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Considering theROSAT detection at a very similar X-ray luminosity of Lx=40x10? ergss! nearly 15 years ago, IQ Aur is likely emitting X-rays at a rather constant level for more than a decade. | Considering the detection at a very similar X-ray luminosity of $L_{\rm X}= 4.0 \times 10^{29}$ $^{-1}$ nearly 15 years ago, IQ Aur is likely emitting X-rays at a rather constant level for more than a decade. |
During the flare phase, IQ Aur has not only significantly brightened, but as shown in refspec, also the associated spectrum is much harder. | During the flare phase, IQ Aur has not only significantly brightened, but as shown in \\ref{spec}, also the associated spectrum is much harder. |
The flare is most prominent at high temperatures, but large excess emission is also present at lower temperatures. | The flare is most prominent at high temperatures, but large excess emission is also present at lower temperatures. |
During the peak of the strong flare the emitted X-ray luminosity of IQ Aur rises to about Lx=3.2x10?! ergss! 110.0 keV), corresponding to a flux increase of nearly a factor 100 comparedto the quasi-quiescent state. | During the peak of the strong flare the emitted X-ray luminosity of IQ Aur rises to about $L_{\rm X}= 3.2 \times 10^{31}$ $^{-1}$ 10.0 keV), corresponding to a flux increase of nearly a factor 100 comparedto the quasi-quiescent state. |
With an average | With an average |
wind collision region is very unstable since in these cases the flux can vary by greater than 450% from one snapshot to the next. | wind collision region is very unstable since in these cases the flux can vary by greater than $\pm 50\%$ from one snapshot to the next. |
As already noted by Corcoran (2001b)}) no variability is seen during the z90ksec exposure of the exatiug spectrum. | As already noted by Corcoran \cite{C2001}) ) no variability is seen during the $\approx 90\; {\rm ksec}$ exposure of the grating spectrum. |
This is to be expected. however. since the dynamical timescale is mich longer. | This is to be expected, however, since the dynamical timescale is much longer. |
Two examples of ‘time-averaged svuthetic spectra areshown iu Fie. 2.. | Two examples of `time-averaged' synthetic spectra areshown in Fig. \ref{fig:hydro}, |
next to density plots of the corresponding calculation. | next to density plots of the corresponding calculation. |
chemical abundance Y=0.98. Z=0.02) is taken to be 0 (Dewietal. 2002). | chemical abundance $Y
= 0.98$ , $ Z = 0.02$ ) is taken to be 0 \cite{dewi02}) ). |
Ina WD + He star system. due to the nuclear evolution or the loss of orbital angular momentum the He star fills its Roche lobe at He main sequence or He subgiant stage. and then begins the mass transfer. | In a WD + He star system, due to the nuclear evolution or the loss of orbital angular momentum the He star fills its Roche lobe at He main sequence or He subgiant stage, and then begins the mass transfer. |
The He-rich material from the He star is acereted by the WD. and is converted into CO heavier elements via thermonuclear burning onto the surface of the WD. | The He-rich material from the He star is accreted by the WD, and is converted into CO heavier elements via thermonuclear burning onto the surface of the WD. |
In input physics. we adopt an optically thick wind scenario (Kato&Hachisu1994.etal. 1996)) and the description for He mass accumulation efficiency onto the surface of the WD eitven by Kato Hachisu (2004). | In input physics, we adopt an optically thick wind scenario \cite{kato94,hach96}) ) and the description for He mass accumulation efficiency onto the surface of the WD given by Kato Hachisu (2004). |
Firstly. if the mass transfer rate |My.| is above a critical rate (Nomoto 1982)) we take that He burning ts steady. and the He-rich material on the surface of the WD Is converted into CO elements at a rate M.. | Firstly, if the mass transfer rate $|\dot{M}_{\rm He}|$ is above a critical rate \cite{nom82}) ) we take that He burning is steady, and the He-rich material on the surface of the WD is converted into CO elements at a rate $\dot{M}_{\rm {cr}}$. |
Secondly. if Αμ. is lower than M,, but higher than M, (Kato Hachisu 2004). which is the minimum accretion rate that the He-shell steadily burn. the He burning is thought to be steady. and all accreting material is converted into CO elements. | Secondly, if $|\dot{M}_{\rm He}|$ is lower than $\dot{M}_{\rm {cr}}$ but higher than $\dot{M}_{\rm {st}}$ (Kato Hachisu 2004), which is the minimum accretion rate that the He-shell steadily burn, the He burning is thought to be steady, and all accreting material is converted into CO elements. |
Thirdly. if |My. is lower than M, but higher than Mis,=4.0x107M.yr!. in which the weak He-shell flashes occur. a part of the envelope material of the WD ts assumed to be blown off from the surface of the WD (Woosleyetal. 1986)). | Thirdly, if $|\dot{M}_{\rm He}|$ is lower than $\dot{M}_{\rm
{st}}$ but higher than $\dot{M}_{\rm {low}}=4.0\times
10^{-8}~M_{\odot}\,\rm yr^{-1}$, in which the weak He-shell flashes occur, a part of the envelope material of the WD is assumed to be blown off from the surface of the WD \cite{woos86}) ). |
Finally. if [My] is lower than Mj... the strong He-shell flashes occur. and no material can be accumulated onto the WD. | Finally, if $|\dot{M}_{\rm He}|$ is lower than $\dot{M}_{\rm {low}}$, the strong He-shell flashes occur, and no material can be accumulated onto the WD. |
Summarizing the above prescriptions. the accumulation efficiency of the acereting He can be written as (Wang et al. | Summarizing the above prescriptions, the accumulation efficiency of the accreting He can be written as (Wang et al. |
2009a) where a” is determined by a linearly interpolated way from a grid computed by Kato&Hachisu(2004). | 2009a) where $\alpha'$ is determined by a linearly interpolated way from a grid computed by \cite{kato04}. |
. Actually. the accumulation efficiency mentioned above is limited to non-rotating WDs. and is inadequate to calculate the mass increase of rotating WD. | Actually, the accumulation efficiency mentioned above is limited to non-rotating WDs, and is inadequate to calculate the mass increase of rotating WD. |
Considering the spin-up of the WDs via accretion. Yoon Langer (2004) found that the rotation can stabilize He-shell burning. and help He-accreting CO WD grow in mass. | Considering the spin-up of the WDs via accretion, Yoon Langer (2004) found that the rotation can stabilize He-shell burning, and help He-accreting CO WD grow in mass. |
Assuming that the CO WDs rigidly rotate. the simulated results by Domínguezetal.(2006) proposed that massive progenitors result in higher ??Ni mass and explosive luminosity. and more massive WDs at the moment of explosion. | Assuming that the CO WDs rigidly rotate, the simulated results by \cite{Domi06} proposed that massive progenitors result in higher $^{56}$ Ni mass and explosive luminosity, and more massive WDs at the moment of explosion. |
To consider the influence of rotation on. the mass accumulation of WDs. we adopt the following input physics (e.g. Chen&Li 2009)). ( | To consider the influence of rotation on the mass accumulation of WDs, we adopt the following input physics (e.g. \cite{chen09}) ). ( |
1) With the mass transfer from the He star. the WDs obtain a large amount of angular momentum from the acereting material. and is spun up to a high rotation velocity 2000)). ( | 1) With the mass transfer from the He star, the WDs obtain a large amount of angular momentum from the accreting material, and is spun up to a high rotation velocity ). ( |
2) Considering the lifting effect in the hydrostatic equilibrium (Domínguezetal.1996)). we introduce an effective mass May of the WD by taking account of the centrifugal force. ( | 2) Considering the lifting effect in the hydrostatic equilibrium \cite{Domi96}) ), we introduce an effective mass $M_{\rm eff}$ of the WD by taking account of the centrifugal force. ( |
3) According to different ranges of the polar angle. we divide the surface of the WD into three zones às follows: the equatorial zone (EZ. @=60°—120°). the middle zone (MZ. 4=30°—60° and 120°—1507). and the polar zone (PZ. η=0°—30° and 150*—1807). ( | 3) According to different ranges of the polar angle, we divide the surface of the WD into three zones as follows: the equatorial zone (EZ, $\theta=60^{\circ}-120^{\circ}$ ), the middle zone (MZ, $\theta=30^{\circ}-60^{\circ} $ and $120^{\circ}-150^{\circ} $ ), and the polar zone (PZ, $\theta=0^{\circ}-30^{\circ}$ and $150^{\circ}-180^{\circ} $ ). ( |
4) Assuming that each zone accretes the transferred material at a rate proportional tc its area. we can obtain aceretion fraction fj,=0.5. 0.366. and 0.134 for EZ. MZ. and PZ. respectively. | 4) Assuming that each zone accretes the transferred material at a rate proportional to its area, we can obtain accretion fraction $f_{\rm i}=0.5$, 0.366, and 0.134 for EZ, MZ, and PZ, respectively. |
Therefore. the mass growth rate of the rotating WD is given by where>) a; is the mass accumulation efficiencies for different zones on the surface of the WD. | Therefore, the mass growth rate of the rotating WD is given by where $\alpha_{\rm i}$ is the mass accumulation efficiencies for different zones on the surface of the WD. |
The mass loss rate of the binary system is M=(1—XaifMg. which is assumed to be ejected in the vicinity of the WD in the form of isotropic winds or outflows. and taking away the specific orbital angular momentum of the WD. | The mass loss rate of the binary system is $\dot{M}=(1-\sum \alpha_{\rm i}f_{\rm
i})\dot{M}_{\rm He}$, which is assumed to be ejected in the vicinity of the WD in the form of isotropic winds or outflows, and taking away the specific orbital angular momentum of the WD. |
In our calculations. we set the initial surface velocity at the WD's equator to be 10kms!. and the radius of the WD changes with ΑοMy. | In our calculations, we set the initial surface velocity at the WD's equator to be $10~ \rm km\,s^{-1}$, and the radius of the WD changes with $R\propto M_{\rm WD}^{-1/3}$. |
The criterion that the WD differentially rotate is a key input physics in this work. | The criterion that the WD differentially rotate is a key input physics in this work. |
We adopt the conclusion derived by Yoon&Langer2004.. in which the WD should differentially rotate when its accretion rate =3x107M« yr. | We adopt the conclusion derived by \cite{yoon04b}, , in which the WD should differentially rotate when its accretion rate $\ga3\times10^{-7}M_\odot$ $\rm {r}^{-1}$. |
As a result of differential rotation. the central carbon ignition of the WD cannot occur even if its mass exceed canonical Chandrasekhar limit of 1.4.M... | As a result of differential rotation, the central carbon ignition of the WD cannot occur even if its mass exceed canonical Chandrasekhar limit of $1.4~M_\odot$. |
Once Mwp>L4M; and M<3xI07/Ma. we stop the calculation. and assume the WD to explode as a SN Ia (we use Msn to denote the explosive mass of the WD) because of no differential rotation to support the massive WD. | Once $M_{\rm WD}\ge1.4~M_\odot$ and $\dot
M<3\times10^{-7}~M_\odot$, we stop the calculation, and assume the WD to explode as a SN Ia (we use $M_{\rm SN}$ to denote the explosive mass of the WD) because of no differential rotation to support the massive WD. |
Considering the prescriptions above for the mass accumulation on the surface of the WD in Eggleton;s stellar evolution code. we have calculated the evolution of WD + He star systems. and obtained the initial parameters of the WD binaries that lead to SNe la. An example of the evolutionary sequences of a WD binary (with My;=1.8Me. Mwoi=1.2Me. and log(Pow ;/day) = -].20) are shown in Figures 2 and 3. | Considering the prescriptions above for the mass accumulation on the surface of the WD in Eggleton¡¯s stellar evolution code, we have calculated the evolution of WD + He star systems, and obtained the initial parameters of the WD binaries that lead to SNe Ia. An example of the evolutionary sequences of a WD binary (with $M_{\rm He, i}=1.8~M_\odot$, $M_{\rm WD, i}=1.2~M_\odot$, and $P_{\rm orb, i}$ /day) = -1.20) are shown in Figures 2 and 3. |
We plot the evolution of Myo. Mwp and Mywp varying with time in Figure 2. | We plot the evolution of $\dot {M}_{\rm He}$, ${\dot M}_{\rm WD}$ and $M_{\rm WD}$ varying with time in Figure 2. |
When the age of the He star is 1.76x10° vr. the Roche lobe overflow occurs. and the WD accretes the material from the He star. | When the age of the He star is $1.76\times 10^{6}$ yr, the Roche lobe overflow occurs, and the WD accretes the material from the He star. |
In the earlier phase of the mass transfer. the orbital period decreases until Pa, log(i/day =—|.24. because the material is transferred from themore massive He star to the less massive WD. and then increases when the WD mass grows and exceeds the He star mass. | In the earlier phase of the mass transfer, the orbital period decreases until $P_{\rm orb, i}$ /day $\approx -1.24$, because the material is transferred from themore massive He star to the less massive WD, and then increases when the WD mass grows and exceeds the He star mass. |
After the massexchange of 2.75x10? yr. the mass transfer rate decline to be 3xI07M.yr !'. and the WD cannot differentially rotate and SN Ia explosion ts triggered. | After the massexchange of $2.75\times 10^{5}$ yr, the mass transfer rate decline to be $3\times 10^{-7}M_{\odot}
\rm yr^{-1}$ , and the WD cannot differentially rotate and SN Ia explosion is triggered. |
At | At |
(z1 Bildsten (2003.2001: hereafter paper Taud IT) aud Cliaug. Arras. Bildsten (2001: hereafter paper IID) that II is casily destroved by diffusive nuclear burning (DNB). | $\approx 10^{-17}\Msun$ $\approx 1$ Bildsten (2003,2004; hereafter paper I and II) and Chang, Arras, Bildsten (2004; hereafter paper III) that H is easily destroyed by diffusive nuclear burning (DNB). |
The ceutral idea behind DNB. which was first proposed by Chin Salpeter (1968: see also Rosen 1970). is that II C‘an diffuse to ereat depth where the temperature and density are suffiiientlv large to consume II bv proton Cecaptures outo heavier clemenuts. | The central idea behind DNB, which was first proposed by Chiu Salpeter (1968; see also Rosen 1970), is that H can diffuse to great depth where the temperature and density are sufficiently large to consume H by proton captures onto heavier elements. |
In paper I aud IL we showed that this process is so effective that we expect NS surfaces to be depleted of any primorcial IT. The observation of II ou the surfaces of NS would then poit to late-time or contiuuous accretion. | In paper I and II, we showed that this process is so effective that we expect NS surfaces to be depleted of any primordial H. The observation of H on the surfaces of NS would then point to late-time or continuous accretion. |
This conclusiou is inscusitive to the streneth of the maeuetic field aud the size of an inert Πο buffer that sits between the IT and the underlying proton capturing clements (paper IIT. | This conclusion is insensitive to the strength of the magnetic field and the size of an inert He buffer that sits between the H and the underlying proton capturing elements (paper III). |
With all the IT constuned. one would expect surfaces of Ie. | With all the H consumed, one would expect surfaces of He. |
However. recent observations by To IIeiuke (2009) sugeest that the NS at the ceuter of the Cassiopeia A supernova remiant has a carbon surface with au effective temperature T.=Ls«101 and radius R=12.11 kan. | However, recent observations by Ho Heinke (2009) suggest that the NS at the center of the Cassiopeia A supernova remnant has a carbon surface with an effective temperature $T_{\rm e} = 1.8\times 10^6\,{\rm K}$ and radius $R = 12-14$ km. |
This paper shows that Πο is also vulucrable to diffusive nuclear burning on the surface of Cas A. All primordial IL/Ile is consumued durius its early cooling history. exposing the uuderbvius material. | This paper shows that He is also vulnerable to diffusive nuclear burning on the surface of Cas A. All primordial H/He is consumed during its early cooling history, exposing the underlying material. |
If subsequent accretion does nof cover this underlying material. we would expect a population of NSs with mid-Z surfaces. | If subsequent accretion does not cover this underlying material, we would expect a population of NSs with mid-Z surfaces. |
We first review the plivsies of DNB in refsec:plivsics..outlining the major results of papers | We first review the physics of DNB in \\ref{sec:physics}, ,outlining the major results of papers |
expectations, we now move to the more pertinent matter of how the reality of these physical properties might be revealed. | expectations, we now move to the more pertinent matter of how the reality of these physical properties might be revealed. |
The model from Fig. | The model from Fig. |
2 can be approached as if it were real data and the process of observation followed to find out if and where pitfalls in the interpretation of the data may lie. | \ref{hist} can be approached as if it were real data and the process of observation followed to find out if and where pitfalls in the interpretation of the data may lie. |
To illustrate the importance of sampling effects in high redshift surveys, the generated population of galaxies from refSSFR can be analysed under the same observational constraints, and using the same techniques as were applied by to the real data. | To illustrate the importance of sampling effects in high redshift surveys, the generated population of galaxies from \\ref{SSFR} can be analysed under the same observational constraints, and using the same techniques as were applied by to the real data. |
Fig. | Fig. |
?? shows the real and model samples on a plot of observed quantities: rest frame visible magnitude vs. rest-frame UVmagnitude*. | \ref{data} shows the real and model samples on a plot of observed quantities: rest frame visible magnitude vs. rest-frame UV. |
. The two sets of points are somewhat offset from each other, but the statistical significance of this difference is low; the majority of the population (i.e. the fainter galaxies) are overlapping. | The two sets of points are somewhat offset from each other, but the statistical significance of this difference is low; the majority of the population (i.e. the fainter galaxies) are overlapping. |
So, for the purposes of this purely illustrative exercise, we consider this model to be an acceptable match to the data. ( | So, for the purposes of this purely illustrative exercise, we consider this model to be an acceptable match to the data. ( |
For a discussion of discrepancies that exist between current semi-analytic models and recent observations, the reader is referred to 2010)). | For a discussion of discrepancies that exist between current semi-analytic models and recent observations, the reader is referred to ). |
With this caveat, we proceed to follow our model sample all the way through from the “real” physical parameters to the magnitudes that would be observed, and then back again to theinferred physical parameters. | With this caveat, we proceed to follow our model sample all the way through from the “real” physical parameters to the magnitudes that would be observed, and then back again to the physical parameters. |
This process from physical quantities to observables, and back, is shown as a sequence of panels in Fig. 4.. | This process from physical quantities to observables, and back, is shown as a sequence of panels in Fig. \ref{mock}. |
Each transition (clockwise) from one panel to the next introduces one part of this chain, as follows: The top left panel of Fig. | Each transition (clockwise) from one panel to the next introduces one part of this chain, as follows: The top left panel of Fig. |
4 shows the star formation rates and stellar masses of the model galaxy sample. | \ref{mock} shows the star formation rates and stellar masses of the model galaxy sample. |
Immediately to the right of this is shown the mapping to absolute rest-frame UV magnitude. | Immediately to the right of this is shown the mapping to absolute rest-frame UV magnitude. |
Whilst the scatter in the relation from SFR toinitial UV emission is worth understanding refUV-SFR)), it is very minimal. | Whilst the scatter in the relation from SFR to UV emission is worth understanding \\ref{UV-SFR}) ), it is very minimal. |
The real problem in any efforts to derive the SFR is the effect of intervening dust on the UV emission. | The real problem in any efforts to derive the SFR is the effect of intervening dust on the UV emission. |
An estimate of this effect is included in the model, after(1999),, by following the radiative transfer of light (at all wavelengths) through dust assumed to be distributed smoothly in the galactic disk. | An estimate of this effect is included in the model, after, by following the radiative transfer of light (at all wavelengths) through dust assumed to be distributed smoothly in the galactic disk. |
Metallicities are included in the calculation, and inclination angles are assigned to each galaxy at random. | Metallicities are included in the calculation, and inclination angles are assigned to each galaxy at random. |
For full details, the reader is directed to and(2010). | For full details, the reader is directed to and. |
. The top right panel shows the correlation between the absolute UV magnitude and star formation rate. | The top right panel shows the correlation between the absolute UV magnitude and star formation rate. |
For comparison, a dashed line refInferredSFR)) shows the relationship that will be assumed when mapping back from the UV to the SFR. | For comparison, a dashed line \\ref{InferredSFR}) ) shows the relationship that will be assumed when mapping back from the UV to the SFR. |
Unsurprisingly, the effect of dust has been both to introduce scatter and to reduce the UV luminosities with respect to this estimate. | Unsurprisingly, the effect of dust has been both to introduce scatter and to reduce the UV luminosities with respect to this estimate. |
The systematic effect of continuing to use this relationship (dashed line) can be appreciated from the remaining panels. | The systematic effect of continuing to use this relationship (dashed line) can be appreciated from the remaining panels. |
The middle right panel of Fig. | The middle right panel of Fig. |
4 shows our model galaxy population in terms of two estimated observables, the UV and visible absolute magnitudes. | \ref{mock} shows our model galaxy population in terms of two estimated observables, the UV and visible absolute magnitudes. |
This is the at which we can turn the process around and analyse the sample to see how well we can recover the physical properties of the population. | This is the at which we can turn the process around and analyse the sample to see how well we can recover the physical properties of the population. |
The choice of observational limits used for this illustrative exercise are taken from(2009),, namely that galaxies are included in the survey if their apparent magnitudes in the rest-frame UV satisfies mzz5« 27. | The choice of observational limits used for this illustrative exercise are taken from, namely that galaxies are included in the survey if their apparent magnitudes in the rest-frame UV satisfies $m_{775}<27$ . |
Those sources that are fainter than m36007:27 are not detected in that filter, but will still be included in the sample | Those sources that are fainter than $m_{3600}\approx27$ are not detected in that filter, but will still be included in the sample |
material. likely associated to dust lanes. is to force the observer to classify the object as al intermediate Sevfert 1 or a Compou-thiu Sevfert 9 | material, likely associated to dust lanes, is to force the observer to classify the object as an intermediate Seyfert 1 or a Compton-thin Seyfert 2. |
To this dual-absorber scenario. basically the sale proposed by ?.. a third maerial should be added. on a scale mich shorter thai the torus. roughly where the DLR is located. | To this dual-absorber scenario, basically the same proposed by \citet{matt00b}, a third material should be added, on a scale much shorter than the torus, roughly where the BLR is located. |
This maeral cannot be seeu iu Conmpton-thick Sevtert 2s. Le. those sources absorved by the torus. Teenπο it Is obscured bv the torus itself. | This material cannot be seen in Compton-thick Seyfert 2s, i.e. those sources absorbed by the torus, because it is obscured by the torus itself. |
It is responsibe for fast variability of the absorbing cohnun cleusity. | It is responsible for fast variability of the absorbing column density. |
From an observational poiut of view. his niaera can be effectively discriminated from the orus if it is patchy. | From an observational point of view, this material can be effectively discriminated from the torus if it is patchy. |
In this case. giveu the close disance to the BID. the chance to see a cloud apALINEo and disappearing along the Lue of sight ix no low. on short timescales. | In this case, given the close distance to the BH, the chance to see a cloud appearing and disappearing along the line of sight is not low, on short timescales. |
In the exceptional case of NGC 1365. a clear case of eclipse POUL a cloud is actually observed (?).. | In the exceptional case of NGC 1365, a clear case of eclipse from a cloud is actually observed \citep{ris07}. |
I£ the cloud is no Comptou-thnick. vou may still observe. Coupton-thin Sevtert 2 wit1 large cobuun deusities (as large as several NMP077" eii7. ofor example). likely. varving. on short timescales. | If the cloud is not Compton-thick, you may still observe Compton-thin Seyfert 2 with large column densities (as large as several $10^{23}$ $^{-2}$, for example), likely varying on short timescales. |
It is inuportau to stress that such a material docs not relax our need for a torus. | It is important to stress that such a material does not relax our need for a torus. |
The latter is needed because nearly aI the observed ACN have a Compton reflection component aud neutral iron Ίνα linc. which do uot siow sieuificant variabiitv up to quite loje timescales. | The latter is needed because nearly all the observed AGN have a Compton reflection component and neutral iron $\alpha$ line, which do not show significant variability up to quite long timescales. |
In order to reproduce this observationa evideace. the material must be Conmpton-thick. with laree covering factoY axl. most of all. quite far your the DII. taleASS he nuclear eniüssion remains constant (but this is iot the case for many ACN). | In order to reproduce this observational evidence, the material must be Compton-thick, with large covering factor and, most of all, quite far from the BH, unless the nuclear emission remains constant (but this is not the case for many AGN). |
Oulv a pe-scale torus las all these characteristics. | Only a pc-scale torus has all these characteristics. |
This scenario (ΜατΊσος in Table 3)) talses a number of predictions. | This scenario (summarized in Table \ref{unifmodel}) ) makes a number of predictions. |
Most of the Coiupton-thick Sevfert 2s are likely still absorbed w the torus and are not expeced o show any flux or spectral change ou timescales lower than vears. | Most of the Compton-thick Seyfert 2s are likely still absorbed by the torus and are not expected to show any flux or spectral change on timescales lower than years. |
ILowever. a fraction of Comptou-thick objecs does uot intercept the torus along the line of sigit. but are caught when a Comptou-thick cloud |ocated in the BLR is passing iu yout of the source | However, a fraction of Compton-thick objects does not intercept the torus along the line of sight, but are caught when a Compton-thick cloud located in the BLR is passing in front of the source. |
Such sources niv change t111 status in a folowing observation. once the cloud has passed. explaine sole of the so-called ‘changine-lools objec ποιο,thenucleus) aud. defiütelv. NGC 1365 (7). | Such sources may change their status in a following observation, once the cloud has passed, explaining some of the so-called `changing-look' objects \citep[e.g.][the alternative being a `switching-off' of the nucleus]{mgm03,gua05,bianchi05c,teng08} and, definitely, NGC 1365 \citep{ris07}. |
. The raction of Comptou-thick sources belonging to he two classes basically desends on the covering ‘actors of the torus aud the Inner absorber. | The fraction of Compton-thick sources belonging to the two classes basically depends on the covering factors of the torus and the inner absorber. |
Most of the C'oniptou-hin Seyfert 2s with colunun censitics of the oider of 1072 7 do iof intercept at all the torus along the ine of sieht. but are absorbe by arge scale dust lanes. which are also respousible or the obscuraion of 1ο optical broad cuison lines. | Most of the Compton-thin Seyfert 2s with column densities of the order of $10^{22}$ $^{-2}$ do not intercept at all the torus along the line of sight, but are absorbed by large scale dust lanes, which are also responsible for the obscuration of the optical broad emission lines. |
On the otos rand. Seyfert 2s with arecr coluun densiles. of re order of-- 107? 5-. are ikelv seen thyroteh he absorbing clouds located at he DLR.iu anaogv to 16 chaneiue-look. oljjects cited above. the oilv aifference being that he iuterveniis clouds are rot Compton-thick. | On the other hand, Seyfert 2s with larger column densities, of the order of $10^{23}$ $^{-2}$, are likely seen through the absorbing clouds located at the BLR, in analogy to the `changing-look' objects cited above, the only difference being that the intervening clouds are not Compton-thick. |
TLOSC sources are probably jo best caididates for imnonitoriung causeus. since they are those with higher xobabilitv of rapid column deusitv variations. | These sources are probably the best candidates for monitoring campaigns, since they are those with higher probability of rapid column density variations. |
NGC 7582 is a clear example of this class. | NGC 7582 is a clear example of this class. |
Again. the fraction of these sources amoung Comstou-thin Sevtert 2s depends ou he covering factor aud the ecouetiy of the imer absorber. | Again, the fraction of these sources among Compton-thin Seyfert 2s depends on the covering factor and the geometry of the inner absorber. |
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