correction methods. identifving an optimum approach. and then use that in a preliminary investigation of SER. histories of galaxies. | We compare and contrast SFRs corrected for dust using different obscuration correction methods, identifying an optimum approach, and then use that in a preliminary investigation of SFR histories of galaxies. |
Phe SEIs are compared with theoretical evolutionary paths to. better unclerstaned the star formation history of the galaxies in our sample. | The SFRs are compared with theoretical evolutionary paths to better understand the star formation history of the galaxies in our sample. |
In 52 we present details of the data used. | In $\S$ 2 we present details of the data used. |
ln §3 we esent a prescription for the derivation ancl the application of the obscuration corrections. | In $\S$ 3 we present a prescription for the derivation and the application of the obscuration corrections. |
5&4 compares Ho. FUY and Ou] derived SEIHs. | $\S$ 4 compares $\alpha$, FUV and ] derived SFRs. |
55 examines the evolution of our SLRs » comparing these SEIts to evolutionary synthesis models. | $\S$ 5 examines the evolution of our SFRs by comparing these SFRs to evolutionary synthesis models. |
We also compare the evolutionary. paths of the SEIs with he predictions of the IGIME theory. | We also compare the evolutionary paths of the SFRs with the predictions of the IGIMF theory. |
§6 is the summary and conclusion of our findings. | $\S$ 6 is the summary and conclusion of our findings. |
Phe assumed. cosmological λεοίος are: Πω=7Okkmss |. Oy=0.3 and Q4=0.7. | The assumed cosmological parameters are: $H_{0}=$ $^{-1}$ $^{-1}$, $\Omega_{M}=0.3$ and $\Omega_{\Lambda}=0.7$. |
All magnitudes are in the AB svstem. | All magnitudes are in the AB system. |
We use cata from the Galaxy and. Mass Assembly (GAAALA) survey (Driveretal. 2009).. | We use data from the Galaxy and Mass Assembly (GAMA) survey \citep{Drv:09}. . |
GAALA is a multi-band imaging and spectroscopic survey covering zz 144. square degrees of skv in three 12«4° regions (Robothametal.2009:Baldryetal. 2009). | GAMA is a multi-band imaging and spectroscopic survey covering $\approx$ 144 square degrees of sky in three $12^{\circ} \times 4^{\circ}$ regions \citep{Rob:09, Bld:09}. |
.. Phe spectroscopy comes from the AXOmoega spectrograph (Sharpetal.2006). at the Anglo Australian Telescope CAAT). | The spectroscopy comes from the AAOmega spectrograph \citep{Shp:06} at the Anglo Australian Telescope (AAT). |
The UV data was obtained from the GALEN-C:AALA survey (Seibertetal.inprep. | The UV data was obtained from the GALEX-GAMA survey \citep{Seb:10}. |
).. A signal-to-noise ratio cut o[2.5 was applied to the UV data. | A signal-to-noise ratio cut of 2.5 was applied to the UV data. |
Phere are 111521 GALEN sources within the GAALA regions of which 110443 matching sources were found⋅ using. a matching. radius. of DRM67. | There are 111521 GALEX sources within the GAMA regions of which 110443 matching sources were found using a matching radius of $6^{\prime \prime}$. |
Such a large matching radius is is necessary due to the increase positional uncertainty inherent in the larger GALEN PSE. | Such a large matching radius is is necessary due to the increased positional uncertainty inherent in the larger GALEX PSF. |
SDSS and spectral line data were matehed to GAALA sources . . . ⋅⋅∕∕r ⋠⋠ ⋅ ⊔⊳∖↓⊔⋏∙≟⋜↧⊔↓⋜↧≼⇍↓⊔⊔⋏∙≟↓⋅⋯⊔⊔⊳∖∪⇂−≻⋅⋅≻↓⋅∢⊾⊳∖⊔∐↓⊔⋏∙≟↓⊔⋜↧∐⊔⋜↧↓⊳∖⋜⋯↓↓≻↓∢⊾∪ 412690 objects. | SDSS and spectral line data were matched to GAMA sources using a matching radius of $2.5^{\prime \prime}$ resulting in a final sample of 47269 objects. |
After removing GN (using the prescription of Ixewley. 2006) and sources without emission lines. bare magnitude and redshift data the final sample size containec 31508 ealaxies. | After removing AGN (using the prescription of Kewley, 2006) and sources without emission lines, $r$ -band magnitude and redshift data the final sample size contained 31508 galaxies. |
All remaining galaxies contain both Lla anc 11:7 fluxes. | All remaining galaxies contain both $\alpha$ and $\beta$ fluxes. |
The final sample has à GALEX FUY magnitude range of IAT«mpeg25.3 and a redshift limit of due to the requirement for Ho. being in the observed spectral range. | The final sample has a GALEX FUV magnitude range of $14.7 < m_{FUV} < 25.3$ and a redshift limit of due to the requirement for $\alpha$ being in the observed spectral range. |
Correcting galaxy emission for dust. requires a case specific approach. that is nebular and continuum. emission. [rom ealaxies require two cilferent treatments. | Correcting galaxy emission for dust requires a case specific approach, that is nebular and continuum emission from galaxies require two different treatments. |
At the core of these obscuration corrections are the obscuration curves and the Jalmer decrement. | At the core of these obscuration corrections are the obscuration curves and the Balmer decrement. |
We apply several different obscuration curves to examine their effects on continuum. (UV) and nebular (lla ancl ΟΠ} emission. | We apply several different obscuration curves to examine their effects on continuum (UV) and nebular $\alpha$ and ]) emission. |
Applying obscuration corrections for continuum. emission requires (wo parts. the reddening sullerecl by the stellar continuum (£2V),,,;/| and an applicable obscuration curve (A(A)). | Applying obscuration corrections for continuum emission requires two parts, the reddening suffered by the stellar continuum $E(B-V)_{cont}$ ) and an applicable obscuration curve $k(\lambda)$ ). |
We use the standard form toapplyour dust corrections. | We use the standard form toapplyour dust corrections. |
£; and £, are the intrinsic ancl observed. galaxy Iuminosities. respectively. | $L_{i}$ and $L_{o}$ are the intrinsic and observed galaxy luminosities respectively. |
Lo obtain AA) for the UV. stellar continuum. the Calzetti(2001) obscuration curve and the FDO5 obscuration curves can be applied. | To obtain $k(\lambda)$ for the UV stellar continuum, the \citet{Cal:01} obscuration curve and the FD05 obscuration curves can be applied. |
When using the Calzetti(2001) curve A(X) must be divided by 0.44. | When using the \citet{Cal:01} curve $k(\lambda)$ must be divided by 0.44. |
“Phe 05 curves can also be applied to correct for the obscuration of nebular emission. | The FD05 curves can also be applied to correct for the obscuration of nebular emission. |
The UV luminosities were derived using the UV. Dluxes from the GALEN survey and the redshifts obtained from the GAALA | The UV luminosities were derived using the UV fluxes from the GALEX survey and the redshifts obtained from the GAMA survey. |
SULVOY, LBων ds derived. from the reddening in the ionized gas. LEB μον | $E(B-V)_{cont}$ is derived from the reddening in the ionized gas, $E(B-V)_{gas}$ . |
ECBVou=0ΕΕVyas (Calzetti2001). and where AC) and Αα) are the obseurations of the nebular emission at the Πα ancl 112 wavelengths derived from a AIW obscuration curve (Seaton1979:Carclelliοἱal.1989) or a theoretically modelled. curve (PD05) as in this paper. | $E(B-V)_{cont} = 0.44E(B-V)_{gas}$ \citep{Cal:01} and where $k(H\beta)$ and $k(H\alpha)$ are the obscurations of the nebular emission at the $\alpha$ and $\beta$ wavelengths derived from a MW obscuration curve \citep{Stn:79,Cdl:89} or a theoretically modelled curve (FD05) as in this paper. |
fa. and fgs are the stellar absorption corrected but not cust corrected Lla and 172 [uxes. | $f_{H\alpha}$ and $f_{H\beta}$ are the stellar absorption corrected but not dust corrected $\alpha$ and $\beta$ fluxes. |
foo and fg; were corrected for stellar absorption using the formalism. outlined. in Hopkinsetal.(2003)© stated below. | $f_{H\alpha}$ and $f_{H\beta}$ were corrected for stellar absorption using the formalism outlined in \citet{Hpk:03} stated below. |
where f], and. f; are the observed and stellar absorption corrected Luxes respectively. | where $f_{o}$ and $f_{s}$ are the observed and stellar absorption corrected fluxes respectively. |
EW is the equivalent. width of the line being corrected and £M is the correction for stellar absorption taken to be which has been shown to be a reasonable assumption (Ciunawardhanaetal.in prep. | EW is the equivalent width of the line being corrected and $EW_{c}$ is the correction for stellar absorption taken to be which has been shown to be a reasonable assumption \citep{Gun:10}. . |
).. Assuming a Case-B recombination with a cdensitv of 100ccnmi and atemperature of KIS. the predicted ratio ol fg to fg «sis 2.86 (Osterbrock 1989)... | Assuming a Case-B recombination with a density of $^{-3}$ and atemperature of K, the predicted ratio of $f_{H\alpha}$ to $f_{H\beta}$ is 2.86 \citep{Ost:89}. . |
ALL Balmer decrements below 2.86 were set equal to2.86 as suggested | All Balmer decrements below 2.86 were set equal to2.86 as suggested |
the main conclusions and in Sect. | the main conclusions and in Sect. |
6 we list main results of the paper. | 6 we list main results of the paper. |
For investigating the long-term radio variability of B0605—085 we use University of Michigan Radio Astronomical Observatory (UMRAO) monitoring data at 4.8 GHz. 8 GHz and 14.5 GHz (Aller et al. | For investigating the long-term radio variability of $-$ 085 we use University of Michigan Radio Astronomical Observatory (UMRAO) monitoring data at 4.8 GHz, 8 GHz and 14.5 GHz (Aller et al. |
1999) complemented with the archival data from 408 MHz to 230 GHz. | 1999) complemented with the archival data from 408 MHz to 230 GHz. |
Table | lists observed frequency. time span of observations. and references. | Table \ref{data} lists observed frequency, time span of observations, and references. |
All references in the table refer to the published earlier data. | All references in the table refer to the published earlier data. |
We analyze total flux-density variability of the quasar at ten frequencies obtained at University of Michigan Radio Astronomical Observatory. Metsühhovi Radio Astronomical Observatory. Haystack observatory. Bologna interferometer. Algonquin Radio Observatory. and SEST telescope. | We analyze total flux-density variability of the quasar at ten frequencies obtained at University of Michigan Radio Astronomical Observatory, Metsähhovi Radio Astronomical Observatory, Haystack observatory, Bologna interferometer, Algonquin Radio Observatory, and SEST telescope. |
The UMRAO data at 4.85 GHz. 8 GHz and 14.5 GHz since March 2001 are unpublished. whereas data from all other telescopes are archival and have been published earlier. | The UMRAO data at 4.8 GHz, 8 GHz and 14.5 GHz since March 2001 are unpublished, whereas data from all other telescopes are archival and have been published earlier. |
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