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This blurred reflection scenario is further corroborated after reviewing a previous high-quality observation of Ark 120.
This blurred reflection scenario is further corroborated after reviewing a previous high-quality observation of Ark 120.
Therefore, as for this single source, the present interpretation stands as the most convincing among those that have been proposed so far, due to the minimal set of geometrical and physical assumptions involved.
Therefore, as for this single source, the present interpretation stands as the most convincing among those that have been proposed so far, due to the minimal set of geometrical and physical assumptions involved.
In general it is difficult to rely on comparably clean lines of sight to the nuclear regions, and we expect that other processes (likely non-thermal Comptonization, warm absorption, reprocessing in winds/outflows) also play a non-negligible role in shaping the X-ray emission of type 1 AGN.
In general it is difficult to rely on comparably clean lines of sight to the nuclear regions, and we expect that other processes (likely non-thermal Comptonization, warm absorption, reprocessing in winds/outflows) also play a non-negligible role in shaping the X-ray emission of type 1 AGN.
As a consequence, our aim in the near future is to extend this study to a large sample of low-obscuration sources, possibly observed by both and to combine the high-energy coverage of the former and the large effective area of the latter.
As a consequence, our aim in the near future is to extend this study to a large sample of low-obscuration sources, possibly observed by both and to combine the high-energy coverage of the former and the large effective area of the latter.
This will allow us to infer the average contribution of blurred disc reflection to the soft excess and to the X-ray luminosity of AGN, and at the same time to test the basic predictions of the light bending model and to explore the nature of the illuminating source itself.
This will allow us to infer the average contribution of blurred disc reflection to the soft excess and to the X-ray luminosity of AGN, and at the same time to test the basic predictions of the light bending model and to explore the nature of the illuminating source itself.
We note that if blurred reflection is really the dominant component of the soft excess, the black hole spin distribution may be biased towards the larger values; indeed, it has already been suggested in many previous works (e.g. Volonteri et al.
We note that if blurred reflection is really the dominant component of the soft excess, the black hole spin distribution may be biased towards the larger values; indeed, it has already been suggested in many previous works (e.g. Volonteri et al.
2005) that the supermassive black holes in AGN should be rapidly rotating.
2005) that the supermassive black holes in AGN should be rapidly rotating.
Even if the spin cannot be firmly constrained in Ark 120 due to the large degeneracy among the blurring parameters (but provisionally a~0.7), our case study is apparently consistent with this scenario.
Even if the spin cannot be firmly constrained in Ark 120 due to the large degeneracy among the blurring parameters (but provisionally $a \simeq 0.7$ ), our case study is apparently consistent with this scenario.
EN acknowledges financial support from the ASI-INAF 1/088/06/0 contract.
EN acknowledges financial support from the ASI-INAF I/088/06/0 contract.
ACF thanks the Royal Society.
ACF thanks the Royal Society.
RCR and DJW acknowledge the financial support provided by STFC.
RCR and DJW acknowledge the financial support provided by STFC.
The authors are also grateful to the anonymous referee for the constructive comments.
The authors are also grateful to the anonymous referee for the constructive comments.
the accelerated particles.
the accelerated particles.
The energy distribution of the accelerated particles determines then yer=1+Por/Ucr; where P. is the cosmic-ray pressure and Ue; the cosmic-ray energy density.
The energy distribution of the accelerated particles determines then $\gamma_{\rm cr}=1 + P_{\rm cr}/u_{\rm cr}$, where $P_{\rm cr}$ is the cosmic-ray pressure and $u_{\rm cr}$ the cosmic-ray energy density.
Hence the kinematic models predict that the e;—w diagram should deviate from ours with Yer= and should lie between our curves of yer=4/3 and 5/3.
Hence the kinematic models predict that the $\epsilon_{\rm esc}-w$ diagram should deviate from ours with $\gamma_{\rm cr}=4/3$, and should lie between our curves of $\gamma_{\rm cr}=4/3$ and $5/3$.
4/3,In practice, however, the cosmic-ray adiabatic index as obtainedby the kinetic models is very close to Yer=4/3 for pas=10°mce, as can be seen in Fig. 6..
In practice, however, the cosmic-ray adiabatic index as obtainedby the kinetic models is very close to $\gamma_{\rm cr}=4/3$ for $p_{max}= 10^5mc$, as can be seen in Fig. \ref{fig:blasi}.
For Pmaz=10?mc we found that a modest increase to Yer=1.4 resulted in a good match between the two-fluid solutions and the kinetic model of Blasietal.(2005).
For $p_{max}= 10^2mc$ we found that a modest increase to $\gamma_{\rm cr}=1.4$ resulted in a good match between the two-fluid solutions and the kinetic model of \citet{blasi05}.
. Interestingly, the two-fluid solutions presented here indicate that a higher energy flux escape is necessary for Yer>4/3, but from the point of view of the spectral energy distribution the energy escape flux is more difficult to achieve for spectra with spectral energy indices T>2 (N(E)« corresponding to 4/3.
Interestingly, the two-fluid solutions presented here indicate that a higher energy flux escape is necessary for $\gamma_{\rm cr}>4/3$, but from the point of view of the spectral energy distribution the energy escape flux is more difficult to achieve for spectra with spectral energy indices $\Gamma > 2$ $N(E)\propto E^{-\Gamma})$, corresponding to $\gamma_{\rm cr}>4/3$ .
The reason is that for ΕΤ),T>2 most of the cosmic- energy is contained by mildly relativistic particles, whereas most of the escaping particles will be near the maximum of the energy distribution in the Blasi model).
The reason is that for $\Gamma > 2$ most of the cosmic-ray energy is contained by mildly relativistic particles, whereas most of the escaping particles will be near the maximum of the energy distribution $p_{max}$ in the Blasi model).
Hence, the escaping particles (pmaxcarry away only a small fraction of the internal energy.
Hence, the escaping particles carry away only a small fraction of the internal energy.
For l'«2 most energy is indeed concentrated around pmax and the necessary escape flux is easily generated.
For $\Gamma < 2$ most energy is indeed concentrated around $p_{max}$ and the necessary escape flux is easily generated.
So from a thermodynamic point of view, a high fractional cosmic-ray pressure w, requires large values of €esc, Which in turn is best achieved with yo,24/3.
So from a thermodynamic point of view, a high fractional cosmic-ray pressure $w$, requires large values of $\epsilon_{\rm esc}$, which in turn is best achieved with $\gamma_{\rm cr} \approx 4/3$.
One of the motivations for this work has been the measurements of the downstream temperature of SNRs with reasonably well known velocities (Helderetal.2009,
One of the motivations for this work has been the measurements of the downstream temperature of SNRs with reasonably well known velocities \citep{helder09,helder10}.
These temperatures were measured using the broad 2010)..component of the Ha line.
These temperatures were measured using the broad component of the $\alpha$ line.
This component is due to charge exchange between neutral hydrogen atoms entering the shock and the downstream population of shock heated protons.
This component is due to charge exchange between neutral hydrogen atoms entering the shock and the downstream population of shock heated protons.
Hence, the width of the line is caused by thermal Doppler broadening and reflects the downstream proton temperature.
Hence, the width of the line is caused by thermal Doppler broadening and reflects the downstream proton temperature.
There is some ambiguity as to how to relate the proton temperature to the overall downstream plasma temperature, as the different plasma constituents (electrons, protons, helium, other ions) may not be in thermal equilibrium.
There is some ambiguity as to how to relate the proton temperature to the overall downstream plasma temperature, as the different plasma constituents (electrons, protons, helium, other ions) may not be in thermal equilibrium.
However, the proton temperature is always expected to be within a factor of 2 of the mean plasma temperature.
However, the proton temperature is always expected to be within a factor of 2 of the mean plasma temperature.
Using Equation (20)) (see also Fig. 3))
Using Equation \ref{eq:beta}) ) (see also Fig. \ref{fig:beta}) )
one can easily estimate from a measured temperatureand shock velocity what the downstream fractional cosmic-ray pressure is and what the required escape flux is.
one can easily estimate from a measured temperatureand shock velocity what the downstream fractional cosmic-ray pressure is and what the required escape flux is.
The only ambiguity that is left for estimating ees. is what the effective cosmic-ray adiabatic index is.
The only ambiguity that is left for estimating $\epsilon_{\rm esc}$ is what the effective cosmic-ray adiabatic index is.
However, γα is not important for determining w.
However, $\gamma_{\rm cr}$ is not important for determining $w$.
For the SNRs under consideration the shock velocities are well in excess of 1000kms~!.
For the SNRs under consideration the shock velocities are well in excess of 1000.
. For a typical sound speed in the interstellar medium of 10kms~!,, we have Mo>100.
For a typical sound speed in the interstellar medium of 10, we have $M_0 > 100$.
This means that the high Mach number approximation is valid as long as w€0.9.
This means that the high Mach number approximation is valid as long as $w \lesssim 0.9$.
This appears to be the case for the SNRs considered below.
This appears to be the case for the SNRs considered below.
For the northeastern region of the TeV gamma-ray emitting remnant RCW 86 Helderetal.(2009) measured a downstream temperature of 0.3 keV for a measured shock velocity of kgT,V,=2800 kms-!.
For the northeastern region of the TeV gamma-ray emitting remnant RCW 86 \citet{helder09} measured a downstream temperature of $k_{\rm B}T_{\rm p}= 2.3\pm 0.3$ keV for a measured shock velocity of $V_s=6000 \pm 2800$ .
Nominally this corresponds to 6=0.055 Equation (20))), but given the systematic uncertainties(see and the ambiguity due to non-equilibration of temperatures 8 could be as high 8=0.31.
Nominally this corresponds to $\beta = 0.055$ (see Equation \ref{eq:beta}) )), but given the systematic uncertainties and the ambiguity due to non-equilibration of temperatures $\beta$ could be as high $\beta = 0.31$.
Using now Equations (9)) and (15)) in the limit for high Mach numbers the measured values of 8 correspond to downstream fractional cosmic-ray pressures and escape fractions of w=0.81,e&,0.59 and w=0.51,€ese 0.20, for y= and 8=0.055 and B=0.31 respectively.
Using now Equations \ref{eq:w}) ) and \ref{eq:epsilon}) ) in the limit for high Mach numbers the measured values of $\beta$ correspond to downstream fractional cosmic-ray pressures and escape fractions of $w=0.81, \epsilon_{\rm esc}=0.59$ and $w=0.51, \epsilon_{\rm esc}=0.20$ , for $\gamma_{\rm cr}=4/3$ and $\beta=0.055$ and $\beta =0.31$ respectively.
For 4/3yer=5/3 this is e.=0.72 and €esc= 0.38, respectively, with w unchanged.
For $\gamma_{\rm cr}=5/3$ this is $\epsilon_{\rm esc}=0.72$ and $\epsilon_{\rm esc}=0.38$ , respectively, with $w$ unchanged.
For the young Large Magellanic Cloud remnant 0509- Helderetal. determined the post-shock temperatures in two (2010)regions.
For the young Large Magellanic Cloud remnant 0509-67.5 \citet{helder10} determined the post-shock temperatures in two regions.
The most constraining measurement was for the southwestern region for which kpTp=28.7 keV, for a conservative velocity estimate of V,=5000 kms-!, corresponding to 8= 0.58.
The most constraining measurement was for the southwestern region for which $k_{\rm B}T_{\rm p}= 28.7$ keV, for a conservative velocity estimate of $V_s=5000$ , corresponding to $\beta=0.58$ .
curvature signal.
curvature signal.
With increased Jeans smoothing. more of the high-curvature regions for the same wowill instead be small undulations on top of broader transmission regions.
With increased Jeans smoothing, more of the high-curvature regions for the same will instead be small undulations on top of broader transmission regions.
Such undulations are more easily obscured than distinct peaks by smoothing or adding noise.
Such undulations are more easily obscured than distinct peaks by smoothing or adding noise.
The curvature at 275 in the T15 runs is therefore lower than in run BIS when measured from. present-quality spectra. and so the temperature in the test runs is over-estimated.
The curvature at $z \sim 5$ in the T15 runs is therefore lower than in run B15 when measured from present-quality spectra, and so the temperature in the test runs is over-estimated.
The impact of thermal history. is more straightforward abo~3.
The impact of thermal history is more straightforward at $z \sim 3$.
Here. changes in the temperature measurement due to noise and. lowered resolution are minimal.
Here, changes in the temperature measurement due to noise and lowered resolution are minimal.
The T15 runs and run DI15 have similar values ofZ5.. but. because the gas in the ΓΕ runs was colder than in run D15 at higher redshifts it has experienced. less Jeans smoothing.
The T15 runs and run D15 have similar values of, but because the gas in the T15 runs was colder than in run D15 at higher redshifts it has experienced less Jeans smoothing.
The curvature in the T15 runs at z~3 is therefore higher bv comparison. which leads to an underestimate of the Lemiperature.
The curvature in the T15 runs at $z \sim 3$ is therefore higher by comparison, which leads to an underestimate of the temperature.
Although the instantaneous temperature is exactly recovered at certain redshifts. this does necessarily mean that the gas has returned to hyclrostatic equilibrium at these points.
Although the instantaneous temperature is exactly recovered at certain redshifts, this does necessarily mean that the gas has returned to hydrostatic equilibrium at these points.
“Lhe florest absorbers at a given density will tend to have physical sizes on order of the local Jeans length. La~ὃςfava. where eis the sound speed. and faa is the dynamical timescale (?)..
The forest absorbers at a given density will tend to have physical sizes on order of the local Jeans length, $L_{\rm J} \sim c_{\rm s} \, t_{\rm dyn}$, where $c_{\rm s}$ is the sound speed and $t_{\rm dyn}$ is the dynamical timescale \citep{schaye2001}.
Following a heating event. therefore. local hydrostatic equilibrium will only be restored on the dynamical timescale. which at the mean density is on the order of a Hubble time: Lea—lptyat
Following a heating event, therefore, local hydrostatic equilibrium will only be restored on the dynamical timescale, which at the mean density is on the order of a Hubble time: $t_{\rm dyn} \equiv 1 / \sqrt{G \rho} \sim t_{\rm H} \, \Delta^{-1}$.
Evenat 2—2.3. where the forest is probing mildly overdense regions. the thermal response time for tvpical absorbers will be significantly longer than the timescales over which hydrogen and helium reionization occur.
Even at $z = 2-3$, where the forest is probing mildly overdense regions, the thermal response time for typical absorbers will be significantly longer than the timescales over which hydrogen and helium reionization occur.
In that case. the small-scale structure of the florest is sensitive to the thermal history of the gas.
In that case, the small-scale structure of the forest is sensitive to the thermal history of the gas.
The redshifts where we exactly recover the temperatures in the test simulations are simply those where the integrated
The redshifts where we exactly recover the temperatures in the test simulations are simply those where the integrated
If the first scenario is ruled out by more precise observations, we are left with the third scenario (Section 5)): planet b created the warp and then had its inclination damped.
If the first scenario is ruled out by more precise observations, we are left with the third scenario (Section \ref{sec:damp}) ): planet b created the warp and then had its inclination damped.
Detailed modeling of scenarios that allow for the damping of planet b's inclination will be necessary.
Detailed modeling of scenarios that allow for the damping of planet b's inclination will be necessary.
Confirmation of the damping scenario, especially if observers discover more systems with planets misaligned with the warp they produced, could shed light on disk properties that are important for planet formation but difficult to measure directly.
Confirmation of the damping scenario, especially if observers discover more systems with planets misaligned with the warp they produced, could shed light on disk properties that are important for planet formation but difficult to measure directly.
We thank Michael Fitzgerald, Paul Kalas, and Philippe Thebault for 8 Pictoris insights, and Thayne Currie and an anonymous referee for helpful comments.
We thank Michael Fitzgerald, Paul Kalas, and Philippe Thebault for $\beta$ Pictoris insights, and Thayne Currie and an anonymous referee for helpful comments.
R.I.D. acknowledges support by NSF Graduate Research Fellowship DGE-1144152 and D.C.F. by NASA Hubble Fellowship HF-51272.01.
R.I.D. acknowledges support by NSF Graduate Research Fellowship DGE-1144152 and D.C.F. by NASA Hubble Fellowship HF-51272.01.
Simulations were run on the Odyssey cluster supported by the Harvard FAS Sciences Division Research Computing Group.
Simulations were run on the Odyssey cluster supported by the Harvard FAS Sciences Division Research Computing Group.
interior mass would expand to re-engulf the companion and restart in-spiral.
interior mass would expand to re-engulf the companion and restart in-spiral.
Note that σα. (17))
Note that Eq. \ref{Delta_E_orb}) )
depends on the stellar structure at the onset of tidal engulfment.
depends on the stellar structure at the onset of tidal engulfment.
To calculate a minimum period gap expected. for a given παν system. we assume a=1.
To calculate a minimum period gap expected for a given binary system, we assume $\alpha = 1$.
This gives an upper round on the orbital radius at which we would expect to ind companions which have survived a common envelope yhase (see Lig. 1))
This gives an upper bound on the orbital radius at which we would expect to find companions which have survived a common envelope phase (see Fig. \ref{thegap}) ).
In 844. we determined the time at which a companion of mass AZ, is engulfed by the giant star.
In 4, we determined the time at which a companion of mass $M_c$ is engulfed by the giant star.
Couplecl with the structure of the star at the time of idal engulfment. we can determine which semimajor axis « satisfies Eq. (14)).
Coupled with the structure of the star at the time of tidal engulfment, we can determine which semimajor axis $a$ satisfies Eq. \ref{CE}) ).
a£ the companion avoids tidal disruption. hen it has successfully ejected the envelope and. survived he CL phase.
If the companion avoids tidal disruption, then it has successfully ejected the envelope and survived the CE phase.
The tical shredding radius can be estimated. by xdancing the dillercntial gravitational force across the companion with its self gravity.
The tidal shredding radius can be estimated by balancing the differential gravitational force across the companion with its self gravity.
This vields a tidal shredcding radius given by ας2I72MM, where 2. is the radius of the companion.
This yields a tidal shredding radius given by $a_s \simeq R_c \sqrt[3]{2M / M_c}$ where $R_c$ is the radius of the companion.
For a ΕΛ companion around a proto-white dwarf core. v,~7.410!" em.
For a $1M_{\rm J}$ companion around a proto-white dwarf core, $a_s \sim 7.4\times 10^{10}$ cm.
For a lO Aj a,23107 em with the actual values dependent on the degenerate. core mass during the CLP.
For a 10 $M_{\rm J}$ $a_s\gtrsim 3\times 10^{10}$ cm with the actual values dependent on the degenerate core mass during the CEP.
Ifthe companion unbinds the envelope exterior to ας. then we savit has survived. common envelope evolution.
If the companion unbinds the envelope exterior to $a_s$ , then we sayit has survived common envelope evolution.
Note that the factor of 2A//M, in the expression for e, neglects the synchronous rotation of the companion and its finite Love number.
Note that the factor of $2M/M_c$ in the expression for $a_s$ neglects the synchronous rotation of the companion and its finite Love number.
Including these ellects beads to a slightly larger tidal shredding radius. implving that slightly. more massive companions are required to unbind the CIS and avoid tidal disruption.
Including these effects leads to a slightly larger tidal shredding radius, implying that slightly more massive companions are required to unbind the CE and avoid tidal disruption.
Since the stellar mass function is heavily weighted towards lower masses. we limit ourselves to a 1 AM. progenitor.
Since the stellar mass function is heavily weighted towards lower masses, we limit ourselves to a 1 $M_\odot$ progenitor.
More massive primaries extend to larger raclii (see Fig. 3))
More massive primaries extend to larger radii (see Fig. \ref{models}) )
and swallow companions at farther distances. leading to wider period gaps.
and swallow companions at farther distances, leading to wider period gaps.
We consider two companions: al Aly planet and a 10 AZ; brown clhwact.
We consider two companions: a 1 $M_{\rm J}$ planet and a 10 $M_{\rm J}$ brown dwarf.
For each evolutionary model and binary configuration. we calculate the minimum and maximum. bounds of the separation gap.
For each evolutionary model and binary configuration, we calculate the minimum and maximum bounds of the separation gap.
Our results for η=0.7. 1 and 5. or=0 and 1 and Qu=10" and f=1 are summarized in Table 2..
Our results for $\eta = 0.7$, $1$ and $5$, $x=0$ and $1$ and $Q_0 = 10^6$ and $f=1$ are summarized in Table \ref{table2}.
The maximum of μμ and minimum of es vield the minimum gap expected for a M,=LM... M,=LAL) system.
The maximum of $a_{\rm min}$ and minimum of $a_{\rm max}$ yield the minimum gap expected for a $M_{\star} = 1 M_\odot$, $M_{\rm c}=1 M_{\rm J}$ system.
Note that for pedagogical purposes we include the results for mass-loss only (no tides) in Table 2..
Note that for pedagogical purposes we include the results for mass-loss only (no tides) in Table \ref{table2}.
However. when determining minimum period gaps. we only consider systems in which tides are acting.
However, when determining minimum period gaps, we only consider systems in which tides are acting.
H0 should. be stressed. that. even in the absence of tides. à minimum period gap exists.
It should be stressed that even in the absence of tides, a minimum period gap exists.
The functional dependence of tides on separation (Iq.
The functional dependence of tides on separation (Eq.
2) largely acts to change the location of the outer boundary. of the gap.
2) largely acts to change the location of the outer boundary of the gap.
Depending on tidal prescription. this shifts the outer boundary of the by a factor of —2-3 at most Clable 2)).
Depending on tidal prescription, this shifts the outer boundary of the by a factor of $\sim$ 2-3 at most (Table \ref{table2}) ).
From Table 2.. we see that there should be a paucity of Jupiter-mass companions with periods S270 days around white cwarls.
From Table \ref{table2}, we see that there should be a paucity of Jupiter-mass companions with periods $\lesssim$ 270 days around white dwarfs.
Actditionally from Table 2.. we see that there should. be a paucity of LO AZ; companions with periods between 0.1. ανν (0.003 AW) and. 380 days. (0.75 AU).
Additionally from Table \ref{table2}, , we see that there should be a paucity of 10 $M_{\rm J}$ companions with periods between 0.1 days (0.003 AU) and 380 days (0.75 AU).
‘This is consistent with the tentative detection of a —2 Mj planet in à Z4 wear (22.75 AU) orbit around. the white dwarl CD 66 (22)..
This is consistent with the tentative detection of a $\sim$ 2 $M_{\rm J}$ planet in a $\gtrsim$ 4 year $\gtrsim$ 2.75 AU) orbit around the white dwarf GD 66 \citep{Mullally:2008fk, Mullally:2009uq}.
Future surveys for low-mass companions around white chvarls will have the ability to confirm or refute our prediction of a gap.
Future surveys for low-mass companions around white dwarfs will have the ability to confirm or refute our prediction of a gap.
Several ellorts are either recently completed or are currently. underway (22777)...
Several efforts are either recently completed or are currently underway \citep{Farihi:2006vn, Tremblay:2007fr, Hoard:2007ys, Farihi:2008rt, Farihi:2009zr}.
Once the samples are suflicicnthy large. observational identification of period gaps could help to constrain aspects of mass-LIoss and tidal theories.
Once the samples are sufficiently large, observational identification of period gaps could help to constrain aspects of mass-loss and tidal theories.
For each Af,=LAL. model. we can calculate the münimunm mass companion able to unbind the envelope and survive a CI. phase.
For each $M_\star = 1 M_\odot$ model, we can calculate the minimum mass companion able to unbind the envelope and survive a CE phase.
This occurs when the binding energy of envelope is at à minimum.
This occurs when the binding energy of envelope is at a minimum.
Note that. in multiple planet systems. several close companions may incur a CLE phase.
Note that, in multiple planet systems, several close companions may incur a CE phase.
In conjunction. multiple lower-mass companions could potentially supply the same energy as a single larger mass object.
In conjunction, multiple lower-mass companions could potentially supply the same energy as a single larger mass object.