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The clearest candidates for such tidally stripped nuclei are the two respective most massive UCDs in the Fornax and Virgo cluster (e.g. ??)), which have extended low-surface brightness envelopes.
The clearest candidates for such tidally stripped nuclei are the two respective most massive UCDs in the Fornax and Virgo cluster (e.g. \citealp{Evstigneeva+07,Hilker+07}) ), which have extended low-surface brightness envelopes.
They have masses M~10°Mo and clearly stand out among all other UCDs in terms of their sizes and luminosities, see also the location in the M—o plane in Fig.
They have masses $M \simeq 10^8 M_{\odot}$ and clearly stand out among all other UCDs in terms of their sizes and luminosities, see also the location in the $- \sigma$ plane in Fig.
8 for the most massive Fornax UCD.
\ref{figfjrmet} for the most massive Fornax UCD.
In contrast, the bulk of fainter, lower mass UCDs show a smooth luminosity and size transition towards the regime of ordinary globular clusters, and their number counts are typically well accounted for by an extrapolation of the globular cluster luminosity function (e.g. ?)).
In contrast, the bulk of fainter, lower mass UCDs show a smooth luminosity and size transition towards the regime of ordinary globular clusters, and their number counts are typically well accounted for by an extrapolation of the globular cluster luminosity function (e.g. \citealp{MHI04}) ).
This is consistent with a scenario where the bulk of these UCDs (in particular the metal-rich ones) formed when the majority of the red GCs and the spheroids of their host galaxies were created (either via a monolithic collapse or multiple wet mergers).
This is consistent with a scenario where the bulk of these UCDs (in particular the metal-rich ones) formed when the majority of the red GCs and the spheroids of their host galaxies were created (either via a monolithic collapse or multiple wet mergers).
IC acknowledges the ESO Visiting Scientist programme.
IC acknowledges the ESO Visiting Scientist programme.
'This study is based on observations made with the European Southern Observatory Very Large Telescope; partially based on observations made with the NASA/ESA Hubble Space Telescope, and obtained from the Hubble Legacy Archive, which is a collaboration between the Space Telescope Science Institute (STScI/NASA), the Space Telescope European Coordinating Facility (ST-ECF/ESA) and the Canadian Astronomy Data Centre (CADC/NRC/CSA).
This study is based on observations made with the European Southern Observatory Very Large Telescope; partially based on observations made with the NASA/ESA Hubble Space Telescope, and obtained from the Hubble Legacy Archive, which is a collaboration between the Space Telescope Science Institute (STScI/NASA), the Space Telescope European Coordinating Facility (ST-ECF/ESA) and the Canadian Astronomy Data Centre (CADC/NRC/CSA).
This research has made use ofDS9,IRAF, and software tools and packages.
This research has made use of, and software tools and packages.
Using the technique described in ?,, we quantified the age and metallicity sensitive information in the absorption-line spectra at intermediate spectral resolution R=10000 in the wavelength range corresponding to the HR09 setup of FLAMES/Giraffe for a single SSP fitting.
Using the technique described in \citet{Chilingarian09}, we quantified the age and metallicity sensitive information in the absorption-line spectra at intermediate spectral resolution $R=10000$ in the wavelength range corresponding to the HR09 setup of FLAMES/Giraffe for a single SSP fitting.
We explored the relative sensitivity of different spectral features to the stellar population parameters defined in the rotated age-metallicity coordinate system η--0 (??) with the η axis parallel to the age-metallicity degeneracy (?):: The contribution of every pixel at every wavelength to the total X? depends on many parameters.
We explored the relative sensitivity of different spectral features to the stellar population parameters defined in the rotated age–metallicity coordinate system $\eta - \theta$ \citep{Chilingarian+08,Chilingarian09} with the $\eta$ axis parallel to the age–metallicity degeneracy \citep{Worthey94}: The contribution of every pixel at every wavelength to the total $\chi^2$ depends on many parameters.
In order to estimate the fitting procedure sensitivity S(A,po,...,pa) to a given parameter p; at a given wavelength A, we should compute the corresponding partial derivative of the template grid removing the global continuum shape, i.e. accountingfor the multiplicative continuum variations.
In order to estimate the fitting procedure sensitivity $S(\lambda, p_0, \dots, p_n)$ to a given parameter $p_i$ at a given wavelength $\lambda$, we should compute the corresponding partial derivative of the template grid removing the global continuum shape, i.e. accountingfor the multiplicative continuum variations.
Hence, the response to a stellar population parameter η (and similarly, 0) at a given point (fo,Zo,co) of the parameter space is expressed as: where £ is the LOSVD, T(),t,Z) is the flux of a given SSP template spectrum characterised by its age and metallicity at a given wavelength, Pj; is the p-th order multiplicative Legendre polynomial.
Hence, the response to a stellar population parameter $\eta$ (and similarly, $\theta$ ) at a given point $(t_0, Z_0, \sigma_0)$ of the parameter space is expressed as: where $\mathcal{L}$ is the LOSVD, $T(\lambda, t, Z)$ is the flux of a given SSP template spectrum characterised by its age and metallicity at a given wavelength, $P_{1p}$ is the $p$ -th order multiplicative Legendre polynomial.
The meaning of this equation is that the squared fitting residuals of a spectrum (mo+An,90) against the model (70,90) at every pixel would correspond to its contribution to the overall x? when varying 7.
The meaning of this equation is that the squared fitting residuals of a spectrum $(\eta_0 + \Delta \eta, \theta_0)$ against the model $(\eta_0, \theta_0)$ at every pixel would correspond to its contribution to the overall $\chi^2$ when varying $\eta$ .
In practice, we compute these derivatives numerically using a “one-side” approach, i.e. computing the function
In practice, we compute these derivatives numerically using a “one-side” approach, i.e. computing the function
Gamma-ray bursts (GRB) afterglows have been ascribed to synchrotron emission from relativistic shocks in an electron-proton plasma.
Gamma-ray bursts (GRB) afterglows have been ascribed to synchrotron emission from relativistic shocks in an electron-proton plasma.
Detailed studies of GRB spectra and light curves have shown that the magnetic field strength in the shocked plasma (the downstream) is a fraction of ej~lor107 of the internal energy. while the energy in the emitting electrons is a fraction of e,~107! of the internal energy (Panaitescu&Kumar2002:Yostetal.2003).
Detailed studies of GRB spectra and light curves have shown that the magnetic field strength in the shocked plasma (the downstream) is a fraction of $\epsilon_B\sim 10^{-2}- 10^{-3}$ of the internal energy, while the energy in the emitting electrons is a fraction of $\epsilon_e\sim 10^{-1}$ of the internal energy \citep{Panaitescu:02,Yost:03}.
. Recently. Eichler&Waxman(2005) have shown that if only a small fraction f of the electrons are accelerated into the high-energy nonthermal tail. the observations can be fitted with values of €p and e, that are smaller by a factor of f£ than the above estimates. as long as f is larger than the electron to proton mass ratio ni,/7.
Recently, \citet{Eichler:05} have shown that if only a small fraction $f$ of the electrons are accelerated into the high-energy nonthermal tail, the observations can be fitted with values of $\epsilon_B$ and $\epsilon_e$ that are smaller by a factor of $f$ than the above estimates, as long as $f$ is larger than the electron to proton mass ratio $m_e/m_p$.
This still implies a magnetic field with ej= 10°.
This still implies a magnetic field with $\epsilon_B\gtrsim 10^{-6}$ .
Furthermore. the measurement of linear polarization at the level of a few percent (Bjórnsson.Gudmundsson.&Jóhannesson2004:Covinoetal.2004. and references therein) implies that the magnetic field in the synchrotron emitting region must deviate from isotropy.
Furthermore, the measurement of linear polarization at the level of a few percent \citealt{Bjornsson:04,Covino:04} and references therein) implies that the magnetic field in the synchrotron emitting region must deviate from isotropy.
Simple compressional hydrodynamic amplification of a pre-existing magnetic field of the unshocked plasma (the upstream) results in eg~107? (Gruzinov2001).
Simple compressional hydrodynamic amplification of a pre-existing magnetic field of the unshocked plasma (the upstream) results in $\epsilon_B\sim 10^{-9}$ \citep{Gruzinov:01}.
. Thus the requisite magnetic field must be generated in the shock itself or in the downstream.
Thus the requisite magnetic field must be generated in the shock itself or in the downstream.
À leading candidate mechanism that produces a magnetic field near equipartition in the shock transition layer is the transverse Weibel instability (Weibel1959:Fried 1959).. as was suggested by Gruzinov&Waxman(1999) and Medvedev&Loeb(1999).
A leading candidate mechanism that produces a magnetic field near equipartition in the shock transition layer is the transverse Weibel instability \citep{Weibel:59,Fried:59}, as was suggested by \citet{GruzinovWaxman:99} and \citet{Medvedev:99}.
.! This instability and the magnetic field it produces are expected to play a crucial role in the thermalization of the upstream and in the shock dynamics.
This instability and the magnetic field it produces are expected to play a crucial role in the thermalization of the upstream and in the shock dynamics.
Recently. a great progress toward understanding unmagnetized collisionless shocks has been— made by means of two- and three-dimensional particle-in-cell (PIC) simulations (e.g.. Lee&Lampe1973:Gruzinov2001:SilvaMedvedevetal.2005:Kato 2005)).
Recently, a great progress toward understanding unmagnetized collisionless shocks has been made by means of two- and three-dimensional particle-in-cell (PIC) simulations (e.g., \citealt{Lee:73,Gruzinov:01,Silva:03, Frederiksen:04,Jaroschek:04,Medvedev:05,Kato:05}) ).
Although the simulations do not resolve the e-p shock. they do show that within a layer ~100 proton skin depths wide. where the upstream and the downstream plasma interpenetrate. the transverse Weibel instability saturates.
Although the simulations do not resolve the $e$ $p$ shock, they do show that within a layer $\sim 100$ proton skin depths wide, where the upstream and the downstream plasma interpenetrate, the transverse Weibel instability saturates.
The saturated state of the transverse Weibel instability consists of magnetically self-pinched current filaments (see. e.g.. Fig.
The saturated state of the transverse Weibel instability consists of magnetically self-pinched current filaments (see, e.g., Fig.
| in Frederiksenetal. 2004)).
1 in \citealt{Frederiksen:04}) ).
The filaments are initially about a proton plasma skin depth in diameter and are parallel to the direction of shock propagation.
The filaments are initially about a proton plasma skin depth in diameter and are parallel to the direction of shock propagation.
As such they come with a magnetic field close to equipartition (ej 0.1) that lies in theplane perpendicular to the direction of shock propagation.
As such they come with a magnetic field close to equipartition $\epsilon_B\sim 0.1$ ) that lies in theplane perpendicular to the direction of shock propagation.
However. the observed emission from GRB afterglows Is expected to be produced at a distance of over 10° plasma skin depths from the shock (e.g.. Piran2005 and references therein).
However, the observed emission from GRB afterglows is expected to be produced at a distance of over $\sim10^9$ plasma skin depths from the shock (e.g., \citealt{Piran:05} and references therein).
Thus. even if the filamentary picture correctly describes the transition layer of GRB shocks. only the late evolution of these filaments is relevant for the observed emission.
Thus, even if the filamentary picture correctly describes the transition layer of GRB shocks, only the late evolution of these filaments is relevant for the observed emission.
Gruzinov(2001) pointed out that there is no obvious theoretical justification for the perseverance of the magnetic field.
\citet{Gruzinov:01} pointed out that there is no obvious theoretical justification for the perseverance of the magnetic field.
Since the field forms only on a small scale—the plasma skin depth—one would expect that it also decays over the small distance comparable to the skin depth.
Since the field forms only on a small scale—the plasma skin depth—one would expect that it also decays over the small distance comparable to the skin depth.
On the other hand. based on a quasi-two-dimensional picture of current filaments. Medvedevetal.(2005) suggest that the interaction between neighboring filaments may result in a magnetic field of an ever growing coherence length. whereby ερ is saturated ata finite value many plasma skin depths in the downstream.
On the other hand, based on a quasi–two-dimensional picture of current filaments, \citet{Medvedev:05} suggest that the interaction between neighboring filaments may result in a magnetic field of an ever growing coherence length, whereby $\epsilon_B$ is saturated at a finite value many plasma skin depths in the downstream.
There are reasons to believe. however. that the quasi- picture is short-lived.
There are reasons to believe, however, that the quasi--two-dimensional picture is short-lived.
For example. shock compression cannot be achieved 1n the region where the filaments make such an ordered structure (Milosavljevic.Nakar.&Spitkovsky 2005).
For example, shock compression cannot be achieved in the region where the filaments make such an ordered structure \citep{Milosavljevic:05}. .
Lacking
Lacking
the 102591. work the potential level of contamination and the location of completeness limit of the data with respect to our claimed. dip. individually these results are of. low statistical significance.
the IC2391 work the potential level of contamination and the location of completeness limit of the data with respect to our claimed dip, individually these results are of low statistical significance.
Nevertheless. when all the available evidence is taken together we believe the case for a drop in the bolometric LE between MT7-MS is compelling.
Nevertheless, when all the available evidence is taken together we believe the case for a drop in the bolometric LF between M7-M8 is compelling.
We have argued above for the existence of a dip in the bolometric LE of star forming regions. voung clusters ancl the field near Tar& 27001Ix. οσο populations span a wide range of ages. hence the dip covers a large range of masses. arguing against it being related to fine structure in the IMESs.
We have argued above for the existence of a dip in the bolometric LF of star forming regions, young clusters and the field near $_{\rm eff}\approx2700$ K. These populations span a wide range of ages, hence the dip covers a large range of masses, arguing against it being related to fine structure in the IMFs.
Instead. the obvious way to explain this is a sharp fall in the Iuminositv-niass relation between spectral types MT-MS.
Instead, the obvious way to explain this is a sharp fall in the luminosity-mass relation between spectral types M7-M8.
The LE can be written as in Equation 1. where dNdim is the AME and dMdin is the slope of the Iuminosity-mass relation where M is the bolometric magnitude and m the mass.
The LF can be written as in Equation 1, where $dN/dm$ is the MF and $dM/dm$ is the slope of the luminosity-mass relation where M is the bolometric magnitude and m the mass.
Logically. i£ dMfern increases. the LE will drop.
Logically, if $dM/dm$ increases, the LF will drop.
Indeed. we believe that what has recently been interpreted as structure in the Trapezium EME (Aleunch ct al.
Indeed, we believe that what has recently been interpreted as structure in the Trapezium IMF (Meunch et al.
2002) in fact arises [rom the feature in the luminosity-mass relation reported. here.
2002) in fact arises from the feature in the luminosity-mass relation reported here.
An qualitative estimate of the form of the 1 band magnitude-mass relation for the Pleiades can be derived from the observations.
An qualitative estimate of the form of the $I-$ band magnitude-mass relation for the Pleiades can be derived from the observations.
For example. we start with the magnitucde-mass relation from the 125Myr NIZNCTGUSN model of Baralle et al. (
For example, we start with the magnitude-mass relation from the 125Myr NEXTGEN model of Baraffe et al. (
1998) and assume this holds true as [ar as the top of the gap in the sequence.
1998) and assume this holds true as far as the top of the gap in the sequence.
We could equally use the models of Burrows et al. (
We could equally use the models of Burrows et al. (
1997) or. D'AXntona Alazzitelli (1997). but for reasons outlined in our previous work (e.g Jameson et al.
1997) or D'Antona Mazzitelli (1997), but for reasons outlined in our previous work (e.g Jameson et al.
2002) we prefer the calculation of the Lyon group.
2002) we prefer the calculation of the Lyon group.
Equation 2 can be readily. derived. by integrating the ME from à mass my to à lower mass mo and again from me to an even lower mass ms. where à is the index of a powerlaw model ME and N45.N25 the number of brown cdwarfs in the mass intervals ni to ni» and me to nis respectively.
Equation 2 can be readily derived by integrating the MF from a mass $_{1}$ to a lower mass $_{2}$ and again from $_{2}$ to an even lower mass $_{3}$, where $\alpha$ is the index of a powerlaw model MF and $_{12}$ $_{23}$ the number of brown dwarfs in the mass intervals $_{1}$ to $_{2}$ and $_{2}$ to $_{3}$ respectively.
The ratio Noy /N42 is determined from observation and my anc me are taken from the NIENTEGISZN mocel just above the dip.
The ratio $_{23}$ $_{12}$ is determined from observation and $_{1}$ and $_{2}$ are taken from the NEXTGEN model just above the dip.
Assuming a ME for the Pleiades. dN/dm x aoe0.4. as recently derived by Jameson et al. (
Assuming a MF for the Pleiades, dN/dm $\propto$ $^{-\alpha}$ , $\alpha\approx0.4$, as recently derived by Jameson et al. (
2002) over a broad mass range (0981. mz 0.035M .).. Equation 2 can be solved. for ma.
2002) over a broad mass range $_{\odot}\geq$ $\geq 0.035$ $_{\odot}$ ), Equation 2 can be solved for $_{3}$.
Thus masses can be estimated down through the dip region and below as a function of 1. and hence the magnitudie-mass relation derived.
Thus masses can be estimated down through the dip region and below as a function of $I$, and hence the magnitude-mass relation derived.
We find our estimates to be relatively insensitive to the assumed index à unless mas «mo.
We find our estimates to be relatively insensitive to the assumed index $\alpha$ unless $_{3}\ll $ $_{2}$.
The result for the Pleiades. using only those candidate members with photometry consistent with them being single objects. is shown in Figure 3 (open circles).
The result for the Pleiades, using only those candidate members with photometry consistent with them being single objects, is shown in Figure 3 (open circles).
Since the component masses of the unresolved binary candidates are not known. a robust treatment of these is cillicult.
Since the component masses of the unresolved binary candidates are not known, a robust treatment of these is difficult.
However. i£ we assume all of these to be equal mass binaries and the binary fraction above the gap. where the data are incomplete. to be the same as below. then the maenitucde-mass relation is found to be nearly identical to that obtained by considering only the single brown clwarls. as above.
However, if we assume all of these to be equal mass binaries and the binary fraction above the gap, where the data are incomplete, to be the same as below, then the magnitude-mass relation is found to be nearly identical to that obtained by considering only the single brown dwarfs, as above.
We choose not to proceed. bevond £=19.5 as the survey is complete over only a small area of sky at fainter magnitudes and the statistics are poor.
We choose not to proceed beyond $I=19.5$ as the survey is complete over only a small area of sky at fainter magnitudes and the statistics are poor.
The overall result. as illastratec in Figure 3 bv the dotted line. is that bevond the beginning of the gap any brown cwarf mass may be greater jun predicted by either the ΛΙΑΝΤΟΝ or the DUSTY models.
The overall result, as illustrated in Figure 3 by the dotted line, is that beyond the beginning of the gap any brown dwarf mass may be greater than predicted by either the NEXTGEN or the DUSTY models.
This procedure has been repeated. for the @ Orionis cluster. using the £ band. LE data reported in Dejar et al. (
This procedure has been repeated for the $\sigma-$ Orionis cluster, using the $I-$ band LF data reported in Bejar et al. (
2001).
2001).
Following our argument above and since the single and binary sequences are not well resolved. in. the Bejar et al.
Following our argument above and since the single and binary sequences are not well resolved in the Bejar et al.
work. we have ignored the influence of unresolved jnaritv. treating all their candidate members as single.
work, we have ignored the influence of unresolved binarity, treating all their candidate members as single.
As 10 models of Burrows et al. (
As the models of Burrows et al. (
L997) are widely used. to oediet the magnitudes of planetary. mass objects we opt rere to use. in turn. both the SAlIwe NENTGEN mocel and 16 5Myr evolutionary. track of the Arizona group as the starting points for this caleulation.
1997) are widely used to predict the magnitudes of planetary mass objects we opt here to use, in turn, both the 5Myr NEXTGEN model and the 5Myr evolutionary track of the Arizona group as the starting points for this calculation.
The latter model has en transformed. onto the observational plane using the xlomoetric corrections and T.ar-colour relation of the former.
The latter model has been transformed onto the observational plane using the bolometric corrections and $_{\rm eff}$ -colour relation of the former.
As shown in FigureD 4 we have derived the magnit5ude-mass relation down to /=24.0. although due to the completeness imit ofthe survey. the point below £=22 provides only an upper limit on mass.
As shown in Figure 4 we have derived the magnitude-mass relation down to $I=24.0$, although due to the completeness limit of the survey, the point below $I=22$ provides only an upper limit on mass.
Except at the lowest masses where he models and our semi-empirical estimates appear to be reasonably consistent. we find once again that beyond. the vinning of the eap the mass of any brown dwarf may be ereater than predicted by the theoretical calculations.
Except at the lowest masses where the models and our semi-empirical estimates appear to be reasonably consistent, we find once again that beyond the beginning of the gap the mass of any brown dwarf may be greater than predicted by the theoretical calculations.
Carclul scrutiny of the evidence cited above suggests that the width. of the LE dip is greater in the vounger σ and @ Orionis clusters than in either the Pleiades or the field populations.
Careful scrutiny of the evidence cited above suggests that the width of the LF dip is greater in the younger $\sigma-$ and $\theta-$ Orionis clusters than in either the Pleiades or the field populations.
This is consistent with our interpretation in which the luminosity-mass relation undergoes a sharp change in form between spectral types MT-MS8.
This is consistent with our interpretation in which the luminosity-mass relation undergoes a sharp change in form between spectral types M7-M8.
. For a population older than LOOMrs. radius ds. virtually independent of mass right across the planetarybrown dwarf nis regime.
For a population older than 100Myrs, radius is virtually independent of mass right across the planetary/brown dwarf mass regime.
However. in a population with an age of only a [ow Myrs. radius decreases rapidly with mass.
However, in a population with an age of only a few Myrs, radius decreases rapidly with mass.
Fherefore. as one moves to lower masses ancl ars. Luminosity decreases more rapidly in vounger populations. and one should. find a wider dip in the LE. representative of the dip spanning a constant range in Tr.
Therefore, as one moves to lower masses and $_{\rm eff}$ s, luminosity decreases more rapidly in younger populations, and one should find a wider dip in the LF, representative of the dip spanning a constant range in $_{\rm eff}$.
From Figure 11 in Burrows οἱ al. (
From Figure 11 in Burrows et al. (
19907) we estimate that the dip in th LE should. be ~22.5. wider (in magnitudes) in à 5 Myr population than ina 125Myr population. in agreement with the observations.
1997) we estimate that the dip in the LF should be $\sim 2-2.5\times$ wider (in magnitudes) in a 5 Myr population than in a 125Myr population, in agreement with the observations.
We now speculate on thepossible. cause of a sharp local changeὃν in the form of the [uminositv-mass relation in
We now speculate on thepossible cause of a sharp local change in the form of the luminosity-mass relation in
(2007) further indicated that the correlation between these two parameters is not so simple.
(2007) further indicated that the correlation between these two parameters is not so simple.
In fact. they observed a positive trend which reverses at F 2. 1.8.. the equivalent width versus Γ plot shows a correlation for E «2 and an anticorrelation for Γ > 2.
In fact, they observed a positive trend which reverses at $\Gamma$ $\sim$ 2, i.e., the equivalent width versus $\Gamma$ plot shows a correlation for $\Gamma$ $<$2 and an anticorrelation for $\Gamma$ $>$ 2.
Interestingly. thetwo sources with a constrained value ofR reflect this behavior: IGR 16482-3036. for which the R value is 1n agreement with the equivalent width of the line. has aT = 1.62 + 0.07. while IGR J17418-1212. in which a low equivalent width of the line is found. has Γ = 2.07 + 0.05.
Interestingly, thetwo sources with a constrained value of $R$ reflect this behavior: IGR J16482-3036, for which the $R$ value is in agreement with the equivalent width of the line, has a $\Gamma$ = 1.62 $\pm$ 0.07, while IGR J17418-1212, in which a low equivalent width of the line is found, has $\Gamma$ = 2.07 $\pm$ 0.05.
Broad-band simultaneous observations. such as those performed bySuzaku. are highly recommended in order to confirm the presence of a strong reflection component in these objects.
Broad-band simultaneous observations, such as those performed by, are highly recommended in order to confirm the presence of a strong reflection component in these objects.
The IBIS instrument on board has allowed the detection of new bright Seyfert galaxies.
The IBIS instrument on board has allowed the detection of new bright Seyfert galaxies.
Here. we have presented the broad-band spectral analysis of nine type | AGN. eight of them never observed before below 10 keV with the high sensitivity ofXMM-Newton.
Here, we have presented the broad-band spectral analysis of nine type 1 AGN, eight of them never observed before below 10 keV with the high sensitivity of.
The simultaneous fitting of the EPIC/pn (or MOS) and IBIS data allowed a description of the continuum in terms of an absorbed power-law. plus a thermal soft component ard an FeK emission line.
The simultaneous fitting of the EPIC/pn (or MOS) and IBIS data allowed a description of the continuum in terms of an absorbed power-law, plus a thermal soft component and an FeK emission line.
In addition to this simple description of the continuum. we have found several additional components. such as partial covering absorption. Gaussian components and/or absorption edges.
In addition to this simple description of the continuum, we have found several additional components, such as partial covering absorption, Gaussian components and/or absorption edges.
We also checked for the presence of an exponential high energy cut-off and à Compton reflection component.
We also checked for the presence of an exponential high energy cut-off and a Compton reflection component.
Bearing in mind the limitations of our. broad-band analysis. mainly due to the non-simultaneity of the and observations. we summarize our findings as follows:
Bearing in mind the limitations of our broad-band analysis, mainly due to the non-simultaneity of the and observations, we summarize our findings as follows:
metals that escape into he loxcer-deusity IGAL
metals that escape into the lower-density IGM.
At densities n2ng or the imiportaut fine-structure coolants. the fraecmentation criterion shifts Z4: to higher values.
At densities $n > n_{\rm cr}$ for the important fine-structure coolants, the fragmentation criterion shifts $Z_{\rm crit}$ to higher values.
Cooling is most efficient at or near the deusitics. DOCEpoy. at which the fine-structure levels reach LTE.
Cooling is most efficient at or near the densities, $n \approx n_{\rm cr}$, at which the fine-structure levels reach LTE.
For this density cuviroument. the nini critical unctallicities. Zain. are [C/Teair=BLS, Παν= 3.78. [Sas=3.5L. and [Fo/TI|azo03,52,
For this density environment, the minimum critical metallicities, $Z_{\rm crit}$ , are $_{\rm crit} \approx -3.48$, $_{\rm crit} \approx -3.78$ , $_{\rm crit} = -3.54$, and $_{\rm crit} \approx -3.52$.
When all netal trausitious are included. we find a minima Z4=LOB5Z, aud Mj~11TAL... with a temperature Zí4,=199 KI. Tn this case. all the netals appear in the cooling function with the sane fractional abuudance. but these results are highly dependent ou the amount of wining in the eas. which may determine the relative abundances of cach inetal.
When all metal transitions are included, we find a minimum $Z_{\rm crit} = 10^{-4.08}~Z_{\odot}$ and $M_J \approx 117~M_{\odot}$, with a temperature $T_{\rm crit} = 199$ K. In this case, all the metals appear in the cooling function with the same fractional abundance, but these results are highly dependent on the amount of mixing in the gas, which may determine the relative abundances of each metal.
An important parameter to consider in more detail is the initial abundance of residual molectlar hydrogen.
An important parameter to consider in more detail is the initial abundance of residual molecular hydrogen.
Iu all the simulations presented in this paper. the initial fractional abundance was taken as the residual ffraction. fip=HDresÜng2.0.tof,
In all the simulations presented in this paper, the initial fractional abundance was taken as the residual fraction, $f_{\rm H2,res} = n_{\rm H2,res}/n_{\rm tot} = 2.0 \times 10^{-4}$.
We also. ran the sinele-1etallicity simulations with fuas=11ν106 (?)..
We also ran the single-metallicity simulations with $f_{\rm H2,res} = 1.1\times10^{-6}$ \citep{GaP98}.
The differences between simulations run with aand residual abundances are πια at low densities.
The differences between simulations run with and residual abundances are small at low densities.
The μια difference in Ζω IS ~505€ at nycm2000 cni5m
The maximum difference in $Z_{\rm crit}$ is $\sim50$ at $n_f \approx 2000$ $^{-3}$ .
For excitation fron the C»eround vibrational state ofIL... the most important rotational cooling lines are at 28.22 µια (JJ=2» 0) and 17.03 jan (J—=5> 1]. labeled (0-0) S(0) and (0-0) S(]). respectively.
For excitation from the ground vibrational state of, the most important rotational cooling lines are at $28.22~\mu$ m $J = 2 \rightarrow 0$ ) and $17.03~\mu$ m $J = 3 \rightarrow 1$ ), labeled (0-0) S(0) and (0-0) S(1), respectively.