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Piled-up rate for the NOS2 camera is —S counts D (0.510 keV). | Piled-up rate for the MOS2 camera is $\sim 8$ counts $^{-1}$ (0.5–10 keV). |
Source position is consistent with that determined with Chandra (Alarkwardt. Juda Swank 2003). | Source position is consistent with that determined with Chandra (Markwardt, Juda Swank 2003). |
A simple way round to pile-up is to extract photons for spectral analysis in au annulus around the source. thus excluding the imuer (piled-up) core. | A simple way round to pile-up is to extract photons for spectral analysis in an annulus around the source, thus excluding the inner (piled-up) core. |
We followed this strategv for analvsine the MOS2 data. whereas for AIOSI data the window is foo snall(see also the analysis carried out on Mni 121. presenting simular problems. bv Ravasio et al. | We followed this strategy for analysing the MOS2 data, whereas for MOS1 data the window is too small (see also the analysis carried out on Mkn 421, presenting similar problems, by Ravasio et al. |
2003. in preparation). | 2003, in preparation). |
We extracted photous from an anmuulus of 30% and 60" immer and outer radii. respectively. | We extracted photons from an annulus of $30''$ and $60''$ inner and outer radii, respectively. |
Background was extracted from two different outer circular regions 2 cach. | Background was extracted from two different outer circular regions $2'$ each. |
We selected only siugle-pixel events and make sure that uo pile-up is prescut (through the tool epatplot). | We selected only single-pixel events and make sure that no pile-up is present (through the tool ). |
This is also in line with prescriptions by Moleudi Seiibay. | This is also in line with prescriptions by Molendi Sembay. |
The count rate within the 0.510 keV cnerev baud is 1.01+0.01 counts 1. | The count rate within the 0.5–10 keV energy band is $1.01\pm0.01$ counts $^{-1}$. |
Respouse matrix aud ancillary files were then ecucrated with audarfgen. | Response matrix and ancillary files were then generated with and. |
Spectral data were rebiuned to have SO counts per bin. | Spectral data were rebinned to have 80 counts per bin. |
Timing mode was selected to have the opportunity to reveal Xrav pulsations. | Timing mode was selected to have the opportunity to reveal X–ray pulsations. |
(ναι the high count rate of the source. this mode is particularly useful to avoid any pile-up. | Given the high count rate of the source, this mode is particularly useful to avoid any pile-up. |
We extracted the source spectrum from the raw μαnaee. taking a column 10-pixel wide aud thebackerouu from two separate boxes 5-pixels wide. | We extracted the source spectrum from the raw image, taking a column 10-pixel wide and the background from two separate boxes 5-pixels wide. |
We cousider only ssingle” ando “double” events. | We consider only “single” and “double” events. |
No OXray bursts were etected. | No X–ray bursts were detected. |
The source count rate in the 0.510 keV. hane is 310280.1 counts s1, | The source count rate in the 0.5–10 keV band is $34.0\pm0.1$ counts $^{-1}$. |
Background is at a level of 1.0 oe1i the full baud. | Background is at a level of $1.4\%$ in the full band. |
Pile-up is ucelicible aud dead time is a a level of ~1.5%. | Pile-up is negligible and dead time is at a level of $\sim 1.5\%$. |
Response matrix and ancillary file for PApectral analysis were generated using tasks. | Response matrix and ancillary file for spectral analysis were generated using tasks. |
Spectra ata were rebiuned to have SO counts per bin. | Spectral data were rebinned to have 80 counts per bin. |
The standard procedure was used to derive the first order RCS spectra. | The standard procedure was used to derive the first order RGS spectra. |
In the 736 bbaud rates are 0.39+0.01 aud 0.5640.01 count | for RGS1 aud RGS2. respectively. | In the 7–36 band rates are $0.39\pm0.01$ and $0.56\pm0.01$ count $^{-1}$ for RGS1 and RGS2, respectively. |
For spectral analvsis together with the other iustruuieuts we rebinned the data to LOO counts per bin. | For spectral analysis together with the other instruments we rebinned the data to 100 counts per bin. |
We further inspect the RCS spectra for euission or absorption lues. | We further inspect the RGS spectra for emission or absorption lines. |
Iu this case we rebin the spectra to 20 counts per bin (see below). | In this case we rebin the spectra to 20 counts per bin (see below). |
The spectral analysis has been carried out in the 0.510 τον enerev range for MOS2 and pu aud in 736 for the two RGSs. using NSPEC (v11.2). | The spectral analysis has been carried out in the 0.5–10 keV energy range for MOS2 and pn and in 7–36 for the two RGSs, using XSPEC (v11.2). |
All spectral uncertainties are given at 90% confidence level for oue deerce of freedom (A4?=2.71). | All spectral uncertainties are given at $90\%$ confidence level for one degree of freedom $\Delta \chi^2=2.71$ ). |
We first tried he standard model for SNRTs. ie. au absorbed soft component (black body) plus a hard compoucut (power aw: ce. Campana ct al. | We first tried the standard model for SXRTs, i.e. an absorbed soft component (black body) plus a hard component (power law; e.g. Campana et al. |
1998b). | 1998b). |
The fit is good with a ra=1.10 (1571 degrees of freedom. d.o.f.). | The fit is good with a $\chi^2_{\rm red}=1.10$ (1571 degrees of freedom, d.o.f.). |
Fit xuanneters can be found in Table 1.. | Fit parameters can be found in Table \ref{spetab}. |
The column density is (6.320.1)<1073cus? which is slightly higher than he Galactic value of 3:«1074cur7. | The column density is $(6.3\pm0.1)\times 10^{21}\cmdue$ which is slightly higher than the Galactic value of $3\times 10^{21}\cmdue$. |
A comparable (but nore plivsical) model is obtained by replacing the power aw with a Comptouization model Tike (Titarchu- un as observed in SAX. J1808.13658 (Goertinshi. Don miBarret 2002 but not iu NTE 1751 Aller et al. | A comparable (but more physical) model is obtained by replacing the power law with a Comptonization model like (Titarchuk 1994) as observed in SAX J1808.4–3658 (Gierlinski, Done Barret 2002 but not in XTE 1751–305, Miller et al. |
Tu | 2003). |
this case we obtain rag=1.10 (1569 d.o.f.. | In this case we obtain $\chi^2_{\rm
red}=1.10$ (1569 d.o.f., |
see Fig. D). | see Fig. \ref{spe}) ). |
Svstematic errors at 1.6% level provide a- acceptable fit. | Systematic errors at $1.6\%$ level provide an acceptable fit. |
The column density is LG!«1023un7. cousistent with the Calactic value. | The column density is $4.6^{+0.3}_{-0.2}\times 10^{21}\cmdue$, consistent with the Galactic value. |
The black. body temperature is ALF=Os0d0.03 keV aud its radius (assuming spherical sviuuetiv) is Ro=(2.2+0.2)f? kan (with ds the unknown distance to NTE JLso7291 iu units of 8 aud f the spectral hardening factor obtained as the ratio of the color temperature to the effective telmpecrature. e.g. Merloni. Fabian Ross 2000). | The black body temperature is $k\,T=0.80\pm0.03$ keV and its radius (assuming spherical symmetry) is $R=(2.2\pm0.2)\,d_{8}^2\,f^2$ km (with $d_{8}$ the unknown distance to XTE J1807–294 in units of 8 and $f$ the spectral hardening factor obtained as the ratio of the color temperature to the effective temperature, e.g. Merloni, Fabian Ross 2000). |
The temperature of the soft input photon is 0.21c0.02 seV. Tn the Comptonization model the temperature of the corona and the plasma optical depth are tightly related. | The temperature of the soft input photon is $0.24\pm0.02$ keV. In the Comptonization model the temperature of the corona and the plasma optical depth are tightly related. |
The best fit is for a relatively cool corona (&T.=9.6!nia seV. Lo. KT>5.6 keV) aud an optical depth (7<12). | The best fit is for a relatively cool corona $k\,T=9.6^{+102}_{-4.0}$ keV, i.e. $k\,T>5.6$ keV) and an optical depth $\tau<4.2$ ). |
A very hot corona (kT=>370 keV) with a negligible optical depth (7~0.02. zr«2.68 at 90% confidence evel. cl.) | A very hot corona $k\,T>370$ keV) with a negligible optical depth $\tau\sim
0.02$, $\tau<2.68$ at $90\%$ confidence level, c.l.) |
is also (formally) acceptable. | is also (formally) acceptable. |
Towever. with a detector sensitive up to 10 keV. it is impossible to test he presence of such a high coronal temperature. | However, with a detector sensitive up to 10 keV, it is impossible to test the presence of such a high coronal temperature. |
The unabsorbed 0.510 keV flux for our best fit spectrmi aumouuts to 2.6«10.erestem7. | The unabsorbed 0.5–10 keV flux for our best fit spectrum amounts to $2.6\times 10^{-10}\ergs\cmdue$. |
Source luminosity is 2.0«10742ereἩ, | Source luminosity is $2.0\times 10^{36}\,d_{8}^2\ergs$. |
The extrapolated Imminosity to the 0.01.100 keV band is a factor ~1.7 larecr. | The extrapolated luminosity to the 0.01–100 keV band is a factor $\sim 1.7$ larger. |
The hard component contributes almost the ecutire (Quunabsorbed) flux (87). | The hard component contributes almost the entire (unabsorbed) flux $87\%$ ). |
The black body. component is detectable ouly at low energies | The black body component is detectable only at low energies. |
Other soft commpoucuts have been tried with similarly good results. | Other soft components have been tried with similarly good results. |
We caunot discriminate between an enmusson from a black body or froiir an accretion disk (Ty~0.2 keV. Ry~10 lan). | We cannot discriminate between an emission from a black body or from an accretion disk $T_{d}\sim0.2$ keV, $R_{d}\sim 10$ km). |
A neutron star atmosphere (Cuinusicke. Braje Romani 2002) aud a thermal spectrum however provided a worse fit to the data. | A neutron star atmosphere (Gännsicke, Braje Romani 2002) and a thermal spectrum however provided a worse fit to the data. |
We searched for an iron lines in the pu aud MOS data with negative results. | We searched for an iron lines in the pn and MOS data with negative results. |
Takine the line energy in 6.10.6.97 keV and 0.1 keV (best fit) or null values for the line width. we derive au upper nit on the equivalent width of 18 aud 25 eV. respectively. | Taking the line energy in 6.40–6.97 keV and 0.1 keV (best fit) or null values for the line width, we derive an upper limit on the equivalent width of 18 and 25 eV, respectively. |
Following Miller et al. ( | Following Miller et al. ( |
2002) we searched the ROS spectra for cluission or absorption lines. | 2002) we searched the RGS spectra for emission or absorption lines. |
We divided the spectra (rebiuned at 20 counts per bin) in slices of 3 aand inspect them for lines in the enerev band 720 Α.. | We divided the spectra (rebinned at 20 counts per bin) in slices of 3 and inspect them for lines in the energy band 7–20 . |
and a J star (Group 4) with a strong lithium line. | and a J star (Group 4) with a strong lithium line. |
The relative smoothness of the residual spectra. bevond the bounds of the lithium line and seen throughout the entire upper panel shows that this is à valid procedure which takes account of the strong star-to-star variations in the overall strengths of the €N bands. presumably depending on both temperature and carbon and nitrogen abundances. | The relative smoothness of the residual spectra beyond the bounds of the lithium line and seen throughout the entire upper panel shows that this is a valid procedure which takes account of the strong star-to-star variations in the overall strengths of the CN bands, presumably depending on both temperature and carbon and nitrogen abundances. |
ltesidual spectra such as those shown as solid. lines in Fie. | Residual spectra such as those shown as solid lines in Fig. |
2 were constructed For every observed star. | \ref{montage}
were constructed for every observed star. |
Lines in the residual spectra were then sought as follows: a second-order »»Ivnomial was fitted. to the residual spectrum: using the east-squares method in order to define the "continuum and all maxima ancl minima (potential emission and absorption ines respectively) about this fit were identified. | Lines in the residual spectra were then sought as follows: a second-order polynomial was fitted to the residual spectrum using the least-squares method in order to define the `continuum' and all maxima and minima (potential emission and absorption lines respectively) about this fit were identified. |
A gaussian. superimposed on the quacwatie fit. was then fitted to each maximum and minimum allowing its central wavelength (A). height and full width at half maximum. (EWIIM) to vary. | A gaussian, superimposed on the quadratic fit, was then fitted to each maximum and minimum allowing its central wavelength $\lambda_{cen}$ ), height and full width at half maximum ) to vary. |
The equivalent widths of the identified (both emission ancl absorption) lines (Wy) were then measured. | The equivalent widths of the identified (both emission and absorption) lines $_{\lambda}$ ) were then measured. |
I should be emphasized that these equivalent widths should be used in a relative sense only. due to their strong dependence on the way the continuum is defined. | It should be emphasized that these equivalent widths should be used in a relative sense only, due to their strong dependence on the way the continuum is defined. |
Other authors use a pseucdo-continuum. so equivalent widths from cilferent sources are not directly comparable. | Other authors use a pseudo-continuum, so equivalent widths from different sources are not directly comparable. |
Clearly. a spectrum with poor signal-to-noise ratio will eencrate random "lines. | Clearly, a spectrum with poor signal-to-noise ratio will generate random `lines'. |
H turns out that the number of detected lines starts to rise significantly once the count level falls below ~500 counts per pixel. | It turns out that the number of detected lines starts to rise significantly once the count level falls below $\sim$ 500 counts per pixel. |
Consequently. 53. stars with fewer than 500 counts per pixel were excluded. from further consideration. | Consequently, 53 stars with fewer than 500 counts per pixel were excluded from further consideration. |
The derived Wy are plotted. againsto A; in Fig.e 3.. | The derived $_{\lambda}$ are plotted against $\lambda_{cen}$ in Fig. \ref{lines}. |
What appears as a gap around Wy=Q0 is merely a consequence of the limit applied. to the line search procedure. | What appears as a gap around $W_{\lambda}=0$ is merely a consequence of the limit applied to the line search procedure. |
Absorption features appear in the lower part of the figure: emission. features appear in the upper half. | Absorption features appear in the lower part of the figure; emission features appear in the upper half. |
The fairly uniform distribution of points with [WA]< | The fairly uniform distribution of points with $|W_{\lambda}| \leq$ |
8M (CCSNe). 1978;Falk1978) (Soderbergetal.2008;al.2009),, Gezarietal.2008), 2008).. (Ims | $M_\odot$ \citep{kc78,f78} \citep{s08,l08,m08,m09}, \citep{sc08,g08}, \citep{g08}. |
hennik&Nadezhin1988;Matzner2010a) (Shigeyamaetal.1988;WoosleyBlinnikov2009). | \citep{in88,mm99,ss10a} \citep{s88,w88,bn91,eb92,to09}. |
. Shigeyamaetal.(1988) Blinnikov&Nadyozhin Ensman&Burrows(1992) the following general agreements: i) the temperature behind the shock wave reaches about 10? - 109 K, ii) the matter behind the shock wave emits UV/soft X-ray photons having a black body spectrum, iii) the light travel time across the progenitor star is crucial to determine the light curves of the emissions. | \cite{s88} \cite{bn91} \cite{eb92} the following general agreements: i) the temperature behind the shock wave reaches about $10^5$ - $10^6$ K, ii) the matter behind the shock wave emits UV/soft X-ray photons having a black body spectrum, iii) the light travel time across the progenitor star is crucial to determine the light curves of the emissions. |
It is noteworthy that the observed spectrum of the shock breakout of SN 2008D is well fitted by a power-law (Soderbergetal.2008;LiMazzaliModjazetal. 2009), which contradicts the theoretical expectation. | It is noteworthy that the observed spectrum of the shock breakout of SN 2008D is well fitted by a power-law \citep{s08,l08,m08,m09}, which contradicts the theoretical expectation. |
Recent investigations etal.2007;Suzuki&Shigeyama suggest(Wang that interactions between thermal photons2010a) and electrons at the shock front via Compton scattering, which is so-called bulk comptonization &Payne1981a,b),, can form a power-law spectrum. | Recent investigations \citep{wwm07,ss10a} suggest that interactions between thermal photons and electrons at the shock front via Compton scattering, which is so-called bulk comptonization \citep{bp81a,bp81b}, can form a power-law spectrum. |
(Blandford These investigations assume spherical symmetry, i.e., the shock breakout is assumed to occur simultaneously at every point on the surface of the progenitor star. | These investigations assume spherical symmetry, i.e., the shock breakout is assumed to occur simultaneously at every point on the surface of the progenitor star. |
However, recent studies on the explosion mechanism of CCSNe strongly indicate aspherical deposition of the explosion energy (MacFadyen&Woosley1999;Blondinetal.2003;Kotake although details of the energy deposition process 2004),,remain unclear. | However, recent studies on the explosion mechanism of CCSNe strongly indicate aspherical deposition of the explosion energy \citep{m99,b03,k04}, although details of the energy deposition process remain unclear. |
From the observational viewpoint, CCSNe are considered to be aspherical in general (e.g.,Maedaetal. | From the observational viewpoint, CCSNe are considered to be aspherical in general \citep[e.g.,][]{ma08}. |
For an aspherical explosion, the shock breakout 2008)..does not occur simultaneously in different radial directions. | For an aspherical explosion, the shock breakout does not occur simultaneously in different radial directions. |
As a result, the light curve during the shock breakout of an aspherical explosion can deviate from those of spherical one. | As a result, the light curve during the shock breakout of an aspherical explosion can deviate from those of spherical one. |
In other words, the light curve during the shock breakout can be used as a probe for the explosion geometry of a CCSN. | In other words, the light curve during the shock breakout can be used as a probe for the explosion geometry of a CCSN. |
Couchetal. performed series of hydrodynamical simulations of (2009)jet-induced explosionsa of a red supergiant progenitor. | \cite{c09} performed a series of hydrodynamical simulations of jet-induced explosions of a red supergiant progenitor. |
They investigated only the difference of light curves during the shock breakout between thermal and kinetic energy-dominated jets. | They investigated only the difference of light curves during the shock breakout between thermal and kinetic energy-dominated jets. |
However, light curves during the shock breakout must contain more information on the explosion geometry of a CCSN, e.g., the viewingangle,the degree of aspherical energy deposition, and so on. | However, light curves during the shock breakout must contain more information on the explosion geometry of a CCSN, e.g., the viewingangle,the degree of aspherical energy deposition, and so on. |
In this letter, we present an approximate method to calculate light curves during the shock breakout assuming axisymmetric energy deposition and show | In this letter, we present an approximate method to calculate light curves during the shock breakout assuming axisymmetric energy deposition and show |
numerical simulations suggest that warm gas in halos extends out to galactocentric distances of ~150 kpe with cloud covering fractions of —0.250.6 (Maller&Bullock2004:Kaufmannetal. 2006). | \citep[e.g.,][]{mo02} \citep[e.g.,][]{white78} numerical simulations suggest that warm gas in halos extends out to galactocentric distances of $\sim 150$ kpc with cloud covering fractions of $\sim 0.25-0.6$ \citep{maller04,kaufmann06}. |
. The association of doublet absorption in quasar spectra with normal. bright. field galaxies has been firmly established (e.g..Bergeron&Boissé1991:Steidel.del 2005). | The association of doublet absorption in quasar spectra with normal, bright, field galaxies has been firmly established \citep[e.g.,][]{bb91,sdp94,cwc-china}. |
In an effort to understand halo sizes and gas distributions. Steidel(1995.hereafterS95) searched. for foreground galaxies associated with absorption— within ~I0" (—65 kpe for z= 0.5) of quasars. | In an effort to understand halo sizes and gas distributions, \citet[][hereafter S95]{steidel95}
searched for foreground galaxies associated with absorption within $\sim10''$ $\sim65$ kpc for $z=0.5$ ) of . |
. The sample consisted of 53 absorbing and 14 non-absorbing galaxies with a A2796 equivalent width sensitivity limit of W,(2796)>0.3A. | The sample consisted of 53 absorbing and 14 non–absorbing galaxies with a $\lambda 2796$ equivalent width sensitivity limit of $W_{r}(2796) > 0.3$. |
S95 directly fitted the data by assuming a Holmberg-like luminosity scaling. and minimizing the number of non-absorbing and absorbing galaxies above and below the A(Z) relation. | S95 directly fitted the data by assuming a Holmberg–like luminosity scaling, and minimizing the number of non–absorbing and absorbing galaxies above and below the $R(L)$ relation. |
The best fit obtained clearly showed that absorbing and non-absorbing galaxies could be separated and that the halo radii R(Lx) and R(Lg) scale with luminosity with .}=0.15 and 120.2. respectively. where an L5 galaxy has a gas halo cross section ofR.=55 kpe. | The best fit obtained clearly showed that absorbing and non--absorbing galaxies could be separated and that the halo radii $R(L_K)$ and $R(L_B)$ scale with luminosity with $\beta = 0.15$ and $\beta = 0.2$, respectively, where an $L_B^{\ast}$ galaxy has a gas halo cross section of $R_{\ast} =
55$ kpc. |
Furthermore. since almost none of the absorbing galaxies were observed abovethe R(L) boundary and that almost none of the non-absorbing galaxies were observed below the R(L) boundary. S95 inferred that | Furthermore, since almost none of the absorbing galaxies were observed abovethe $R(L)$ boundary and that almost none of the non–absorbing galaxies were observed below the $R(L)$ boundary, S95 inferred that |
Furthermore. since almost none of the absorbing galaxies were observed abovethe R(L) boundary and that almost none of the non-absorbing galaxies were observed below the R(L) boundary. S95 inferred that | Furthermore, since almost none of the absorbing galaxies were observed abovethe $R(L)$ boundary and that almost none of the non–absorbing galaxies were observed below the $R(L)$ boundary, S95 inferred that |
Furthermore. since almost none of the absorbing galaxies were observed abovethe R(L) boundary and that almost none of the non-absorbing galaxies were observed below the R(L) boundary. S95 inferred that « | Furthermore, since almost none of the absorbing galaxies were observed abovethe $R(L)$ boundary and that almost none of the non–absorbing galaxies were observed below the $R(L)$ boundary, S95 inferred that |
Furthermore. since almost none of the absorbing galaxies were observed abovethe R(L) boundary and that almost none of the non-absorbing galaxies were observed below the R(L) boundary. S95 inferred that «/ | Furthermore, since almost none of the absorbing galaxies were observed abovethe $R(L)$ boundary and that almost none of the non–absorbing galaxies were observed below the $R(L)$ boundary, S95 inferred that |
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