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Finally, we must keep in mind the evolutionary status of open clusters included in our sample.
Finally, we must keep in mind the evolutionary status of open clusters included in our sample.
Since cluster membership is based on the proper motion data mainly obtained in the optical spectral range, we consider our sample as representative of optical clusters or “classical” open clusters.
Since cluster membership is based on the proper motion data mainly obtained in the optical spectral range, we consider our sample as representative of optical clusters or “classical” open clusters.
The embedded objects (e.g., the nearby cluster NGC 1333) are not included in our statistics since their members would be fainter than the limiting magnitude of V«11.5 of the if observed in the optical.
The embedded objects (e.g., the nearby cluster NGC 1333) are not included in our statistics since their members would be fainter than the limiting magnitude of $V\approx11.5$ of the if observed in the optical.
Therefore, we assume the beginning of the transparency phase after the removal of the bulk of the placental matter to be a starting point of the evolution of a classical open cluster.
Therefore, we assume the beginning of the transparency phase after the removal of the bulk of the placental matter to be a starting point of the evolution of a classical open cluster.
The corresponding age fo is defined by the lowest age of our clusters, that is, about 4 Myr.
The corresponding age $t_0$ is defined by the lowest age of our clusters, that is, about 4 Myr.
a large sample of GRBs. Vergani et al. (
a large sample of GRBs, Vergani et al. (
2009) confirm the presence of this effect. ancl set the cliserepaney to a factor of ~2.
2009) confirm the presence of this effect and set the discrepancy to a factor of $\sim 2$.
Even if the reason for this discrepancy is still not fully understood. Bucezvoski Hewett (2011) show that this cliscrepaney is likely related to a lack of quasars heavily absorbed along their line of sight.
Even if the reason for this discrepancy is still not fully understood, Budezynski Hewett (2011) show that this discrepancy is likely related to a lack of quasars heavily absorbed along their line of sight.
Given this observational result. we artificially increased the number of intervening systems based. on quasar studies by a factor of 2.
Given this observational result, we artificially increased the number of intervening systems based on quasar studies by a factor of 2.
The resulting mean contribution derived. from the intervening systems is shown in Fig.
The resulting mean contribution derived from the intervening systems is shown in Fig.
2 with a dashed line.
2 with a dashed line.
This line follows nicely the increase of the intrinsic column clensity with redshift’) providing a plausible explanation for this ollect.
This line follows nicely the increase of the intrinsic column density with redshift providing a plausible explanation for this effect.
Given our complete sample of bright GRBs we investigate the connection between the NX.ray absorption and the CRB darkness.
Given our complete sample of bright GRBs we investigate the connection between the X–ray absorption and the GRB darkness.
Phe GRB darkness can be caused by several effects that can be divided into two main classes: 7) intrinsic. i.e. cue to some physical mechanism hampering the optical emission or 4) environmental. i.e. due to intrinsic absorption and/or to high redshift.
The GRB darkness can be caused by several effects that can be divided into two main classes: $i)$ intrinsic, i.e. due to some physical mechanism hampering the optical emission or $ii)$ environmental, i.e. due to intrinsic absorption and/or to high redshift.
Considering the ον values computed in Alelandri et al. (
Considering the $\beta_{OX}$ values computed in Melandri et al. (
2011). there are 19 GRBs in our complete sample that can be classified as dark according to Jakobsson et al. (
2011), there are 19 GRBs in our complete sample that can be classified as dark according to Jakobsson et al. (
2004) and 12 according to van der Horst et al. (
2004) and 12 according to van der Horst et al. (
2009).
2009).
Out of them. 4 (common to both definitions) do not have any redshift. information.
Out of them, 4 (common to both definitions) do not have any redshift information.
In Fig.
In Fig.
3 we show the distribution of sox as a function of the intrinsic column density for the GRBs of our complete sample (sve did no include the few GRBs without a redshift determination).
3 we show the distribution of $\beta_{OX}$ as a function of the intrinsic column density for the GRBs of our complete sample (we did not include the few GRBs without a redshift determination).
lt is apparent that for bursts with soy«0.5 (ie. dark according to the Jakobsson's definition) all but three (al with 0.45<Sox=0.5) have an intrinsic column density larger than log(Nyfem7)z22.
It is apparent that for bursts with $\beta_{OX}<0.5$ (i.e. dark according to the Jakobsson's definition) all but three (all with $0.45<\beta_{OX}=0.5$ ) have an intrinsic column density larger than $\log(N_H/{\rm cm^{-2}})\gsim 22$.
This sume Ng limit is vali for al 1ο bursts that are dark according to the van der Lorst’s definition (Fig.
This same $N_H$ limit is valid for all the bursts that are dark according to the van der Horst's definition (Fig.
3).
3).
Comparing the intrinsic column density. clistribution of dark and non-dark GRBs (taking only the ones with known redshift. ie. 1538 and 845 for the Jakobssons and van der Ilorst's definition. respectively) we obtain a WS probability of 2.10" CESo) for the Jakobsson’s definition and 110. (4.130) for the van der Horsts definition (the lower value is due to the smaller number of dark GRBs).
Comparing the intrinsic column density distribution of dark and non-dark GRBs (taking only the ones with known redshift, i.e. 15–38 and 8–45 for the Jakobsson's and van der Horst's definition, respectively) we obtain a KS probability of $2\times 10^{-6}$ $4.8\,\sigma$ ) for the Jakobsson's definition and $1\times 10^{-5}$ $4.4\,\sigma$ ) for the van der Horst's definition (the lower value is due to the smaller number of dark GRBs).
We also note that if the 4 GRBs classified as dark ancl without recshilt information would have a redshift in line with the mean of the sample. then they would. have an intrinsic column density Ng(z)z1077 7.
We also note that if the 4 GRBs classified as dark and without redshift information would have a redshift in line with the mean of the sample, then they would have an intrinsic column density $N_H(z)\gsim 10^{22}$ $^{-2}$.
‘Phese results indicate that the intrinsic absorption as evaluated in the X.ray band is highly correlated with the darkness ofa GIU.
These results indicate that the intrinsic absorption as evaluated in the X–ray band is highly correlated with the darkness of a GRB.
Salvaterra et al. (
Salvaterra et al. (
2011) selected a complete sample of bright GRBs with a high degree of completeness in redshift.
2011) selected a complete sample of bright GRBs with a high degree of completeness in redshift.
In a series of papers we investigate the impact. of this sample on GRB studies.
In a series of papers we investigate the impact of this sample on GRB studies.
Here. we locus on the properties of the sample with respect to the intrinsic XNrav absorption.
Here we focus on the properties of the sample with respect to the intrinsic X–ray absorption.
We found that the intrinsic column density clistribution of our complete sample is consistent with the total distribution of column clensities presented in Campana et al. (
We found that the intrinsic column density distribution of our complete sample is consistent with the total distribution of column densities presented in Campana et al. (
2010).
2010).
The mean of the two distributions are in fact 21.7£0.5 and 21.9dEO.1. respectively.
The mean of the two distributions are in fact $21.7\pm0.5$ and $21.9\pm0.1$, respectively.
This likely indicates that the GIU brightness. as well as any other bias present in the total sample of CRBs with redshift (c.g. dust). does not heavily influence the total distribution of intrinsic column densities.
This likely indicates that the GRB brightness, as well as any other bias present in the total sample of GRBs with redshift (e.g. dust), does not heavily influence the total distribution of intrinsic column densities.
At variance with the total distribution presented. in Campana ct al. (
At variance with the total distribution presented in Campana et al. (
2010). we see in the complete. sample presented here that the region at high column densities and low redshift is now more populated by GRBs.
2010), we see in the complete sample presented here that the region at high column densities and low redshift is now more populated by GRBs.
This clearly reveals a bias present in the non-complete sample. where this region is heavily uncerpopulated due to the lack of a redshift determination of dark bursts.
This clearly reveals a bias present in the non-complete sample, where this region is heavily underpopulated due to the lack of a redshift determination of dark bursts.
Even if not statistically compelling there is an increase of the intrinsic column density with redshift. (this is more apparent in the full sample ofGRBs with redshift. Campana et al.
Even if not statistically compelling there is an increase of the intrinsic column density with redshift (this is more apparent in the full sample of GRBs with redshift, Campana et al.
2010).
2010).
We evaluate the mean contribution to INg(z) due to the intervening svstems along the GRB line of sight.
We evaluate the mean contribution to $N_H(z)$ due to the intervening systems along the GRB line of sight.
We find that. if we take into account the larger number of observed svstems allecting the line of sight of GRBs with respect to the quasar one (Verganiet al.
We find that, if we take into account the larger number of observed systems affecting the line of sight of GRBs with respect to the quasar one (Verganiet al.
2009). the population of sub-Damped Lsman-o and Daniped Lyman-a systems can account for the increase with recdshift of Ny(2).
2009), the population of sub-Damped $\alpha$ and Damped $\alpha$ systems can account for the increase with redshift of $N_H(z)$.
Lt would be interesting to confirm this directly through the study of high-z CRB lines of sight.
It would be interesting to confirm this directly through the study of $z$ GRB lines of sight.
Unfortunately this elfect plavs a significant role at very high redshift. where the number of GRB afterglow spectra is very. low.
Unfortunately this effect plays a significant role at very high redshift, where the number of GRB afterglow spectra is very low.
Ht is indeed difficult to measure absorption from Lyman-a.
It is indeed difficult to measure absorption from $\alpha$ .
However. the
However, the
We account for radiative cooling via a simple approach with a locally constant cooling time.
We account for radiative cooling via a simple approach with a locally constant cooling time.
The cooling term in the energy equation is where 58 is the internal energy of an SPILL particle. and Ris the radial location of the SPLL particle. i.e. the distance to the SAIBLE.
The cooling term in the energy equation is where $u$ is the internal energy of an SPH particle, and $R$ is the radial location of the SPH particle, i.e. the distance to the SMBH.
We parameterize the cooling time as a fixed fraction of the local cvnamical time. where fas.=1/8 and 3 is a parameter of the simulations.
We parameterize the cooling time as a fixed fraction of the local dynamical time, where $t_{\rm dyn}= 1/\Omega$ and $\beta$ is a parameter of the simulations.
“This simple model allows for a conyenicnt validation of the simulation methodology as such locally constant cooling time models. were investigated in detail in previous literature.
This simple model allows for a convenient validation of the simulation methodology as such locally constant cooling time models were investigated in detail in previous literature.
Canunie(2001) has shown that self-eravitating discs are bound to collapse if the cooling time. ωμή. is shorter than about 3/Q (Le. 3< 3).
\cite{Gammie01} has shown that self-gravitating discs are bound to collapse if the cooling time, $t_{\rm cool}$ , is shorter than about $3/\Omega$ (i.e., $\beta < 3$ ).
Rice(2005) presented a range of runs that tested: this fragmentation criterion in detail. ancl found that. for the aciabatic index of 4=5/3 (as used throughout this paper) the cisk fragments as long as ος6.
\cite{Rice05} presented a range of runs that tested this fragmentation criterion in detail, and found that, for the adiabatic index of $\gamma=5/3$ (as used throughout this paper), the disk fragments as long as $\beta \le 6$.
Note that most of the simulations explored in this paper are performed for circular initial gas orbits. ancl also for relatively small disc (total gas) masses as compared: with that of the SALBLL.
Note that most of the simulations explored in this paper are performed for circular initial gas orbits, and also for relatively small disc (total gas) masses as compared with that of the SMBH.
As will become clear later. this implies that during the simulations gas particles continue to follow nearly circular orbits. and thus their cooling time is constant aroun their orbits.
As will become clear later, this implies that during the simulations gas particles continue to follow nearly circular orbits, and thus their cooling time is constant around their orbits.
One exception to this is a test done with eccentric initial gas orbits (Section 77)).
One exception to this is a test done with eccentric initial gas orbits (Section \ref{sec:ecc}) ).
‘To compare our numerical approach with known results. we ran several tests (ο).S5 in Table 1) in a setup reminiscent of that of Tticectal.(2005).. who mocelled fragmentation of proto-stellar disces.
To compare our numerical approach with known results, we ran several tests (S1–S5 in Table 1) in a setup reminiscent of that of \cite{Rice05}, who modelled fragmentation of proto-stellar discs.
One key dillerence between marginally sel-eravitating proto-stellar ancl AGN clises is the relative disc mass.
One key difference between marginally self-gravitating proto-stellar and AGN discs is the relative disc mass.
Whereas the former become self-gravitating [or a disc to central object mass ratio of Adi/AM.~0.10.5 (ος...Ralikov.2005).. where M. isthe mass of the central star. the latter become self-gravitating for a mass ratio as small as ΑμAinc0.0030.01 (e.g...Ciaunmie. 2006)..
Whereas the former become self-gravitating for a disc to central object mass ratio of $M_{\rm disc}/M_* \sim 0.1-0.5$ \citep[e.g.,][]{Rafikov05}, where $M_*$ isthe mass of the central star, the latter become self-gravitating for a mass ratio as small as $M_{\rm disc}/\mbh \simeq 0.003-0.01$ \citep[e.g.,][]{Gammie01,Goodman03,NC05,Levin06}. .
For mareinally self-gravitating discs. the ratio of disc height to racius. £4/I. is of order of the mass ratio. HAR~MassfAdy (Gammic.2001).
For marginally self-gravitating discs, the ratio of disc height to radius, $H/R$, is of order of the mass ratio, $H/R \sim M_{\rm disc}/\mbh$ \citep{Gammie01}.
. Phe disc viscous time 15 where àO.1.Ll is the effective viscosity parameter (Lin&Pringle.1987:Gamamie.2001) for scll-eravitating discs.
The disc viscous time is where $\alpha \sim 0.1-1$ is the effective viscosity parameter \citep{Lin87,Gammie01} for self-gravitating discs.
πας. for 47/1X;0.01. the disc viscous time is some + orders of magnitude longer than the disc dynamical time.
Thus, for $H/R \simlt 0.01$, the disc viscous time is some 4--6 orders of magnitude longer than the disc dynamical time.
"his is in fact one of the reasons why the cises become selt-eravitating. as they are not able to heat up quickly enough via viscous energy dissipation (Shlosman&Begelman.1989:Gamnie. 2001).
This is in fact one of the reasons why the discs become self-gravitating, as they are not able to heat up quickly enough via viscous energy dissipation \citep{Shlosman89,Gammie01}.
. Also. since the cise viscous time is many orders of magnitude longer than the orbital time. as well as the total simulation time. we expect no radial redistribution of gas in the simulations. and indeed: very little occurs.
Also, since the disc viscous time is many orders of magnitude longer than the orbital time, as well as the total simulation time, we expect no radial re-distribution of gas in the simulations, and indeed very little occurs.
Our moclels are essentially local. as emphasized by Navakshin(2006).. and hence it sullices to simulate à small racial region of the disc.
Our models are essentially local, as emphasized by \cite{Nayakshin06a}, and hence it suffices to simulate a small radial region of the disc.
For definiteness. we simulate a gaseous. disc of mass Aly—35LOYAL.&OOLANy extending from rg=1 to rou=de where vis the dimensionless distance fromA*.
For definiteness, we simulate a gaseous disc of mass $M_{\rm d} = 3 \times 10^4 \msun \approx 0.01 \mbh$ extending from $r_{\rm in}=1$ to $r_{\rm out}=4$, where $r$ is the dimensionless distance from.
. The total initial number of ΟΡΟΙ particles used in runs S1 S5is 10°.
The total initial number of SPH particles used in runs S1–S5 is $4\times 10^6$.
Gravitational softening is adaptive. with a minimun 4.gravitational softening length of 3LO! for gas. ancl 0.001 for sink particles.
Gravitational softening is adaptive, with a minimum gravitational softening length of $3\times 10^{-4}$ for gas, and $0.001$ for sink particles.
The disk is in circular rotation around a SMDII with mass of Aij=3.5«10"M..
The disk is in circular rotation around a SMBH with mass of $\mbh = 3.5\times 10^6\msun$.
The disc is extended: vertically to a height of f/f=0.02 in theinitial conditions. whieh renders it stable to self-gravity.
The disc is extended vertically to a height of $H/R = 0.02$ in theinitial conditions, which renders it stable to self-gravity.
As the gas settles into hyelrostatic equilibrium. it heats up due to compressional heating. and then cools according to equation (4)).
As the gas settles into hydrostatic equilibrium, it heats up due to compressional heating, and then cools according to equation \ref{tcool}) ).
We used five different values of 3=£5,,4. in the runs. namely 3=0.3. 2. 3. 4.5 and 6 (see Table 1).
We used five different values of $\beta = t_{\rm cool} \Omega$, in the runs, namely $\beta = 0.3$, $2$, $3$, $4.5$ and $6$ (see Table 1).
‘The disces in runs with 7=0.3.2 and 3 fragmented and formed stars. as expected. based on the results of ice (2005).. but our tests with J=4.5 and 6 did not.
The discs in runs with $\beta=0.3, 2$ and 3 fragmented and formed stars, as expected based on the results of \cite{Rice05}, but our tests with $\beta=4.5$ and $6$ did not.
However. the run with οὐ=4.5 and 6 did fragment in the sense of forming high density gas clumps. some of which can benoted in Figure 1..
However, the run with $\beta=4.5$ and $6$ did fragment in the sense of forming high density gas clumps, some of which can benoted in Figure \ref{fig:fragm}. .
The maximunm density in these clumps fuctuated between a few to a few dozen times pp. i.c. the clumps were dense enough to be οσους. but not dense enough for our sink particlecriterion. since shoo. was set to 30 for these tests (ef.
The maximum density in these clumps fluctuated between a few to a few dozen times $\rho_{\rm BH}$ , i.e. the clumps were dense enough to be self-bound, but not dense enough for our sink particlecriterion, since $A_{\rm col}$ was set to 30 for these tests (cf.
Section ??))
Section \ref{sec:collapse}) ).
the PN detector, also separately in a soft and hard energy band.
the PN detector, also separately in a soft and hard energy band.
The shape of the light curve with its fast rise time of 1000 s and e-folding decay time of ~ 11400s points to a single, rather compact flaring structure.
The shape of the light curve with its fast rise time of 1000 s and e-folding decay time of $\sim$ 1400 s points to a single, rather compact flaring structure.
Since for X-ray emitting structures/loops the radiative energy loss scales with the square of the particle density, compact, i.e. small and dense structures, radiate their energy much faster away than larger tenuous structures of comparable X-ray luminosity.
Since for X-ray emitting structures/loops the radiative energy loss scales with the square of the particle density, compact, i.e. small and dense structures, radiate their energy much faster away than larger tenuous structures of comparable X-ray luminosity.
Further, the harder X-ray emission increases by a larger factor and its peak precedes the one of the softer emission, a behavior that is reminiscent of coronal flares.
Further, the harder X-ray emission increases by a larger factor and its peak precedes the one of the softer emission, a behavior that is reminiscent of coronal flares.
We do not detect any spatial shift between the photons detected during the quasi-quiescent and the flaring phase at an accuracy of x0.5"..
We do not detect any spatial shift between the photons detected during the quasi-quiescent and the flaring phase at an accuracy of $\lesssim 0.5$ .
In refleqq we show the quasi-quiescent part with two different time binnings.
In \\ref{lcqq} we show the quasi-quiescent part with two different time binnings.
Notably, the average quasi-quiescent X-ray count rate has increased by roughly the spectral hardness given by the ratio of the count rates in the 00.8 keV vs. 22.0 keV energy band is virtually identical compared to the pre-flare phase.
Notably, the average quasi-quiescent X-ray count rate has increased by roughly the spectral hardness given by the ratio of the count rates in the 0.8 keV vs. 2.0 keV energy band is virtually identical compared to the pre-flare phase.
We do not see a strong modulation of the X-ray light curve that might be expected in an oblique rotator model (Pj=213 ks).
We do not see a strong modulation of the X-ray light curve that might be expected in an oblique rotator model $P_{\rm rot}= 213$ ks).
Significant, roughly sinusoidal shaped rotational modulation would be present, if the star partially occults X-ray emitting material depending on rotational phase or when a disk around the star is viewed under different angles.
Significant, roughly sinusoidal shaped rotational modulation would be present, if the star partially occults X-ray emitting material depending on rotational phase or when a disk around the star is viewed under different angles.
An example are the X-ray brightness variations of the massive magnetic star 6! CC (?)..
An example are the X-ray brightness variations of the massive magnetic star $\theta ^1$ C \citep{gag05}.
This model assumes cylindrical symmetry for the system as is likely appropriate for a star with a disk and what is commonly observed when e.g. magnetic field variations are monitored, but clearly we cannot completely exclude unexpected phenomena during the not covered phase.
This model assumes cylindrical symmetry for the system as is likely appropriate for a star with a disk and what is commonly observed when e.g. magnetic field variations are monitored, but clearly we cannot completely exclude unexpected phenomena during the not covered phase.
Brightness variations on timescales of 20 ks (~ 0.1 phase) between consecutive bins during the pre-flare or the post-flare phase have an amplitude of x10 However, the strength of the modulation depends on the plasma location and the geometry of the system and for a small obliquity the effects are likely minor in any case.
Brightness variations on timescales of 20 ks $\sim$ 0.1 phase) between consecutive bins during the pre-flare or the post-flare phase have an amplitude of $\lesssim 10$ However, the strength of the modulation depends on the plasma location and the geometry of the system and for a small obliquity the effects are likely minor in any case.
The global spectral properties of IQ Aur are determined from the EPIC spectra by using photo-absorbed multi-temperature models with free metallicity.
The global spectral properties of IQ Aur are determined from the EPIC spectra by using photo-absorbed multi-temperature models with free metallicity.
Given the large variations in count rate over the observation, we separate the data into a flaring phase (PN: 558.1 ks) and a quasi-quiescent (QQ) phase for spectra analysis.
Given the large variations in count rate over the observation, we separate the data into a flaring phase (PN: 58.1 ks) and a quasi-quiescent (QQ) phase for spectra analysis.
As an example we show the respective PN spectra with the applied spectral models in refspec;; the derived spectral properties from our modeling are summarized in reffit..
As an example we show the respective PN spectra with the applied spectral models in \\ref{spec}; the derived spectral properties from our modeling are summarized in \\ref{fit}.
The description of the quasi-quiescent spectrum clearly requires, beside the 0.3 keV component seen byROSAT, the presence of additional hotter plasma components.
The description of the quasi-quiescent spectrum clearly requires, beside the 0.3 keV component seen by, the presence of additional hotter plasma components.
Applying a one temperature model for example to the MOS spectra results in an average temperature of kT=0.75 keV, albeit with xà,4>2 the fit is quite poor.
Applying a one temperature model for example to the MOS spectra results in an average temperature of $kT=0.75$ keV, albeit with $\chi^2_{\rm red}>2$ the fit is quite poor.
Models with three temperature components describe the individual data sets rather well.
Models with three temperature components describe the individual data sets rather well.
The absorption component required to model the X-ray spectra of IQ Aur is rather weak, compatible with the findings from theROSAT data.
The absorption component required to model the X-ray spectra of IQ Aur is rather weak, compatible with the findings from the data.
We again find strong emission from rather cool plasma at ~3 MK, but overall the plasma temperatures derived from the spectra are significantly higher than those determined withROSAT, where only weak indications for hotter plasma were found due to the limited sensitivity of the instrument.
We again find strong emission from rather cool plasma at $\sim3$ MK, but overall the plasma temperatures derived from the spectra are significantly higher than those determined with, where only weak indications for hotter plasma were found due to the limited sensitivity of the instrument.
We detect a significant contribution from z10 MK plasma even in the quasi-quiescent phase of IQ Aur, here it accounts for about We find that a subsolar metallicity best describes the data.
We detect a significant contribution from $\gtrsim10$ MK plasma even in the quasi-quiescent phase of IQ Aur, here it accounts for about We find that a subsolar metallicity best describes the data.
Using solar abundances reduces the emission measure, but also worsens the model quality Q,47 1.24), albeit the temperature structure remains nearly identical.
Using solar abundances reduces the emission measure, but also worsens the model quality $\chi^2_{red}=1.24$ ), albeit the temperature structure remains nearly identical.
Allowing individual abundances to vary independently does not lead to robust results, therefore a global metallicity is used in all models.
Allowing individual abundances to vary independently does not lead to robust results, therefore a global metallicity is used in all models.
The best fit model corresponds to a source flux of 1.7x107? ccm? ss“! and an X-ray luminosity of Lx=42x10? ss“! in the keV band; for theROSAT 22.4 keV band we derive a roughly This corresponds to an activity level of log Lx/Lpo1~—6.4, assuming that the X-rays are emitted by IQ Aur.
The best fit model corresponds to a source flux of $1.7 \times 10^{-13}$ $^{-2}$ $^{-1}$ and an X-ray luminosity of $L_{\rm X}= 4.2 \times 10^{29}$ $^{-1}$ in the keV band; for the 2.4 keV band we derive a roughly This corresponds to an activity level of $\log L_{\rm X}$ $L_{\rm bol} \approx -6.4$, assuming that the X-rays are emitted by IQ Aur.