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From the velocity sample we extracted. those objects waving estimates of the metal abuudance. ou the basis of the catalogue by Cavrel de Strobel et al. (
From the velocity sample we extracted those objects having estimates of the metal abundance, on the basis of the catalogue by Cayrel de Strobel et al. (
1997). to which we added 121 stars from Carney ct al. (
1997), to which we added 124 stars from Carney et al. (
1991). for a total of 1511 stars (σπα 0.15).
1994), for a total of 1541 stars $\sigma_\pi$ $\pi$ $\leq$ 0.15).
Cavrels catalogue
Cayrel's catalogue
From the era of the Einstein observatory. it was known through N-rav imaging that
From the era of the ${\it Einstein}$ observatory, it was known through X-ray imaging that
Tay is the duration of emission of the CRB ceutral engine in seconds (so that A/e=Tay).
$\tau_{\rm dur}$ is the duration of emission of the GRB central engine in seconds (so that $\Delta/c=\tau_{\rm dur}$ ).
When Eq.
When Eq.
1 is satisfied the protous will keep accelerating while the neutrous are left behind. (
\ref{decouple} is satisfied the protons will keep accelerating while the neutrons are left behind. (
If Eq.
If Eq.
Lo is not satisfied the protous and neutrons still decouple. but have the same Loreutz factors because the fireball is no longer accelerating.)
\ref{decouple} is not satisfied the protons and neutrons still decouple, but have the same Lorentz factors because the fireball is no longer accelerating.)
When Eq.
When Eq.
[is satisfied aud the ucutrous dvuamically decouple during the acceleration stage of the fireballs evolution. the fraction of initial energy goiug to the ueutrons is and the final Loreutz factor of the protons aud neutrons is given by (See Devishev.Koclirovsky.&Kocharovsky.(1999)... Fuller.Pruet.&Abazajian (2000)..0r. Meszaros(2000). for a detailed derivation of the above equations).
\ref{decouple} is satisfied and the neutrons dynamically decouple during the acceleration stage of the fireball's evolution, the fraction of initial energy going to the neutrons is and the final Lorentz factor of the protons and neutrons is given by (See \cite{derishev}, \cite{fuller}, ,or \cite{bahcall} for a detailed derivation of the above equations).
Iu this equation 2,,5 aud 2,3 are. respectively. the final Loreutz factors of the proton aud neutron shells in uuits of 100,
In this equation $\gamma_{p,3}$ and $\gamma_{n,3}$ are, respectively, the final Lorentz factors of the proton and neutron shells in units of $10^3$.
Of course. if the ueutrons initially prescut iu the fireball are going to shock aud lead to an observable photon signature they iust first decay iuto protous.
Of course, if the neutrons initially present in the fireball are going to shock and lead to an observable photon signature they must first decay into protons.
Even after they have decayed we will still refer to the slower iuner shell as the “neutron” shell.
Even after they have decayed we will still refer to the slower inner shell as the “neutron” shell.
We note that it was argued in Fuller.Pruct.&Abazajian(2000) that in fireballs with very low initial 35. the electron fraction will be driven to 3;~0.05 during neutron decoupling.
We note that it was argued in \cite{fuller} that in fireballs with very low initial $Y_e$, the electron fraction will be driven to $Y_e\sim 0.05$ during neutron decoupling.
A useful relation. derivable from Eqs.
A useful relation, derivable from Eqs.
57 and δν, aud valid when ueutrou decoupling occurs. is 3)
\ref{fn} and \ref{gammap}, and valid when neutron decoupling occurs, is $\gamma_{p,3}^{4/3}=(1/5)(1-f_n)^{1/3}(\gamma_p/\gamma_n) (E_{51}/r_6\tau_{\rm dur})^{1/3}$ .
Using this relation and Eq. 1..
Using this relation and Eq. \ref{xi},
note that thore is a very interesting connection between the couditious in the neutron decoupled fireball aud the iuitial value of the parameter &. This implics that for a given £. neutron decoupling occurs iu the progenitor fireball unless ©.
note that there is a very interesting connection between the conditions in the neutron decoupled fireball and the initial value of the parameter $\xi$, This implies that for a given $\xi$, neutron decoupling occurs in the progenitor fireball unless $n_{\rm ISM} E_{51} \tau_{\rm dur} Y_e/r_6^2>(3)^6 \xi^{-6}$ .
Therefore. We are interested im the couditious wader which a slow ueutrou shell decavs aud shocks with the outer proton shell.
Therefore, We are interested in the conditions under which a slow neutron shell decays and shocks with the outer proton shell.
Di order for this to occur. aud in order for observable enissiou to result. the following couditious iust be moet: 1) neutron decoupling occurs aud leads to a final ratio of proton to neutron Loreutz factors of 5,73>6s
In order for this to occur, and in order for observable emission to result, the following conditions must be met: i) neutron decoupling occurs and leads to a final ratio of proton to neutron Lorentz factors of $\gamma_{p,3}/\gamma_{n,3}>\alpha$.
Tere a deteruunes the strength of neutron decoupling aud is chosen so that the emission from the two peaks is distiuguislhable.
Here $\alpha$ determines the strength of neutron decoupling and is chosen so that the emission from the two peaks is distinguishable.
a) the fraction of energy. eoine to the neutrons is non-uceleible (for definiteness we nupose f,2 0.2). aud ii) the neutrons decay before colliding aud shocking with the outer decelerating shell.
ii) the fraction of energy going to the neutrons is non-negligible (for definiteness we impose $f_n>0.2$ ), and iii) the neutrons decay before colliding and shocking with the outer decelerating shell.
We denote with a subscript p properties of the observed euission from the (faster) proton shell (so that T5, is the observed curation of the first peak). aud with a subscript » properties of the observed enission from the decaved neutron shell.
We denote with a subscript $p$ properties of the observed emission from the (faster) proton shell (so that $T_{b,p}$ is the observed duration of the first peak), and with a subscript $n$ properties of the observed emission from the decayed neutron shell.
When &71. the expression for the duration of the euission frou the proton shell can be written in the useful form Eq.
When $\xi>1$, the expression for the duration of the emission from the proton shell can be written in the useful form Eq.
8 allows us to relate the observed proton shell burst duration to the condition that the neutrous strouglv decouple aud carry a substantial fraction of the fireball energy. (conditions (1) aud Gi) above).
\ref{usefultbp} allows us to relate the observed proton shell burst duration to the condition that the neutrons strongly decouple and carry a substantial fraction of the fireball energy (conditions (i) and (ii) above).
Noting that f£,=Vaal(spiteLV) we see that 2,75,>a and f£,>0.2 whe
Noting that $f_n=(1-Y_e)/(1+(\gamma_p/\gamma_n-1)Y_e)$ we see that $\gamma_p/\gamma_n>\alpha$ and $f_n>0.2$ when
other determinations of AY/AZ available in the literature.
other determinations of $\Delta Y/\Delta Z$ available in the literature.
This work is partly supported by NSF grant ATM-0348837 to SD.
This work is partly supported by NSF grant ATM-0348837 to SB.
fields are produced al regular intervals during operation in all fillers of each camera.
fields are produced at regular intervals during operation in all filters of each camera.
We used FLIOW and FIGOW camera 3 flat field observations created on Sept. 9. 2003 [rom proposal 9640.
We used F110W and F160W camera 3 flat field observations created on Sept. 9, 2003 from proposal 9640.
The flat fields are analvzed in the identical manner as described in the preceding steps.
The flat fields are analyzed in the identical manner as described in the preceding steps.
The ΤΙ reference flat fields were not used in this analvsis because tliev were based on [lat fields observed previous to the NCS safing event.
The STScI reference flat fields were not used in this analysis because they were based on flat fields observed previous to the NCS safing event.
One of the effects that the flat field corrects is a slight vignetting along the lower edge ol camera 3.
One of the effects that the flat field corrects is a slight vignetting along the lower edge of camera 3.
For (wo reasons this correction was not ellective in the UDF fields.
For two reasons this correction was not effective in the UDF fields.
First. the net ellect has (wo components. vignetting of the incoming astronomical flux. ancl emission from the vignetüng component which is thought to be the edge of the mount for the field division mirror for camera 3.
First, the net effect has two components, vignetting of the incoming astronomical flux and emission from the vignetting component which is thought to be the edge of the mount for the field division mirror for camera 3.
For bright sources the vignetting is the dominant effect and the flat field properly corrects the field.
For bright sources the vignetting is the dominant effect and the flat field properly corrects the field.
For very faint images. such as the UDF. emission can be a significant component which varies due to the natural temperature variations in (he alt shroud.
For very faint images, such as the UDF, emission can be a significant component which varies due to the natural temperature variations in the aft shroud.
Second. variations in geometry due to temperature changes can affect the degree of vienet(ing.
Second, variations in geometry due to temperature changes can affect the degree of vignetting.
Again the effect is slight for bright sources but can be significant for the UDF signal levels.
Again the effect is slight for bright sources but can be significant for the UDF signal levels.
For these reasons the lower 20 rows of all UDF images were masked off in the drizzle procedure described in section 3.7..
For these reasons the lower 20 rows of all UDF images were masked off in the drizzle procedure described in section \ref{ss-dp}.
Bad pixels are defined as pixels with quantum efficiencies less than 1054 of the average QE or with dark currents high enough (o reach nonlinear signal levels in 1000 seconds or less.
Bad pixels are defined as pixels with quantum efficiencies less than $10\%$ of the average QE or with dark currents high enough to reach nonlinear signal levels in 1000 seconds or less.
In the post NCS era the list of bad. pixels has increased over the evele 7 listing due to the higher temperature of the detector creating more high dark current pixels.
In the post NCS era the list of bad pixels has increased over the cycle 7 listing due to the higher temperature of the detector creating more high dark current pixels.
Pixels which satis[v neither criterion but are highly variable in their dark eiurrent were also added to the list.
Pixels which satisfy neither criterion but are highly variable in their dark current were also added to the list.
All bad pixel signals are replaced with the median of the eight pixels surrounding them.
All bad pixel signals are replaced with the median of the eight pixels surrounding them.
In the case of adjacent bad. pixels (hat number is reduced by the number of adjoining bad pixels.
In the case of adjacent bad pixels that number is reduced by the number of adjoining bad pixels.
All bad. pixels are listed in the data quality array which is an extension of the image or SCI array.
All bad pixels are listed in the data quality array which is an extension of the image or SCI array.
Table 3. exgives the decimal codes for each of the steps described above.
Table \ref{tb-dqc} gives the decimal codes for each of the steps described above.
They are each a single. different. binary. bit. so each combination of actions performed on a pixel has a unique output code.
They are each a single, different, binary bit, so each combination of actions performed on a pixel has a unique output code.
Note that only a few of the 16 bits available lor pixel actions are used in (his analvsis.
Note that only a few of the 16 bits available for pixel actions are used in this analysis.
A full set of data quality codes can be found in the NICMOS Handbook but only the ones listed here are used in the NICMOS UDF Treasury data.
A full set of data quality codes can be found in the NICMOS Handbook but only the ones listed here are used in the NICMOS UDF Treasury data.
Even for the codes used they may in many cases differ from the codes returned by the ST5cl pipeline analysis.
Even for the codes used they may in many cases differ from the codes returned by the STScI pipeline analysis.
As an example. the lists of bad. pixels differ between the pipeline analvsis and the analysis described here.
As an example, the lists of bad pixels differ between the pipeline analysis and the analysis described here.
Note that the data equality extensions only exist for the individual NICMOS UDF images.
Note that the data quality extensions only exist for the individual NICMOS UDF images.
The drizzle procedure does not preserve these codes since many input. pixels
The drizzle procedure does not preserve these codes since many input pixels
11996: Lewis et 12998b).
1996; Lewis et 1998b).
This situation would arise very naturally if the hotter dust. cloucs were smaller and. more central. as in the AGN moclel deseribed in equation (2).
This situation would arise very naturally if the hotter dust clouds were smaller and more central, as in the AGN model described in equation (2).
A similar cHeet can be produced. by the microlensing cllect of individual stars within the lensing galaxy (Lewis ct L199Sa).
A similar effect can be produced by the microlensing effect of individual stars within the lensing galaxy (Lewis et 1998a).
lor a large magnilication to occur. a clistant source must lie very close to à caustic curve of a geavitational lens.
For a large magnification to occur, a distant source must lie very close to a caustic curve of a gravitational lens.
The magnification is formally infinite on such a curve. but an upper limit μμ is imposed to the magnification if the source has a finite size d (Peacock 1982).
The magnification is formally infinite on such a curve, but an upper limit $A_{\rm max}$ is imposed to the magnification if the source has a finite size $d$ (Peacock 1982).
Ifthe lensing galaxy can be modeled as a singular isothermal sphere (SIS). then thinsxd+.
If the lensing galaxy can be modeled as a singular isothermal sphere (SIS), then $A_{\rm max} \propto d^{-1}$.
Assuming an SIS lens. the image geometry and magnifications expected in such a situation are illustrated in Kies 2 ane 3.
Assuming an SIS lens, the image geometry and magnifications expected in such a situation are illustrated in Figs 2 and 3.
Eisenhardt et. al. (
Eisenhardt et al. (
1996) present more sophisticated lens models that account for the geometry of high-resolution images of FE10214|4724.
1996) present more sophisticated lens models that account for the geometry of high-resolution images of F10214+4724.
Ellipticity in the lens moclifics the magnificationsize relation to chinsoxdο on scales between 0.001. and Laaresecc.
Ellipticity in the lens modifies the magnification–size relation to $A_{\rm max} \propto d^{-0.63}$ on scales between 0.001 and arcsec.
In general. a similar magnificationsize relationship holds regardless of both the gcometry of the source ancl whether one or more objects is responsible for producing the lensing ellect. (I&neib et 11998).
In general, a similar magnification–size relationship holds regardless of both the geometry of the source and whether one or more objects is responsible for producing the lensing effect (Kneib et 1998).
The diagnostic feature of such à situation is the production of multiple images of comparable brightness.
The diagnostic feature of such a situation is the production of multiple images of comparable brightness.
"This is clearly the case for both APALOOS279|5255 (Lewis et 11998b) and 11413| (Να et
This is clearly the case for both 08279+5255 (Lewis et 1998b) and H1413+117 (Kneib et
sphericall density àistributions suggested by cosmological simulations.
l density distributions suggested by cosmological simulations.
the binary orbit of RRo where cooler plasma resides.
the binary orbit of $_\odot$ where cooler plasma resides.
Hotter plasma further inside may have been blocked during the times of the flares, as both occurred during Closer inspection of the spectral changes with the dips in the first observation gives insights into the geometry of the emitting gas at this time.
Hotter plasma further inside may have been blocked during the times of the flares, as both occurred during Closer inspection of the spectral changes with the dips in the first observation gives insights into the geometry of the emitting gas at this time.
The brightness map in the top panel of Fig.
The brightness map in the top panel of Fig.
12 illustrates that during the dips, the continuum and the emission lines are both fainter, but neither component is completely occulted.
\ref{smap} illustrates that during the dips, the continuum and the emission lines are both fainter, but neither component is completely occulted.
Both components must therefore originate from both, the inner orbit and from further outside.
Both components must therefore originate from both, the inner orbit and from further outside.
Thompson scattering of the continuum component has already been discussed in Sect.
Thompson scattering of the continuum component has already been discussed in Sect.
4.1.0 to resolve the conflict between high temperature from the Wien tail and the required small size of the central source plus the absence of absorption lines.
\ref{rgsspec} to resolve the conflict between high temperature from the Wien tail and the required small size of the central source plus the absence of absorption lines.
A comparison of the emission lines during and outside dips on day 22.9 can be studied in Fig. 14..
A comparison of the emission lines during and outside dips on day 22.9 can be studied in Fig. \ref{diffspec}.
The difference spectrum represents the plasma within the inner orbit, revealing the spectral signatures of resonance line scattering.
The difference spectrum represents the plasma within the inner orbit, revealing the spectral signatures of resonance line scattering.
While the strong 1s-2p resonance lines of, e.g., AA)) and AA)) originate only from the outer regions, lines belonging to transitions involving higher principal quantum numbers are stronger in the difference spectrum, e.g., the He-like series lines Hef, 7, and ó.
While the strong 1s-2p resonance lines of, e.g., ) and ) originate only from the outer regions, lines belonging to transitions involving higher principal quantum numbers are stronger in the difference spectrum, e.g., the He-like series lines $\beta$, $\gamma$, and $\delta$ .
The respective oscillator strengths are f=0.17, 0.05, and 0.02, which are low compared to the oscillator strength of the Hea line of f=0.71.
The respective oscillator strengths are $f=0.17$, $0.05$, and $0.02$, which are low compared to the oscillator strength of the $\alpha$ line of $f=0.71$.
Line photons with high values of f can be absorbed and re-emitted in a different direction.
Line photons with high values of $f$ can be absorbed and re-emitted in a different direction.
In a non-spherical plasma, strong resonance lines can be amplified or reduced, compared to lines of lower oscillator strengths.
In a non-spherical plasma, strong resonance lines can be amplified or reduced, compared to lines of lower oscillator strengths.
The fact that we are seeing the high-f lines from the inner regions reduced is consistent with the geometry of an accretion disk.
The fact that we are seeing the $f$ lines from the inner regions reduced is consistent with the geometry of an accretion disk.
Forbidden transitions can not be absorbed, and the contributions from the inner region to the ratio of the intercombination and forbidden lines at 29.1 and iis unaltered by resonance scattering.
Forbidden transitions can not be absorbed, and the contributions from the inner region to the ratio of the intercombination and forbidden lines at 29.1 and is unaltered by resonance scattering.
Owing to the proximity to the central source, the low f/i ratio reported in Sect.
Owing to the proximity to the central source, the low f/i ratio reported in Sect.
4.1.0 has to be interpreted as photoexcitation rather than high density.
\ref{rgsspec} has to be interpreted as photoexcitation rather than high density.
In the top panel of Fig.13, the evolution of the Wien tail can be studied that is a qualitative indicator for the photospheric temperature.
In the top panel of \ref{smap_epic}, the evolution of the Wien tail can be studied that is a qualitative indicator for the photospheric temperature.
During dips, the continuum emission has gone down at all wavelength, and no shift of the Wien tail can be seen that would indicate that the continuum component from the outside regions has a different temperature.
During dips, the continuum emission has gone down at all wavelength, and no shift of the Wien tail can be seen that would indicate that the continuum component from the outside regions has a different temperature.
This supports the interpretation of Thompson scattering.
This supports the interpretation of Thompson scattering.
The sudden change in the variability pattern at phase 1.25 is accompanied by a sudden increase in brightness above kkeV, extending to ~1.5 kkeV, including the ls-2p transition at kkeV. While solar flare activity could be held responsible for the brighter emission between phases 1.2-1.35, the spectra taken after phase 1.35 are not contaminated and are still harder.
The sudden change in the variability pattern at phase 1.25 is accompanied by a sudden increase in brightness above keV, extending to $\sim 1.5$ keV, including the 1s-2p transition at keV. While solar flare activity could be held responsible for the brighter emission between phases 1.2-1.35, the spectra taken after phase 1.35 are not contaminated and are still harder.
This could be interpreted in the general context of increasing temperature on longer time scales as discussed above.
This could be interpreted in the general context of increasing temperature on longer time scales as discussed above.
Also, a slight increase in the high-ionization line is suggestive in the top panel of Fig. 12,,
Also, a slight increase in the high-ionization line is suggestive in the top panel of Fig. \ref{smap},
but itO may also be part of the general increase in 'The spectral evolution during the second observation is displayed in the bottom panels of Figs.
but it may also be part of the general increase in The spectral evolution during the second observation is displayed in the bottom panels of Figs.
12 and 13..
\ref{smap} and \ref{smap_epic}.
The continuum and the emission line components experience similar changes during the eclipse as during the dips on day 22.9.
The continuum and the emission line components experience similar changes during the eclipse as during the dips on day 22.9.
The total emission decreases by ~50% during eclipse, indicative of significant scattering processes.
The total emission decreases by $\sim 50$ during eclipse, indicative of significant scattering processes.
The emission lines decrease by roughly the same amount as the continuum.
The emission lines decrease by roughly the same amount as the continuum.
This is illustrated with the difference spectra shown with blue color in Fig. 14..
This is illustrated with the difference spectra shown with blue color in Fig. \ref{diffspec}.