Datasets:
ReadingTimeMachine
/
rtm-sgt-ocr-v1

Dataset card Data Studio Files Community
1
source
stringlengths
1
2.05k
⌀
target
stringlengths
1
11.7k
2000) found the /i. in their data violates the usual hemispheric helicity sign rule whereas « not.
2000) found the $h_{c}$ in their data violates the usual hemispheric helicity sign rule whereas $\alpha$ not.
Since its launch in September 2006, Hinode has provided us with high spatial-resolution vector magnetograms for both the descending phase of solar cycle 23 and the ascending phase of solar cycle 24.
Since its launch in September 2006, Hinode has provided us with high spatial-resolution vector magnetograms for both the descending phase of solar cycle 23 and the ascending phase of solar cycle 24.
This gives us a unique chance in this Letter to use these so-far most accurate vector magnetic field measurements to shed a light on above arguments.
This gives us a unique chance in this Letter to use these so-far most accurate vector magnetic field measurements to shed a light on above arguments.
We organize our paper as follows.
We organize our paper as follows.
In section 2. we describe the observations and data reduction.
In section 2, we describe the observations and data reduction.
In section 3, we present our analysis and results.
In section 3, we present our analysis and results.
We conclude with a discussion in the last section.
We conclude with a discussion in the last section.
We used vector manctiograms obtained by the Spectro-polarimeter (SP) aboard Hinode (Kosugi et al.
We used vector manetiograms obtained by the Spectro-polarimeter (SP) aboard Hinode (Kosugi et al.
2007).
2007).
SP/Hinode obtains line profiles of two magnetically sensitive Fe lines at 630.15 and 630.25 nm and nearby continuum, using à 0.16"x164" shit.
SP/Hinode obtains line profiles of two magnetically sensitive Fe lines at 630.15 and 630.25 nm and nearby continuum, using a $0.16''\times164''$ slit.
There are four mapping mode of operation: Normal Map. Fast Map. Dynamics and Deep Magnetogram (Tsuneta et al.
There are four mapping mode of operation: Normal Map, Fast Map, Dynamics and Deep Magnetogram (Tsuneta et al.
2008).
2008).
In this study we only use the normal maps and fast maps.
In this study we only use the normal maps and fast maps.
The resolution of these magnetograms is about 0.32"/pixel for fast maps and 0.16"/pixel for normal maps.
The resolution of these magnetograms is about $''$ /pixel for fast maps and $''$ /pixel for normal maps.
Thus. the flux of radiation enerey density along the 1-direction is constructed usine where the mass fluxes iu the L-direction. Ad1. are defined by equation [55] in SN.
Thus, the flux of radiation energy density along the 1-direction is constructed using where the mass fluxes in the 1-direction, ${\dot M}^1_{i,j}$, are defined by equation [55] in SN.
These fluxes are used to update the radiation energy uncer L-trausport via Advection [Iuxes of radiation energy density in the 2-direction are computed by where the mass fluxes in" the i2-direction.⋅ Al).⋅⋅↽≻ are defined⋅ by equation⋅ ⋉⊌⋅⋅[56] in SN.
These fluxes are used to update the radiation energy under 1-transport via Advection fluxes of radiation energy density in the 2-direction are computed by where the mass fluxes in the 2-direction, ${\dot M}^2_{i,j}$, are defined by equation [56] in SN.
(NT These Iluxes are then used to update the radiation energy due to advection iu the 2-direction via To reduce the systematic effects of the directional splitting discussed iu section 2.3.. the order in which the two advection substeps are applied is reversed each tiniestep.
These fluxes are then used to update the radiation energy due to advection in the 2-direction via To reduce the systematic effects of the directional splitting discussed in section \ref{sec:testtransport}, the order in which the two advection substeps are applied is reversed each timestep.
Since the advection terius in the RHD equations are treated time-explicitly. the caleulations may be numerically uustable unless the timestep is less than the time for waves or advection to carry energy across a erid zone (BRichtiuver Morton 1957).
Since the advection terms in the RHD equations are treated time-explicitly, the calculations may be numerically unstable unless the timestep is less than the time for waves or advection to carry energy across a grid zone (Richtmyer Morton 1957).
The wave [amily with the largest eroup speed may be either aclabatic souud waves or radiation acoustic waves. depending ou the ratio of gas aud radiation euergy deusities.
The wave family with the largest group speed may be either adiabatic sound waves or radiation acoustic waves, depending on the ratio of gas and radiation energy densities.
The souud speed lor purposes of computing tle timestep is therefore chosen to be where £P, is the stun of the gas pressure aud the largest component of the radiation pressure tensor.
The sound speed for purposes of computing the timestep is therefore chosen to be where $P_{tot}$ is the sum of the gas pressure and the largest component of the radiation pressure tensor.
The remaiucder of the calculation of the timestep proceeds as in the hydrodsuazimie portion ol ZEUS-2D (SN).
The remainder of the calculation of the timestep proceeds as in the hydrodynamic portion of ZEUS-2D (SN).
Though resulting steps may be lounger than the time for radiation to diffuse across a ZOLLE GN)?D. the implicit differencing of the radiation diffusion term re[sec:delflux)) ensures stability of the methoc.
Though resulting steps may be longer than the time for radiation to diffuse across a zone $(\Delta x)^2\over D$, the implicit differencing of the radiation diffusion term \\ref{sec:delflux}) ) ensures stability of the method.
The problems used to test the radiation module fall into three categories.
The problems used to test the radiation module fall into three categories.
Calculations outliued in the first three sections below test in isolation the heating and cooling terms
Calculations outlined in the first three sections below test in isolation the heating and cooling terms
(ARBs.2000:Fender2001.2006).. (Blandford&Ikónigl.1979:ΤΠήπιοJolius (Stiliugctuf.2001).. spectrum(iie.Yuanefal.2007).. Dropertie 2009).
\citep[XRBs, e.g.,][]{Hynes00, Fender01, Fender06}, \citep{BK79, HJ88, FB99}, \citep{Stirling+01}, \citep[e.g.,][]{Titarchuk94, MZ95, Gierlinski+97, Esin97, Esin01, Poutanen98, Cadolle+06, YZXW07}. \citep{MiRo94, MFF01, MNCFF03, MNW05, BRP06, PBR06, Kaiser06, GBD06, Kylafis+08, Maitra+09}.
y-ray (sec.e.g.Nowakefal2011).. (1keV:Shakura&Suuvaev1019], the hard X-ray cutting plasma will be mostly ionized. and thus be subject to clectromagnetic forces and iu poteutially near relativistic motion.
$\gamma$ \citep[see, e.g.,][]{Nowak+11}, \citep[$\sim 1$ keV;][]{SS73}, the hard X-ray emitting plasma will be mostly ionized, and thus be subject to electromagnetic forces and in potentially near relativistic motion.
This motion together with the size of the dominant cutting region iuplics a characteristic dynamical time that cau play au important role in the radiative properties of the svstem. and leave its mark in the spectrum.
This motion together with the size of the dominant emitting region implies a characteristic dynamical time that can play an important role in the radiative properties of the system, and leave its mark in the spectrum.
Additional clues come from the spectral behavior at the X-ray baud.
Additional clues come from the spectral behavior at the X-ray band.
In several outbursts. the N-ray spectrum shows a break betweena few to teus of keV. where the spectral iudex steepeus from a~0.5 (where FL,x(e7") toa1.
In several outbursts, the X-ray spectrum shows a break betweena few to $\sim$ tens of keV, where the spectral index steepens from $\alpha\sim0.5$ (where $F_\nu \propto \nu^{-\alpha}$ ) to $\alpha\sim1$.
Such a break can be seeu i nmmcrousobjects.
Such a break can be seen in numerousobjects.
For example. it is seen at 2 keV in the 2000 outburst of NTE J1lls8]|bso (ITIvnesetαἱ.2000:Einefαἱ,200 at 10 keV in the 2000 outburst of NTE 1950-50 (Rodriguezefal2003)... at ~ TO keV in the 2 outburst of CUN. 339-0 (Ποιαetαἱ.2005).. atom keV in the 2005 outburst of GRO 1655-10.- (Joineal.2 Ws). at o{0κο in the 1991 out] » Cre X-1 (Cuerlnskietαἱ.LOOT) and at ~7 keV iu the 2002 outburst of GRD 1758-258 (Soriaetal2011).. to mame oulv few examples.
For example, it is seen at 2 keV in the 2000 outburst of XTE J1118+480 \citep{Hynes00, Esin01}, at $\lesssim$ 40 keV in the 2000 outburst of XTE J1550-564 \citep{RCT03}, at $\sim$ 70 keV in the 2002 outburst of GX 339-4 \citep{Homan+05}, at $\approx$ 40 keV in the 2005 outburst of GRO J1655-40 \citep{JKS08}, at $\sim 40 \keV$ in the 1991 outburst of Cyg X-1 \citep{Gierlinski+97} and at $\sim 7$ keV in the 2002 outburst of GRB 1758-258 \citep{Soria+11}, to name only few examples.
These breaks are most pronounced in the “hardest” state.
These breaks are most pronounced in the 'hardest' state.
Often. these breaks disappear aud the spectrum becomes softer with the increase of the Iuimuinositv. aud the transition to the high/soft state (c.g.Dunncfαἱ. 2011).
Often, these breaks disappear and the spectrum becomes softer with the increase of the luminosity, and the transition to the high/soft state \citep[e.g.,][]{Dunn+11}.
. The existence of this class of spectral breaks is not casily explained in the framework of current models.
The existence of this class of spectral breaks is not easily explained in the framework of current models.
Several ideas include Compton reflection by cold matter. which is expected to harden the N-vav spectrum above ~10 (Lightiun&White1988:TaardtA\laraschi1995). Another idea is metal absorption bv partially ionized eas{οιmmEsinef2001).. Alternatively. the break may frou colmbination of svuchrotron spectrum at low energies. and svuchrotron solf-Comipton at hieher energies 2005)..
Several ideas include Compton reflection by cold matter, which is expected to harden the X-ray spectrum above $\sim 10$ \citep{LW88, HM93, MZ95}.. Another idea is metal absorption by partially ionized gas\citep[e.g.][]{Esin01}.. Alternatively, the break may result from combination of synchrotron spectrum at low energies, and synchrotron self-Compton at higher energies \citep[e.g.,][]{MNW05,Homan+05}. .
Nowadays. the chemical composition of globular cluster (GC) stars cannot be regarded anymore as strictly homogeneous. as implied from the classical paradigm that these old stellar aggregates were the best approximation in nature of simple stellar populations (see Gratton. Sneden and Carretta 2004 for an extensive review and references).
Nowadays, the chemical composition of globular cluster (GC) stars cannot be regarded anymore as strictly homogeneous, as implied from the classical paradigm that these old stellar aggregates were the best approximation in nature of simple stellar populations (see Gratton, Sneden and Carretta 2004 for an extensive review and references).
Recent investigations using large samples of stars observed with multi-object spectrographs at large telescopes highlight that every GC studied so far harbours at least two different stellar generations. distinct in chemical composition and age (Carretta et al.
Recent investigations using large samples of stars observed with multi-object spectrographs at large telescopes highlight that every GC studied so far harbours at least two different stellar generations, distinct in chemical composition and age (Carretta et al.
2009a. 2009b).
2009a, 2009b).
These different populations are found to differ in abundances of light elements (C. N. O. Na. Μα. ΑΙ. St. F; Smith and Martell 2003; Smith et al.
These different populations are found to differ in abundances of light elements (C, N, O, Na, Mg, Al, Si, F; Smith and Martell 2003; Smith et al.
2005: Carretta et al.
2005; Carretta et al.
2009a.b: Yong et al.
2009a,b; Yong et al.
2008a.b: Melendez and Cohen 2009. to quote a few studies: see also Gratton et al.
2008a,b; Melendez and Cohen 2009, to quote a few studies; see also Gratton et al.
2004 and references therein) involved in proton-capture reactions of H-burning at high temperature (Denisenkov and Denisenkova 1989: Langer et al.
2004 and references therein) involved in proton-capture reactions of H-burning at high temperature (Denisenkov and Denisenkova 1989; Langer et al.
1993).
1993).
Star to star abundance variations in these elements are expected to also come with differences in the main outcome of the H burning. the He content (Gratton et al.
Star to star abundance variations in these elements are expected to also come with differences in the main outcome of the H burning, the He content (Gratton et al.
2009: Bragaglia et al.
2009; Bragaglia et al.
in preparation: Prantzos and Charbonnel 2006; Ventura et al.
in preparation; Prantzos and Charbonnel 2006; Ventura et al.
2001).
2001).
Apart from a few alterations presently well understood in term of an extra-mixing episode after the red giant branch (RGB) bump (Charbonnel 1994. 1995; Charbonnel and Zahn 2007; Eggleton et al.
Apart from a few alterations presently well understood in term of an extra-mixing episode after the red giant branch (RGB) bump (Charbonnel 1994, 1995; Charbonnel and Zahn 2007; Eggleton et al.
2007) the abundance variations are inherited by currently observed. long lived GC stars from a previous stellar component/generation: both spectroscopic (Gratton et al.
2007) the abundance variations are inherited by currently observed, long lived GC stars from a previous stellar component/generation: both spectroscopic (Gratton et al.
2001: Ramirez and Cohen 2002: Carretta et al.
2001; Ramirez and Cohen 2002; Carretta et al.
2004: Piotto et al.
2004; Piotto et al.
2005) and photometric (e.g. Bedin et al.
2005) and photometric (e.g. Bedin et al.
2004) observations convincingly showed that the observed pattern of chemical composition is present also among unevolved stars on the subgiant branch and the main sequence.
2004) observations convincingly showed that the observed pattern of chemical composition is present also among unevolved stars on the subgiant branch and the main sequence.
On the other hand. apart from a few notable GCs are still found mono-metallie objects. as far as abundances of heavier elements are concerned (see Gratton et al.
On the other hand, apart from a few notable GCs are still found mono-metallic objects, as far as abundances of heavier elements are concerned (see Gratton et al.
2004 for à recent review on this subject).
2004 for a recent review on this subject).
Their heavy (Z>13) elements metallicity. usually represented by the ratio[Fe/H".. is found to be extremely homogeneous from star to star in each cluster.
Their heavy $>$ 13) elements metallicity, usually represented by the ratio, is found to be extremely homogeneous from star to star in each cluster.
The sites of production of heavy elements (in particular «—capture elements. Fe-group elements) are stars with large and intermediate initial masses. exploding as core-collapse or thermonuclear supernovae (see Wheeler. Sneden and Truran 1989).
The sites of production of heavy elements (in particular $\alpha-$ capture elements, Fe-group elements) are stars with large and intermediate initial masses, exploding as core-collapse or thermonuclear supernovae (see Wheeler, Sneden and Truran 1989).
By studying the level and the dispersion of their yields inherited by stars in GCs we have the possibility of investigating the past history of the precursors from whom the present-day globular clusters formed (see Carretta et al.
By studying the level and the dispersion of their yields inherited by stars in GCs we have the possibility of investigating the past history of the precursors from whom the present-day globular clusters formed (see Carretta et al.
20006).
2009c).
The recently completed analysis of high resolution spectra of almost 2.000 RGB stars in 19 Galactic GCs (Carretta et al.
The recently completed analysis of high resolution spectra of almost 2,000 RGB stars in 19 Galactic GCs (Carretta et al.
2009a.b and references therein) provides an unprecedented sample of stars analysed in a fully homogeneous way.
2009a,b and references therein) provides an unprecedented sample of stars analysed in a fully homogeneous way.
By exploiting these data we can examine in detail the issue of the cosmic scatter in the metallicity (hereafter the abundance of Fe-peak elements) of GCs and hopefully provide clues to the early evolution of their progenitors.
By exploiting these data we can examine in detail the issue of the cosmic scatter in the metallicity (hereafter the abundance of Fe-peak elements) of GCs and hopefully provide clues to the early evolution of their progenitors.
Moreover. abundances of iron from high resolution spectra are traditionally the calibrating points for metallicity scales based on photometric and/or low-resolution spectroscopic indices that are sensitive to metal abundances (see the excellent discussion and historical review or this issue by Kraft and Ivans 2003).
Moreover, abundances of iron from high resolution spectra are traditionally the calibrating points for metallicity scales based on photometric and/or low-resolution spectroscopic indices that are sensitive to metal abundances (see the excellent discussion and historical review on this issue by Kraft and Ivans 2003).
Thus. our sample is perfectly suited to provide robust (for statistics) and homogeneous (for technique of analysis) "pillars". onto which anchor several existing metallicity scales.
Thus, our sample is perfectly suited to provide robust (for statistics) and homogeneous (for technique of analysis) "pillars" onto which anchor several existing metallicity scales.
dips. or a factor of 16.
dips, or a factor of 16.
Thaulss in part to the exceleut stability of the specroeraph. the shape «f the correlation xofile for a eiven sar is very stable.
Thanks in part to the excellent stability of the spectrograph, the shape of the correlation profile for a given star is very stable.
We could thereOre obtain au excelleut estimate of the wings of the iutrimsic correlation profile o each sar. by averagine all profiles of he system. after alicone them at the iie:wured velocity of he star aud. bla1 all pixels wihin two profile wilis of the velocity of the other sar.
We could therefore obtain an excellent estimate of the wings of the intrinsic correlation profile of each star, by averaging all profiles of the system, after aligning them at the measured velocity of the star and blanking all pixels within two profile widths of the velocity of the other star.
The residuals of a gaussiau adjustinent to thex| average profiles are then subtraced ποια all correlatioji profiles. a he measured velocity of cach star.
The residuals of a gaussian adjustment to these average profiles are then subtracted from all correlation profiles, at the measured velocity of each star.
This «lecreases the fluctuation level iu the xofile baseline to a level of E0.05%.
This decreases the fluctuation level in the profile baseline to a level of $\pm$.
. The vyacal velocities neasured by a double gaissa fit to these correc4( xofiles have typical accuracies ofToe 0 in/s for the primary and 100 ii/s for the faimter secondary.
The radial velocities measured by a double gaussian fit to these corrected profiles have typical accuracies of 40 m/s for the primary and 100 m/s for the fainter secondary.
These residials are twice larger than would be ueasured for single-li1ος spectra with equivalent S/N ratios aud correlation dip paralucters. but show no svste1uatfic phase. dependency.
These residuals are twice larger than would be measured for single-lined spectra with equivalent S/N ratios and correlation dip parameters, but show no systematic phase dependency.
They should thus cause no svstenatic errors on nieasure poriuueters.
They should thus cause no systematic errors on measured parameters.
of 107*em iinNy. e.g.. equivalent absorption to the NLR reddening) on the reflection component does not make any significant difference in the quality of the overall fit or spectral parameters of the thermal component.
of $10^{21}$ in, e.g., equivalent absorption to the NLR reddening) on the reflection component does not make any significant difference in the quality of the overall fit or spectral parameters of the thermal component.
The observed 0.5—7 keV flux is estimated to be 2.510ere l which can be translated to the rest-frame keV luminosity of 61075Ferg
The observed 0.5–7 keV flux is estimated to be $2.5\times 10^{-13}$ , which can be translated to the rest-frame 0.7--10 keV luminosity of $6\times 10^{43}h^{-2}$.
The BeppoSAX PDS data (14-40 keV range) were taken from the archive to compare with the Chandra data.
The BeppoSAX PDS data (14–40 keV range) were taken from the archive to compare with the Chandra data.
Franceschini et al (2000) attributed the hard X-ray emission detected with the BeppoSAX PDS to a strongly absorbed. transmitted component.
Franceschini et al (2000) attributed the hard X-ray emission detected with the BeppoSAX PDS to a strongly absorbed, transmitted component.
Using the spectral fit to the Chandra data. we find that this conclusion depends on the assumed spectral slope of the source illuminating the reflecting matter.
Using the spectral fit to the Chandra data, we find that this conclusion depends on the assumed spectral slope of the source illuminating the reflecting matter.
Given the low signal-to-noise ratio of the PDS data. no good constraints can be obtained for the spectral slope however.
Given the low signal-to-noise ratio of the PDS data, no good constraints can be obtained for the spectral slope however.
We thusshow only two cases of =2 and l=1. fora primary source spectrum.
We thusshow only two cases of $\Gamma = 2$ and $\Gamma = 1.4$ for a primary source spectrum.
When Lo—2. typical for a quasar. is assumed. even with an almost face-on setting of the reflection slab in mode (which yields the hardest spectrum above 10 keV). the PDS data lie a factor of ~5 above the extrapolation of the cold reflection mode for the Chandra nuclear spectrum.
When $\Gamma = 2$, typical for a quasar, is assumed, even with an almost face-on setting of the reflection slab in model (which yields the hardest spectrum above 10 keV), the PDS data lie a factor of $\sim 5$ above the extrapolation of the cold reflection model for the Chandra nuclear spectrum.
This suggests the presence of an extra high energy component due to transmitted radiation. as proposed by Franceschini et al (2000).
This suggests the presence of an extra high energy component due to transmitted radiation, as proposed by Franceschini et al (2000).
If an absorbed power-law is fitted to this excess component. the best-fit value of the column density Is found to be 3.3l0!em ((cf..
If an absorbed power-law is fitted to this excess component, the best-fit value of the column density is found to be $3.3\times 10^{24}$ (cf.,
Franceschini et al 2000 obtained 6.6l0?!em 7).
Franceschini et al 2000 obtained $6.6\times 10^{24}$ ).
We note a large statistical error to the value and further ambiguities arising from Compton scattering and iron metallieity in such a high column density range (Matt. Pompilio La Franca 1999: Wilman Fabian 1999).
We note a large statistical error to the value and further ambiguities arising from Compton scattering and iron metallicity in such a high column density range (Matt, Pompilio La Franca 1999; Wilman Fabian 1999).
Despite all the uncertainties. the Thomson depth of the X-ray absorber should exceed unity.
Despite all the uncertainties, the Thomson depth of the X-ray absorber should exceed unity.
The apparent lack of the transmitted component in the Chandra energy range sets the lower limit of the column density to be 2.-107em7.. while the detection of the hard X-ray excess means iis no larger than 107em7...
The apparent lack of the transmitted component in the Chandra energy range sets the lower limit of the column density to be $2\times 10^{24}$, while the detection of the hard X-ray excess means is no larger than $10^{25}$.
Ajoint fit to the Chandra ACIS and BeppoSAX PDS spectra with this model gives V=10.0 for 18 degrees of freedom (see Fig.
Ajoint fit to the Chandra ACIS and BeppoSAX PDS spectra with this model gives $\chi^2=10.0$ for 18 degrees of freedom (see Fig.
4).
4).
In the case of 1=1.4. which is on the flatter side of the photon index distribution of quasars (e.g. Reeves Turner 2000). the PDS data are explained well by cold reflection best-fitting the ACTS data Gis mentioned in the previous section. the quality of the fit to the ACTS data is not sensitive to the selection of slope) with X.=10.8 for 20 degrees of freedom.
In the case of $\Gamma = 1.4$, which is on the flatter side of the photon index distribution of quasars (e.g., Reeves Turner 2000), the PDS data are explained well by cold reflection best-fitting the ACIS data (as mentioned in the previous section, the quality of the fit to the ACIS data is not sensitive to the selection of slope) with $\chi^2=10.8$ for 20 degrees of freedom.
In this case. no transmitted component is required and the column density of the X-ray absorber exceeds 0οι
In this case, no transmitted component is required and the column density of the X-ray absorber exceeds $10^{25}$.
A distinctive cold reflection feature is observed from the nucleus of IRAS 0910444109 in the Chandra spectrum whilst no primary radiation is visible.
A distinctive cold reflection feature is observed from the nucleus of IRAS 09104+4109 in the Chandra spectrum whilst no primary radiation is visible.
Detection of the primary emission at higher energies with BeppoSAX (Franceschini et al 2000) depends. as demonstrated above. on the assumed spectrum of the source.
Detection of the primary emission at higher energies with BeppoSAX (Franceschini et al 2000) depends, as demonstrated above, on the assumed spectrum of the source.
If the PDS-detected X-rays are entirely. due to reflection. only a lower limit of the primary source luminosity can be obtained which is a minimum requirement to produce the observed luminosity through cold reflection.
If the PDS-detected X-rays are entirely due to reflection, only a lower limit of the primary source luminosity can be obtained which is a minimum requirement to produce the observed luminosity through cold reflection.
The albedo (77) in the 2-10 kev band is calculated using thepzxrav model for two inclination angles. /=20° and 60°. for the reflecting slab subtending 2a in solid angle.
The albedo $\eta $ ) in the 2–10 kev band is calculated using the model for two inclination angles, $i=20^{\circ}$ and $^{\circ}$, for the reflecting slab subtending $2\pi$ in solid angle.
For Fo=1.4. 4j;=0.066 ¢= 20°) and 0.052 G= 60°).
For $\Gamma = 1.4$, $\eta = 0.066$ $i=20^{\circ}$ ) and 0.052 $i=60^{\circ}$ ).
Therefore the luminosity of the primary source is larger han=Loyte7.6«lMeres aand 9.610"erg for respective inclinations.
Therefore the luminosity of the primary source is larger than $5\times 10^{43}\eta^{-1}\simeq 7.6\times 10^{44}$ and $9.6\times 10^{44}$ for respective inclinations.
Since the fraction of reflecting surface visible to us is likely to be less than unity. the true luminosity of the orimary source should be larger than these values.
Since the fraction of reflecting surface visible to us is likely to be less than unity, the true luminosity of the primary source should be larger than these values.
We discuss below the ease in which the primary source has a xower-law of Lo=2 and hence its transmitted radiation is detected with the PDS.
We discuss below the case in which the primary source has a power-law of $\Gamma = 2$ and hence its transmitted radiation is detected with the PDS.
The 2-10 keV luminosity of the primary source corrected for the absorption (and effects of Compton scattering. Tatt et al 1999) is estimated to be 7.41074hemsE l for a spherical obscuration. and it can go up to ~3 if absorption occurs at a small cloud in the line of sight when ης3]O07em 74).
The 2–10 keV luminosity of the primary source corrected for the absorption (and effects of Compton scattering, Matt et al 1999) is estimated to be $7.4\times 10^{45}\gamma h^{-2}$ $\gamma\sim 1$ for a spherical obscuration, and it can go up to $\sim 3$ if absorption occurs at a small cloud in the line of sight when $=3\times 10^{24}$ ).