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We manipulate the familiar expression (Frank. King. Raine 2002). substituting azimuthal velocity for the radius via equation 7.. and obtain where Αγιο is the mass accretion rate in units 10Ιόσς T
We manipulate the familiar expression (Frank, King, Raine 2002), substituting azimuthal velocity for the radius via equation \ref{rnot}, and obtain where ${\dot{M}_{16}}$ is the mass accretion rate in units $10^{16}$ g $^{-1}$ .
For the X-ray luminosities given in Table |.. we find M222x10? x107M. yr |).
For the X-ray luminosities given in Table \ref{tab:data}, we find $\dot{M}\approx2\times10^{19}$ g $^{-1}$ $3\times10^{-7}$ $^{-1}$ ).
If we assume the “edge-on” velocities discussed in relation to Figure 3 and @—0.l. equations 13 and I+ require masses 70 and 4000ως or CXOJO33831.8-352604 and RZ2109. respectively.
If we assume the “edge-on” velocities discussed in relation to Figure \ref{fig:maxL} and $\alpha=0.1$, equations \ref{tvisc} and \ref{tvisc2} require masses $70$ and $4000\Msun$ for CXOJ033831.8-352604 and RZ2109, respectively.
These numbers are surprisingly similar to the minimum masses found in Section 2..
These numbers are surprisingly similar to the minimum masses found in Section \ref{maxlum}. .
The minimum masses are 40 times larger if the disc has Z220.3Z..
The minimum masses are $\sim40$ times larger if the disc has $Z\approx0.3 Z_\odot$.
The viscous time-scale A44.300 vr.
The viscous time-scale $t_{\mathrm{visc}}\gtrsim300$ yr.
Temperatures derived with the @-dise prescription. standard assumptions. and the above limits are too cold by a factor ~10. Mj. M.
Temperatures derived with the $\alpha$ -disc prescription, standard assumptions, and the above limits are too cold by a factor $\sim10$. $M_1$, $\dot{M}$,
and a all enter as fairly weak powers and offer no obvious solution.
and $\alpha$ all enter as fairly weak powers and offer no obvious solution.
Irradiation of a surface laver (e.g.. Dubus et al.
Irradiation of a surface layer (e.g., Dubus et al.
1999) might vield the necessary temperatures.
1999) might yield the necessary temperatures.
Can a globular cluster form and sustain the massive accretion discs implied by the observed luminosities?
Can a globular cluster form and sustain the massive accretion discs implied by the observed luminosities?
Combining the minimum masses obtained in Section 2. with the rotational velocities derived above. we obtain ro222.x10 and 5x1014 em. for CXOJO33831.8-352604 and. RZ2109. respectively.
Combining the minimum masses obtained in Section \ref{maxlum} with the rotational velocities derived above, we obtain $r_{0} \approx 2\times10^{14}$ and $5\times10^{14}$ cm, for CXOJ033831.8-352604 and RZ2109, respectively.
We take this as the distance from the center of the compact object to the Lagrange point Lj.
We take this as the distance from the center of the compact object to the Lagrange point $L_1$.
We then assume a 1M. donor star (giving mass ratios g=0.01 and 1.7x10 *y and employ approximations from Eggleton (1983) and Frank. King Raine (2002). and we obtain donor Roche lobe radii rp1075 and 5x107 em. respectively.
We then assume a $1 \Msun$ donor star (giving mass ratios $q = 0.01$ and $1.7\times10^{-4}$ ) and employ approximations from Eggleton (1983) and Frank, King Raine (2002), and we obtain donor Roche lobe radii $r_L \gtrsim10^{13}$ and $5\times10^{12}$ cm, respectively.
These radii are comfortably within the size limits of the red giants that are common in globular clusters. but they represent lower limits.
These radii are comfortably within the size limits of the red giants that are common in globular clusters, but they represent lower limits.
Larger central object masses would correspond to larger overflow radii (because the dependence of rpο on M via q is not enough to overcome roος AY).
Larger central object masses would correspond to larger overflow radii (because the dependence of $r_L/r_0$ on $M$ via $q$ is not enough to overcome $r_0\propto M$ ).
White dwarfs have been mentioned because of the apparent high-metallicity in both systems.
White dwarfs have been mentioned because of the apparent high-metallicity in both systems.
With prototypical My.=0.6M. and r4=I0? em. we can exclude white dwarfs as Roche-lobe overflow donors in either system.
With prototypical $M_{wd} = 0.6\Msun$ and $r_{wd} = 10^9$ cm, we can exclude white dwarfs as Roche-lobe overflow donors in either system.
Does the envelope of a red giant contain enough oxygen to produce the observed line luminosities?
Does the envelope of a red giant contain enough oxygen to produce the observed line luminosities?
Carretta Gratton (1997) considered the metallicities of red giants in 24+ galactic globular clusters.
Carretta Gratton (1997) considered the metallicities of red giants in 24 galactic globular clusters.
They found sub-solar metallicities for every star in their study.
They found sub-solar metallicities for every star in their study.
The majority were sub-solar by at least an order of magnitude.
The majority were sub-solar by at least an order of magnitude.
The most metal-rich had Z20.32..
The most metal-rich had $Z\approx0.3 Z_\odot$.
The [ΟΠΗ/ΗΡ and ΠΝΠΙΗα values IIO. observed in CXOJ033831.8-352604 Co5 and =7. respectively) and the [ΙΟΠΠ/ΠΡ value ZO08 observed in RZ2109 (~ 30) are not necessarily indicative of high-metallicity gas.
The $\beta$ and $\alpha$ values I10 observed in CXOJ033831.8-352604 $\gtrsim5$ and $\gtrsim7$, respectively) and the $\beta$ value Z08 observed in RZ2109 $\sim30$ ) are not necessarily indicative of high-metallicity gas.
The photoionized gas near the center of the giant HIT region 30 Doradus has Z=0.3. and [ΩΠΠ/ΗΡz6 (Pellegrini et al..
The photoionized gas near the center of the giant HII region 30 Doradus has $Z\approx0.3 Z_\odot$ and $\beta\approx6$ (Pellegrini et al.,
. submitted).
submitted).
Elevated ratios are often thought to be signs of radiative shock heating (e.g.. Dopita Sutherland 1995). as was noted by 205.
Elevated ratios are often thought to be signs of radiative shock heating (e.g., Dopita Sutherland 1995), as was noted by Z08.
Detailed attempts to distinguish between enhanced abundances and non-equilibrium heating processes are probably unwarranted without better observational constraints.
Detailed attempts to distinguish between enhanced abundances and non-equilibrium heating processes are probably unwarranted without better observational constraints.
A further complication pertains to whether the line-emitting gas is radiation- or matter-bound (McCall. Rybski. Shields 1985). although the high observed [NIT]/H@ would seem to make the latter less likely in CXOJO033831.8-352604.
A further complication pertains to whether the line-emitting gas is radiation- or matter-bound (McCall, Rybski, Shields 1985), although the high observed $\alpha$ would seem to make the latter less likely in CXOJ033831.8-352604.
If Z08 are correct in their black-hole wind explanation of RZ2109. we offer without comment the suggestion. that the observed very broad line-widths may be the low-resolution appearance of narrow emission lines atop broaderones as in a structure (e.g.. Fernandes 1999).
If Z08 are correct in their black-hole wind explanation of RZ2109, we offer without comment the suggestion that the observed very broad line-widths may be the low-resolution appearance of narrow emission lines atop broaderones as in a two-wind structure (e.g., Fernandes 1999).
RLP thanks referee John Raymond for excellent suggestions and criticisms and Joel Bregman. Gary Ferland. Jimmy Irwin. Jon Miller. EricPellegrini. Mark Reynolds. and Pete Storey for fruitful discussions.
RLP thanks referee John Raymond for excellent suggestions and criticisms and Joel Bregman, Gary Ferland, Jimmy Irwin, Jon Miller, EricPellegrini, Mark Reynolds, and Pete Storey for fruitful discussions.
The variability is à common phenomenon in quasi-stellar objects (QSOs) and provides à powerful constraint on their central engines.
The variability is a common phenomenon in quasi-stellar objects (QSOs) and provides a powerful constraint on their central engines.
In the past two decades. the optical variability research. focused. on the spectral monitoring. instead. of the pure photometric monitoring.
In the past two decades, the optical variability research focused on the spectral monitoring instead of the pure photometric monitoring.
With the active galactic nuclei (AGNs) watch and the Palomer-Green (PG) QSOs spectrophotometrical monitoring projects. the reverberation mapping method. 1.e.. exploring the correlation between the emission lines and the continuum variations. is used to investigate the inner structure in AGNs (e.g..Blandfordetal.1982:Peterson 1993).
With the active galactic nuclei (AGNs) watch and the Palomer-Green (PG) QSOs spectrophotometrical monitoring projects, the reverberation mapping method, i.e., exploring the correlation between the emission lines and the continuum variations, is used to investigate the inner structure in AGNs \citep[e.g.,][]{Blandford82, Peterson93}.
. It is found that motions of clouds in the broad line regions (BLRs) are virialized (ο.σ..Kaspietal.2000.2005;Petersonetal. 2004).
It is found that motions of clouds in the broad line regions (BLRs) are virialized \citep[e.g.,][]{Kaspi00, Kaspi05, Peterson04}.
. With the line width ofH./.... from BLRs. the empirical size-luminosity relation derived from the mapping method is used to calculate the masses of their central supermassive black holes (SMBHs:2004:Petersonetal.Greene&Ho 2005).
With the line width of, from BLRs, the empirical size-luminosity relation derived from the mapping method is used to calculate the masses of their central supermassive black holes \citep[SMBHs; e.g.,][]{Kaspi00, McLure04, Bian04, Peterson04, Greene05}.
. It is found that the emission contributes significantly to the optical and ultraviolet spectra of most AGNs.
It is found that the emission contributes significantly to the optical and ultraviolet spectra of most AGNs.
Thousands of UV emission lines blend together to form à pseudocontinuum. resulting in the "small blue bump” around wwhen they are combined with Balmer continuum emission (e.g..Willsetal.1985).
Thousands of UV emission lines blend together to form a pseudocontinuum, resulting in the “small blue bump” around when they are combined with Balmer continuum emission \citep[e.g.,][]{Wills85}.
. The optical would lead to two bumps in two sides around the AM4861À citep[e.g..]| Boroson92..
The optical would lead to two bumps in two sides around the $\lambda 4861$ \\citep[e.g.,][]{Boroson92}.
It is found that the flux ratio of to.. Rpg. where the optical flux is the flux of the eemission between A4434 and A4684. strongly correlates with the so-called Eigenvector 1. which is suggested to be driven by the accretion rate (e.g..Boroson&Green.1992;Marzianietal. 2003a).
It is found that the flux ratio of to, $R_{\rm Fe}$, where the optical flux is the flux of the emission between $\lambda$ 4434 and $\lambda$ 4684, strongly correlates with the so-called Eigenvector 1, which is suggested to be driven by the accretion rate \citep[e.g.,][]{Boroson92, Marziani03a}.
. The origin of the optical/UV emission is still an open question.
The origin of the optical/UV emission is still an open question.
It is found that photoionized BLRs cannot produce the observed shape and strength of the optical emission and that the strength of UV cannot be explained unless considering the micro-turbulence of hundreds of or the collisional excitation in warm. dense gas (Baldwinetal. 2004).
It is found that photoionized BLRs cannot produce the observed shape and strength of the optical emission and that the strength of UV cannot be explained unless considering the micro-turbulence of hundreds of or the collisional excitation in warm, dense gas \citep{Baldwin04}.
. However. Vestergaard&Peterson(2005) found the correlation between the optical variance and the continuum variance and suggested that the optical His due to the line fluorescent in a photoionized plasma.
However, \citet{Vestergaard05} found the correlation between the optical variance and the continuum variance and suggested that the optical is due to the line fluorescent in a photoionized plasma.
It suggests that the optical line do not come from the same region as the UV emission (e.g..Kuehnetal.2008).
It suggests that the optical line do not come from the same region as the UV emission \citep[e.g.,][]{Kuehn08}.
. Maozetal.(1993) found that the reverberation time lag of UV in NGC 5548 is about 10 days. similar to time lag. smaller than the time lag.
\citet{Maoz93} found that the reverberation time lag of UV in NGC 5548 is about 10 days, similar to time lag, smaller than the time lag.
The reverberation measurement for the optical emission has not fared so well.
The reverberation measurement for the optical emission has not fared so well.
Some suggested that the optical emission is produced in the same region as the other broad emission lines. and some suggested that it is in the outerportion of the BLRs because of narrower FWHM of with respect to (e.g..Laoretal.1997:Marziant 2005).
Some suggested that the optical emission is produced in the same region as the other broad emission lines, and some suggested that it is in the outerportion of the BLRs because of narrower FWHM of with respect to \citep[e.g.,][]{Laor97, Marziani03a, Vestergaard05, Kuehn08}. .
Recently. Huetal...(2008a.b) did à systematic analysis of
Recently, \citet{Hu2008a, Hu2008b} did a systematic analysis of
as shown in Table 1.
, as shown in Table \ref{table1}.
We adopt the stellar mass as 0.5 M. "EMnGaudietal.2008).. and assume its radius to be ~ ON RH.
We adopt the stellar mass as 0.5 $M_{\odot}$ \citep{gaudi08}, and assume its radius to be $\sim$ 0.8 $R_{\odot}$.
each group. the simply difference between individual runs is theinclination of the outer giant. OGLE-n006-DLC-1ü9Lc.
For each group, the simply difference between individual runs is the inclination of the outer giant, OGLE-2006-BLG-109Lc.
Some details of cach group are listed in 2.. which is also summarized as follows. 3oth embrvos ancl planctesimals were considered. to eravitationally interact with cach other in our simulations.
Some details of each group are listed in Table \ref{table2}, which is also summarized as follows, Both embryos and planetesimals were considered to gravitationally interact with each other in our simulations.
Objects were allowed to collide. and we assume that hey merge into a single body with no fragmentation after a collision.
Objects were allowed to collide, and we assume that they merge into a single body with no fragmentation after a collision.
All 46. runs. were performed using a ivbrid svmplectic algorithm provided by the ALERCURY --nteeration package (Chambers1999).
All 46 runs were performed using a hybrid symplectic algorithm provided by the MERCURY integration package \citep{chambers99}.
. Each run evolved for 400 Myr. with a time step length. no more than 3.0 days ( a twentieth of a period for the body at 0.3 AUD.
Each run evolved for 400 Myr, with a time step length no more than 3.0 days $\sim$ a twentieth of a period for the body at 0.3 AU).
The Dulirsch-Stoer3 tolerance for all runs is of 1077
The Bulirsch-Stoer tolerance for all runs is of $10^{-12}$.
|n most of 16 simulations. energy is conserved to better than 1 part in )". and the angular momenta conserved in 10+
In most of the simulations, energy is conserved to better than 1 part in $10^{-3}$, and the angular momenta conserved in $10^{-11}$.
Next. we will briellv introduce simulation outcomes.
Next, we will briefly introduce simulation outcomes.
From the simulations. we can note that a typical accretion scenario occurs in the late stage planet formation (ChambersPOOL:Jietal. 2011).
From the simulations, we can note that a typical accretion scenario occurs in the late stage planet formation \citep{cham01, ji11}.
. In the beginning. the planetary embryos. anc planctesimals are quickly excited to highlv-eccentric orbits due to dramatic perturbations arising from OCGLIZ-2006-DLCi-109Lb.c. Subsequently. frequent orbital crossings emerge when the objects approach cach other. which may result in violate collisions amongst planetesimals ancl embryos.
In the beginning, the planetary embryos and planetesimals are quickly excited to highly-eccentric orbits due to dramatic perturbations arising from OGLE-2006-BLG-109Lb,c. Subsequently, frequent orbital crossings emerge when the objects approach each other, which may result in violate collisions amongst planetesimals and embryos.
Herein a large portion of the runs show that the planetesimal disk becomes quite turbulent within 1. Myr. while in some cases their chaotic period of movement could
Herein a large portion of the runs show that the planetesimal disk becomes quite turbulent within 1 Myr, while in some cases their chaotic period of movement could
derivative with respect to a.
derivative with respect to $a$.
The equations are coupled via ó=a’/(ayana3)—-1,b j, and Ίοχι/.
The equations are coupled via $\delta=a^3/(a_1 a_2 a_3)-1$, $b_j$, and $\lambda_{\mathrm{ext,} j}$.
The latter two are the internal and external shear contributions with and respectively, where the subscript ‘i’ here and in the following denotes the initial time.
The latter two are the internal and external shear contributions with and respectively, where the subscript `i' here and in the following denotes the initial time.
We use a combination of both shear models, called themodel,, which describes the external shear by a linear approximation until the turn-around of an axis, and then switches smoothly to the non-linear approximation.
We use a combination of both shear models, called the, which describes the external shear by a linear approximation until the turn-around of an axis, and then switches smoothly to the non-linear approximation.
The initial conditions for the evolution of the axes are derived from the Zel’dovich approximation and given by The initial ellipticity e; and prolaticity p; of the model are related to the initial overdensity ó; by where o? is the variance of the matter power spectrum.
The initial conditions for the evolution of the axes are derived from the Zel'dovich approximation and given by The initial ellipticity $e_\mathrm{i}$ and prolaticity $p_\mathrm{i}$ of the model are related to the initial overdensity $\delta_\mathrm{i}$ by where $\sigma^2$ is the variance of the matter power spectrum.
These values follow from the probability distribution of the eigenvalues of the Zel'dovich deformation tensor (?)..
These values follow from the probability distribution of the eigenvalues of the Zel'dovich deformation tensor \citep{Doroshkevich1970}.
To stop the collapse, we use the following virialisation conditions for each axis, derived from the tensor virial theorem, The most important difference compared to the spherical collapse model is the circumstance that the parameters 6, and A, become mass- or scale-dependent, respectively, and can belarger (6, for small masses) or smaller (A, for large masses) by a factor of ~2 compared to the canonical values for the ACDM cosmology.
To stop the collapse, we use the following virialisation conditions for each axis, derived from the tensor virial theorem, The most important difference compared to the spherical collapse model is the circumstance that the parameters $\delta_\cc$ and $\Delta_\vv$ become mass- or scale-dependent, respectively, and can belarger $\delta_\cc$ for small masses) or smaller $\Delta_\vv$ for large masses) by a factor of $\sim2$ compared to the canonical values for the $\Lambda$ CDM cosmology.
We refer to Fig.
We refer to Fig.
5 of ? for their detailed dependence on mass and redshift.
5 of \citet{Angrick2010} for their detailed dependence on mass and redshift.
In the following, we will use the results of the ellipsoidal-collapse model and implement it in our formalism for the X-ray temperature function where results from the spherical collapse were used.
In the following, we will use the results of the ellipsoidal-collapse model and implement it in our formalism for the X-ray temperature function where results from the spherical collapse were used.
We have to modify Eq.
We have to modify Eq.
since in the ellipsoidal-collapse case, the critical Laplacian A®, is now a function of the variable AO, which one has to integrate over, through Υ--2Φι/ΔΦ (see Sect. 2.1)).
since in the ellipsoidal-collapse case, the critical Laplacian $\Delta\Phi_\cc$ is now a function of the variable $\Delta\Phi$, which one has to integrate over, through $R=\sqrt{-2\Phi_\mathrm{l}/\Delta\Phi}$ (see Sect. \ref{subsec:tempFuncSph}) ).
It thus becomes where 0 is Heaviside's step function.
It thus becomes where $\uptheta$ is Heaviside's step function.
Note that it is still a good approximation to smooth the density field with an isotropic top-hat of size R and not to introduce an anisotropic smoothing function for the ellipsoidal-collapse model at the initial time, when we are well within the linear regime, and the deviation from sphericity is of order a few times 10:53.
Note that it is still a good approximation to smooth the density field with an isotropic top-hat of size $R$ and not to introduce an anisotropic smoothing function for the ellipsoidal-collapse model at the initial time, when we are well within the linear regime, and the deviation from sphericity is of order a few times $10^{-5}$.
The potential in the centre of an ellipsoid is given by (?,p.57)..
The potential in the centre of an ellipsoid is given by \citep[][p.~57]{Binney1987}.
For a sphere, αι=a»a3R/Rpx, and the integral can be solved analytically yielding 2Rpx/R, hence the result for the sphere, Eq.(8),
For a sphere, $a_1=a_2=a_3=R/R_\mathrm{pk}$, and the integral can be solved analytically yielding $2R_\mathrm{pk}/R$, hence the result for the sphere, Eq.,
, is reproduced.
is reproduced.
We proceed exactly in the same way as in the previous section, calculating the ratio between linear and non-linear potential at the time of collapse.
We proceed exactly in the same way as in the previous section, calculating the ratio between linear and non-linear potential at the time of collapse.
Again, quantities at a small initial scale factor a; are labelled with the index ‘i’ and quantities at the time of collapse with ‘c’.
Again, quantities at a small initial scale factor $a_\mathrm{i}$ are labelled with the index `i' and quantities at the time of collapse with `c'.
Using the approximations aj)3dioai3*αι, and Δι~1, we arrive at where we have also used the fact that for the virial overdensity, we can approximate ὃν=A,—1& since Ay is of order 100.
Using the approximations $a_\mathrm{i,1}\approx a_\mathrm{i,2} \approx a_\mathrm{i,3} \approx a_\mathrm{i}$, and $\Delta_\mathrm{i}\approx 1$, we arrive at where we have also used the fact that for the virial overdensity, we can approximate $\delta_\vv=\Delta_\vv-1\approx\Delta_\vv$ since $\Delta_\vv$ is of order 100.
All quantities that are necessary to evaluate Eq.
All quantities that are necessary to evaluate Eq.
can be calculated using the ellipsoidal-collapse model by ?..
can be calculated using the ellipsoidal-collapse model by \citet{Angrick2010}. .
Note that in the ellipsoidal case, the ratio of non-linear and linear potential becomes dependent on both Φι and Ad.
Note that in the ellipsoidal case, the ratio of non-linear and linear potential becomes dependent on both $\Phi_\mathrm{l}$ and $\Delta\Phi$.
To infer an averaged linear potential for a given non-linear one, we marginalise over the dependence on A® weighted by Π(Φι,A®) as follows, where Φι=Φι(Φμι,AD)via Eq.
To infer an averaged linear potential for a given non-linear one, we marginalise over the dependence on $\Delta\Phi$ weighted by $\tilde{n}(\Phi_\mathrm{l},\Delta\Phi)$ as follows, where $\Phi_\mathrm{l}=\Phi_\mathrm{l}(\Phi_\mathrm{nl},\Delta\Phi)$via Eq.
(21).. We now compare the analytic results for both X-ray temperature functions, using the spherical- and the ellipsoidal-collapse dynamics, to a hydrodynamical simulation by ? for a flat concordance ACDM model with matter density 0.3, baryon density Qpao=0.04, Hubble constant Hp=100h km s! Mpc"! with 7=0.7, and normalisation of the power spectrum σε=0.8 in a box of side-length 192ho! Mpc, starting at redshift Zstar~46.
We now compare the analytic results for both X-ray temperature functions, using the spherical- and the ellipsoidal-collapse dynamics, to a hydrodynamical simulation by \citet{Borgani2004} for a flat concordance $\Lambda$ CDM model with matter density $\Omega_\mathrm{m0}=0.3$ , baryon density $\Omega_\mathrm{bar0}=0.04$, Hubble constant $H_0=100~h$ km $^{-1}$ $^{-1}$ with $h=0.7$, and normalisation of the power spectrum $\sigma_8=0.8$ in a box of side-length $192~h^{-1}$ Mpc, starting at redshift $z_\mathrm{start}\simeq46$.
The gas physics was implemented usingGADGET-2,, a massively parallel N-body/SPH tree code with fully adaptive time-resolution (?)..
The gas physics was implemented using, a massively parallel $N$ -body/SPH tree code with fully adaptive time-resolution \citep{Springel2005}. .
The density field was sampled with 480? dark matter and an equal amount of gas particleswith masses Mpm=6.6x10°Μο and Mea=9.9x108 Mo, respectively.
The density field was sampled with $480^3$ dark matter and an equal amount of gas particleswith masses $M_\mathrm{DM}=6.6\times10^9~\mathrm{M}_\odot$ and $M_\mathrm{gas}=9.9\times10^{8}~\mathrm{M}_\odot$ , respectively.
During the time evolution, the number of gas particles decreases due to their conversion into star particles, which have slightly lower mass than the gas particles.
During the time evolution, the number of gas particles decreases due to their conversion into star particles, which have slightly lower mass than the gas particles.
((2000) stummarize the results of optical identification for a sub-sample of the LSS sources. consisting of 31 sources detected in the 27 keV. band with the SIS.
(2000) summarize the results of optical identification for a sub-sample of the LSS sources, consisting of 34 sources detected in the 2–7 keV band with the SIS.
The najor advantage of this sample compared with other ssurvevs is eood position accuracy: it ix 0.6 arciuiu im radius from the ddata alouc. thauks to superior positional resolution of the SIS.
The major advantage of this sample compared with other surveys is good position accuracy; it is 0.6 arcmin in radius from the data alone, thanks to superior positional resolution of the SIS.
To improve the osition accuracy further. we made follow-up obscrvations withROSAT IIRI over a part of the LSS field in 11997.
To improve the position accuracy further, we made follow-up observations with HRI over a part of the LSS field in 1997.
Optical spectroscopic observations were made using the University of Παπ 887 telescope. the Calar Alto 3.512. telescope. aud he Witt Peak National Observatories Mavall [ui auc 2.lin telescopes.
Optical spectroscopic observations were made using the University of Hawaii $''$ telescope, the Calar Alto 3.5m telescope, and the Kitt Peak National Observatories Mayall 4m and 2.1m telescopes.
Out of the 31 sources. 30 are identified as Αννα, 2 are clusters of galaxies. 1 is a Galactic star. and oulv 1 object remains unidentified.
Out of the 34 sources, 30 are identified as AGNs, 2 are clusters of galaxies, 1 is a Galactic star, and only 1 object remains unidentified.
The identification as ACNs is based on existence of a broad emission line or the line ratios of narrow enissiou lines ([NITJ6583 4/TI aud/or 9): κου Alivama 22000 and references therein.
The identification as AGNs is based on existence of a broad emission line or the line ratios of narrow emission lines $\AA$ $\alpha$ and/or $\AA$ $\beta$ ); see Akiyama 2000 and references therein.
Figure Ἱ(α) shows the correlation between the redshift and the apparent photou iudex in the 0.710 keV range. which is obtained from a spectral fit assuiiug no intrinsic absorption. for the ideutified objects.
Figure 1(a) shows the correlation between the redshift and the apparent photon index in the 0.7–10 keV range, which is obtained from a spectral fit assuming no intrinsic absorption, for the identified objects.
The 5 sources hat have an apparent photon iudex smaller than 1.0 are identified as { narrow-line ACUNS and 1 weak broad-line ACN. all ave located at redshift siualler than 0.5.
The 5 sources that have an apparent photon index smaller than 1.0 are identified as 4 narrow-line AGNs and 1 weak broad-line AGN, all are located at redshift smaller than 0.5.
Ou the other hand. N-rav spectra of the other ACNs are consistent with those of jearby type 1 Sevfert galaxies.
On the other hand, X-ray spectra of the other AGNs are consistent with those of nearby type 1 Seyfert galaxies.
Four Ligh redshift broad-lue ACNs show somewhat apparently hard spectra with an appareut photon index of 1.3+0.3. although it nay be still mareimal due to the limited. statistics.
Four high redshift broad-line AGNs show somewhat apparently hard spectra with an apparent photon index of $1.3\pm0.3$, although it may be still marginal due to the limited statistics.
To avoid complexity in classifviug the ACNs by the optical spectra. we divide 1ο identified ACNs into two using the N-vay data: the "absorbed? ACNs which show intrinsic absorption with a column deusity of Ng>107 7 aud the AGNs with Ng<1077.
To avoid complexity in classifying the AGNs by the optical spectra, we divide the identified AGNs into two using the X-ray data: the “absorbed” AGNs which show intrinsic absorption with a column density of $N_{\rm H} > 10^{22}$ $^{-2}$ and the ``less-absorbed'' AGNs with $N_{\rm H} < 10^{22}$.