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In the case of SN 2006V, NED lists a Galactic color excess value of E(B—V)mw=0.029 mag et(SchlegelaI.|1998).
In the case of SN 2006V, NED lists a Galactic color excess value of $E(B-V)_{\rm MW} = 0.029$ mag \citep{schlegel98}.
. Close examination of our spectroscopic sequence of this object (Sect. ??))
Close examination of our spectroscopic sequence of this object (Sect. \ref{sec:spectra}) )
shows no evidence for D absorption.
shows no evidence for D absorption.
The presence of D is often considered to be a proxy for dust attenuation, therefore the lack of D combined with the position of SN 2006V in the outskirts of its face-on galaxy suggests minimal host extinction.
The presence of D is often considered to be a proxy for dust attenuation, therefore the lack of D combined with the position of SN 2006V in the outskirts of its face-on galaxy suggests minimal host extinction.
In the following we therefore assume zero host extinction for SN 2006V. Adopting the standard total-to-selective extinction value of Ry=3.1 (Cardellietal][I989) the total color excess in the direction of SN 2006V corresponds to a modest Ay—0.09 mag.
In the following we therefore assume zero host extinction for SN 2006V. Adopting the standard total-to-selective extinction value of $R_V = 3.1$ \citep{cardelli89} the total color excess in the direction of SN 2006V corresponds to a modest $A_{V} = 0.09$ mag.
The Galactic color excess in the direction of SN 2006au is E(B—V)uw=0.172 mag ⋅⋅
The Galactic color excess in the direction of SN 2006au is $E(B-V)_{\rm MW} = 0.172$ mag \citep{schlegel98}. .
Conspicuous D absorption lines at the redshift of UGC 11057 are detected in the first spectrum of SN 2006au.
Conspicuous D absorption lines at the redshift of UGC 11057 are detected in the first spectrum of SN 2006au.
This is shown in the inset of Fig. [5].
This is shown in the inset of Fig. \ref{spectra06au}.
Two Gaussian profiles were fit to the (unresolved) D lines anda total equivalent width of 0.88+0.11 iis measured.
Two Gaussian profiles were fit to the (unresolved) D lines and a total equivalent width of $\pm$ 0.11 is measured.
Using the (2003) correlation between the D equivalent width and host galaxy reddening suggests a host galaxy color excesses E(B—V)nost=0.141 mag.
Using the \citet{turatto03} correlation between the D equivalent width and host galaxy reddening suggests a host galaxy color excesses $E(B-V)_{host}=0.141$ mag.
This value might be affected by large uncertainty and lead to an overestimated host extinction, since we know that the[Turatto etal](2003) relation presents large scatter (see[Poznanskiet al.2011).
This value might be affected by large uncertainty and lead to an overestimated host extinction, since we know that the \citet{turatto03} relation presents large scatter \cite[see][]{poznanski11}.
. However, we combined this value with the Galactic color excess, obtaining E(B—V),5,=0.312 mag, which corresponds to Ay=0.97 mag.
However, we combined this value with the Galactic color excess, obtaining $E(B-V)_{tot}=0.312$ mag, which corresponds to $A_{V}=0.97$ mag.
Armed with estimates of E(B—V)tot, the absolute magnitudes of SNe 2006V and 2006au are computed for each observed bandpass.
Armed with estimates of $E(B-V)_{tot}$, the absolute magnitudes of SNe 2006V and 2006au are computed for each observed bandpass.
Their peak values are given in Table [7].
Their peak values are given in Table \ref{absphotPEAK}.
Compared to SN 1987A, SNe 2006V and 2006au are brighter in all bands.
Compared to SN 1987A, SNe 2006V and 2006au are brighter in all bands.
To demonstrate the differences, in Fig.
To demonstrate the differences, in Fig.
BJ the absolute B-, V- and J-band light curves of SNe 2006au and 2006V are compared to those of SN 1987A. SNe 2006V and 2006au show an absolute B magnitude difference to SN 1987A of ~ 1.5 mag and ~ 1.3 mag respectively.
\ref{absb} the absolute $B$ -, $V$ - and $J$ -band light curves of SNe 2006au and 2006V are compared to those of SN 1987A. SNe 2006V and 2006au show an absolute $B$ magnitude difference to SN 1987A of $\sim$ 1.5 mag and $\sim$ 1.3 mag respectively.
The differences in the V band are smaller ( 1.0 and ~ 0.7 mag) and in the J band are further reduced (~ 0.8 and ~ 0.5 mag).
The differences in the $V$ band are smaller $\sim$ 1.0 and $\sim$ 0.7 mag) and in the $J$ band are further reduced $\sim$ 0.8 and $\sim$ 0.5 mag).
To gain insight into the photospheric temperatures, we plot in Fig.
To gain insight into the photospheric temperatures, we plot in Fig.
the B—V, V—r and J—H color curves of SNe 2006V and 2006au.
\ref{color} the $B-V$ , $V-r$ and $J-H$ color curves of SNe 2006V and 2006au.
[6 The plot also includes the colors of SN 1987A. The photometry of each SN has been corrected for extinction adopting the values previously mentioned.
The plot also includes the colors of SN 1987A. The photometry of each SN has been corrected for extinction adopting the values previously mentioned.
Overall the optical colors of SNe 2006V and 2006au are bluer than for SN 19874, particularly during the earliest epochs where the B—V (top panel) and V—r (middle panel) color differences amount to ~ 0.7 mag and ~ 0.3 mag, respectively.
Overall the optical colors of SNe 2006V and 2006au are bluer than for SN 1987A, particularly during the earliest epochs where the $B-V$ (top panel) and $V-r$ (middle panel) color differences amount to $\sim$ 0.7 mag and $\sim$ 0.3 mag, respectively.
Later, around Pax, the B—V colors of our two objects evolve towards the red, and are at day +35 comparable to the B—V color of SN 1987A. The V—r color of ffollows a similar evolution (at the earliest epoch it is bluer than SN 1987A by ~ 0.3 mag).
Later, around $B_{\rm max}$, the $B-V$ colors of our two objects evolve towards the red, and are at day $+$ 35 comparable to the $B-V$ color of SN 1987A. The $V-r$ color of follows a similar evolution (at the earliest epoch it is bluer than SN 1987A by $\sim$ 0.3 mag).
In contrast, the V—r color of SN 2006V is bluer than SN 19874 at all epochs.
In contrast, the $V-r$ color of SN 2006V is bluer than SN 1987A at all epochs.
The near-infrared color J—H (bottom panel) of SN 2006V is found to be similar to SN 19874 at all epochs, whereas for SN 2006au it is slightly bluer atthe earliest phase (by ~ 0.1 mag).
The near-infrared color $J-H$ (bottom panel) of SN 2006V is found to be similar to SN 1987A at all epochs, whereas for SN 2006au it is slightly bluer atthe earliest phase (by $\sim$ 0.1 mag).
The color comparisons suggest that the photospheres of SNe 2006V and 2006au are at higher temperatures thanfor SN 19874. This is confirmed by black-body fits to the spectral energy distributions (SEDs) obtained from each photometric epoch; seemiddle panel in Fig. [[4].
The color comparisons suggest that the photospheres of SNe 2006V and 2006au are at higher temperatures thanfor SN 1987A. This is confirmed by black-body fits to the spectral energy distributions (SEDs) obtained from each photometric epoch; seemiddle panel in Fig. \ref{bolo}. .
As shown in this plot, the black-body temperatures of SNe 2006V and 2006au are higher
As shown in this plot, the black-body temperatures of SNe 2006V and 2006au are higher
With the use of equations (13)) and (142). this restriction can be expressed as Hence. independent of the detailed physics (1.e.. the value of egIsin @/e.). it is not possible to self-consistently move the emission site to a distance corresponding to the light cylinder.
With the use of equations \ref{eq:1.13}) ) and \ref{eq:1.14}) ), this restriction can be expressed as Hence, independent of the detailed physics (i.e., the value of $\epsilon_{\rm B} \Gamma \sin \theta / \epsilon_{\rm e}$ ), it is not possible to self-consistently move the emission site to a distance corresponding to the light cylinder.
One implicit assumption underlying the results in equations (13) and (14)) is that the pulse profile is independent of frequency.
One implicit assumption underlying the results in equations \ref{eq:1.13}) ) and \ref{eq:1.14}) ) is that the pulse profile is independent of frequency.
Although the minimum pulse structure as well às the Full Width Half Maximum (FWHM) are consistent with being constant in the UBV-range. there Is an indication of a decreasing FWHM with frequency2000).
Although the minimum pulse structure as well as the Full Width Half Maximum (FWHM) are consistent with being constant in the UBV-range, there is an indication of a decreasing FWHM with frequency.
The suggested variation is small enough to be caused by the relativistic streaming itself. sice it induces an anti-correlation between frequency and angular width.
The suggested variation is small enough to be caused by the relativistic streaming itself, since it induces an anti-correlation between frequency and angular width.
Hence. if real. this variation could indicate that the value of Fis not much larger than Fi.
Hence, if real, this variation could indicate that the value of $\Gamma$ is not much larger than $\Gamma_{\rm min}$.
Furthermore. a pulse profile varying with frequency would result in an observed spectrum flatter than the intrinsic one due to overlapping emission regions.
Furthermore, a pulse profile varying with frequency would result in an observed spectrum flatter than the intrinsic one due to overlapping emission regions.
Since observations of the Crab pulsar are consistent with a v! ?-spectrum. this limits the magnitude of possible frequency variations of the pulsar profile in a synchrotror scenario.
Since observations of the Crab pulsar are consistent with a $\nu^{1/3}$ -spectrum, this limits the magnitude of possible frequency variations of the pulsar profile in a synchrotron scenario.
It should also be remembered that the low frequency part of the pulsar spectrum shown in Fig. l..
It should also be remembered that the low frequency part of the pulsar spectrum shown in Fig. \ref{fig:ss09},
and. hence. the value of Vana USed above. 1s from Spitzer-observations.
and, hence, the value of $\nu_{\rm amax}$ used above, is from -observations.
Since these data cannot resolve individual pulses. it could be that this emission is not pulsed.
Since these data cannot resolve individual pulses, it could be that this emission is not pulsed.
Although there are no indications that this should be the case. a conservative approach is therefore to use only the NACO-data. since have shown that the pulse profiles in /HK are similar to the optical ones.
Although there are no indications that this should be the case, a conservative approach is therefore to use only the NACO-data, since have shown that the pulse profiles in $JHK$ are similar to the optical ones.
Taking the K-band observations as the lower limit to the pulsed emission would increase the value of Vanay by a factor = 3.
Taking the K-band observations as the lower limit to the pulsed emission would increase the value of $\nu_{\rm amax}$ by a factor $\approx 3$ .
As can be seen from equations (11)) and (12)). the limits of Roomy and Venorm are then changed only by a factor «2.
As can be seen from equations \ref{eq:1.11}) ) and \ref{eq:1.12}) ), the limits of $R_{\rm norm}$ and $\nu_{\rm B,norm}$ are then changed only by a factor $\approx 2$.
Hence. neither the possible frequency dependence of the pulsar profile nor the unconstrained nature of the emission measured bySpitzer is likely to substantially change the conclusion about the location of the emission region.
Hence, neither the possible frequency dependence of the pulsar profile nor the unconstrained nature of the emission measured by is likely to substantially change the conclusion about the location of the emission region.
The above discussion assumed that »7,. is either the Lorentz factor of a mono-energetic electron distribution or the low energy cut-off in a power-law distribution of electron energies.
The above discussion assumed that $\gamma_{\rm me}'$ is either the Lorentz factor of a mono-energetic electron distribution or the low energy cut-off in a power-law distribution of electron energies.
The reason for this choice is that the observed infrared-optical flux from the Crab pulsar increases with frequency consistent with à v H3.spectrum.
The reason for this choice is that the observed infrared-optical flux from the Crab pulsar increases with frequency consistent with a $\nu^{1/3}$ -spectrum.
However. as discussed in Sect.
However, as discussed in Sect.
?? a flatter spectrum cannot be excluded.
\ref{Obs} a flatter spectrum cannot be excluded.
The brightness temperature at vas, is proportional to yj...
The brightness temperature at $\nu_{\rm abs}$ is proportional to $\gamma_{\rm me}'$.
A spectrum flatter than v* would. therefore. decrease the brightness temperature at vas. since it will be determined by electrons with Lorentz factors y<yj,
A spectrum flatter than $\nu^{1/3}$ would, therefore, decrease the brightness temperature at $\nu_{\rm abs}$, since it will be determined by electrons with Lorentz factors $\gamma < \gamma_{\rm me}'$.
Furthermore. for a given value of peak luminosity (1.8.. £4). the flux below the peak frequency (1.e.. v) would be larger than in the v! case.
Furthermore, for a given value of peak luminosity (i.e., $L_{\rm o}$ ), the flux below the peak frequency (i.e., $\nu_{\rm m}$ ) would be larger than in the $\nu^{1/3}$ case.
Both of these effects result in an increase of the lower limit of Rayon, as well as a decrease of the upper limit of vy.
Both of these effects result in an increase of the lower limit of $R_{\rm norm}$ as well as a decrease of the upper limit of $\nu_{\rm norm}$.
Since a y *-speetrum is the hardest possible for optically thin synchrotron radiation. a spectrum flatter than this would strengthen the above conclusions.
Since a $\nu^{1/3}$ -spectrum is the hardest possible for optically thin synchrotron radiation, a spectrum flatter than this would strengthen the above conclusions.
An alternative site for the pulsed emission ts close to the light cylinder.
An alternative site for the pulsed emission is close to the light cylinder.
The physical setting ts here considerably more uncertain than for emission from within the light cylinder: for example. the energy densities in particles and electric field are expected to be of the same magnitude as that in the magnetic field.
The physical setting is here considerably more uncertain than for emission from within the light cylinder; for example, the energy densities in particles and electric field are expected to be of the same magnitude as that in the magnetic field.
This could affect the emission process as well as the beaming of the radiation in several ways.
This could affect the emission process as well as the beaming of the radiation in several ways.
Both of these effects are crucially dependent on the details of the particle motion.
Both of these effects are crucially dependent on the details of the particle motion.
In order to allow a simplified discussion of these issues. the motion of the particles will be divided into three components: (1) Random motion in the zero-momentum (1.e.. primed) frame in addition to the gyration of the particles around the magnetic field (tang). (
In order to allow a simplified discussion of these issues, the motion of the particles will be divided into three components: (1) Random motion in the zero-momentum (i.e., primed) frame in addition to the gyration of the particles around the magnetic field $u_{\rm rand}$ ). (
2) Streaming or bulk motion as measured in the frame co-rotating with the neutron star (44).
2) Streaming or bulk motion as measured in the frame co-rotating with the neutron star $u_{\rm co}$ ).
In the observer's frame. this component differs from that in Sect.
In the observer's frame, this component differs from that in Sect.
??. mainly due to the importance of aberration effects. which make it likely that the motion is at large angles to the magnetic field. (
\ref{Synch1} mainly due to the importance of aberration effects, which make it likely that the motion is at large angles to the magnetic field. (
3) The motion due to the rotation of the neutron star (44).
3) The motion due to the rotation of the neutron star $u_{\rm rot}$ ).
The origin of thang could be small scale irregularities 1n the magnetic and/or the electric field.
The origin of $u_{\rm rand}$ could be small scale irregularities in the magnetic and/or the electric field.
Such deviations from pure gyration in the large scale magnetic field affect mainly the low frequency emission from an individual particle.
Such deviations from pure gyration in the large scale magnetic field affect mainly the low frequency emission from an individual particle.
An example of such a situation is “jitter” radiation2000)... which corresponds to small scale magnetic irregularities.
An example of such a situation is "jitter" radiation, which corresponds to small scale magnetic irregularities.
As shown by (2000).. the low frequency part of the spectrum is determined by the statistical properties of the magnetic field: in particular. the spectral distribution is expected to differ substantially from the synchrotron case (e.g.. for "jitter radiation ay,xI as compared to «,=1/3 for the synchrotron case).
As shown by , the low frequency part of the spectrum is determined by the statistical properties of the magnetic field; in particular, the spectral distribution is expected to differ substantially from the synchrotron case (e.g., for "jitter" radiation $\alpha_{\nu} \approx 1$ as compared to $\alpha_{\nu} = 1/3$ for the synchrotron case).
Although the statistical properties of possible small scale irregularities close to the light cylinder are hard to predict. it is unlikely that they would result in a spectral index close to that of synchrotron radiation.
Although the statistical properties of possible small scale irregularities close to the light cylinder are hard to predict, it is unlikely that they would result in a spectral index close to that of synchrotron radiation.
Hence. the observation of a spectral index for the Crab pulsar close to αν=1/3 suggests that t does not seriously affect the emission from individual particles and. in particular. that the basic frequency is the cyclotron one.
Hence, the observation of a spectral index for the Crab pulsar close to $\alpha_{\nu} = 1/3$ suggests that $u_{\rm rand}$ does not seriously affect the emission from individual particles and, in particular, that the basic frequency is the cyclotron one.
A situation where d; plays a significant role is similar to the one discussed in Sect. ??..
A situation where $u_{\rm co}$ plays a significant role is similar to the one discussed in Sect. \ref{Synch1}.
The main differences are: (1) The angle 8 does not correspond to the angular distance of the emission site from the magnetic axis: instead. sin4 is now a free parameter determining the value of Κο. (
The main differences are: (1) The angle $\theta$ does not correspond to the angular distance of the emission site from the magnetic axis; instead, $\sin \theta$ is now a free parameter determining the value of $R_{\rm c}$. (
2) The energy density associated with the magnetic field (eg) is not necessarily an invariant even for Lorentz boosts along the magnetic field lines.
2) The energy density associated with the magnetic field $\epsilon_{\rm B}$ ) is not necessarily an invariant even for Lorentz boosts along the magnetic field lines.
Both of these effects are contained in the factor egTIsin6/e; (cf.
Both of these effects are contained in the factor $\epsilon_{\rm B} \Gamma \sin \theta / \epsilon_{\rm e}$ (cf.
eqns |11]] and [12].
eqns \ref{eq:1.11}] ] and \ref{eq:1.12}] ]).
As already emphasized. by suitably varying the value of this factor. the limit for either Room or VBnorm Can be made to come closer to that expected for a source distance corresponding to the light eylinderradius: however. this occurs at the expense of the other limit. which will then correspond to an even larger distance tothe emission region (cf.
As already emphasized, by suitably varying the value of this factor, the limit for either $R_{\rm norm}$ or $\nu_{\rm B,norm}$ can be made to come closer to that expected for a source distance corresponding to the light cylinderradius; however, this occurs at the expense of the other limit, which will then correspond to an even larger distance tothe emission region (cf.
eq. I5].
eq. \ref{eq:1.14a}] ]).
Hence. the physical differences in this case are not expected to affect the main conclusions from Sect. ??..
Hence, the physical differences in this case are not expected to affect the main conclusions from Sect. \ref{Synch1}. .
These points show a clear parabolic trend that we fitted to the equation where the correction to the adopted value of the eclipse time, 675, and to the adopted value of the orbitalperiod, 5P.rpo, and the orbital period derivative, P5, are the fit parameters.
These points show a clear parabolic trend that we fitted to the equation where the correction to the adopted value of the eclipse time, $\delta T^e_0$, and to the adopted value of the orbitalperiod, $\delta P_{\rm orb \; 0}$, and the orbital period derivative, $\dot P_{\rm orb}$, are the fit parameters.
We get a very good fit with a x?/d.o.f.=38.69/25— 1.5.
We get a very good fit with a $\chi^2 / d.o.f. = 38.69 / 25 = 1.5$ .
In agreement with previous results, we find a highly significant derivative of the orbital period, which indicates that the orbital period in this system is increasing at a rate of Py,=(1.499£0.071)x10719 s/s. 'The best-fit values for the orbital parameters, calculated with the corrections we found from the fit of the parabolic trend of the eclipse epochs with Eq. 1,,
In agreement with previous results, we find a highly significant derivative of the orbital period, which indicates that the orbital period in this system is increasing at a rate of $\dot P_{\rm orb} = (1.499 \pm 0.071) \times 10^{-10}$ s/s. The best-fit values for the orbital parameters, calculated with the corrections we found from the fit of the parabolic trend of the eclipse epochs with Eq. \ref{eq1},
are shown in 'Table 2..
are shown in Table \ref{tabps}.
Note that a similar orbital period derivative was recently found with new measures of optical eclipses by "
Note that a similar orbital period derivative was recently found with new measures of optical eclipses by \citet{Bayless_09}.
Apart from mass transfer between the companion and the neutron star, the orbital evolution of this binary system is expected to be driven by the emission of gravitational waves and by magnetic braking.
Apart from mass transfer between the companion and the neutron star, the orbital evolution of this binary system is expected to be driven by the emission of gravitational waves and by magnetic braking.
Under the further assumption of conservative mass transfer, orbital evolution calculations show that the orbital period derivative should be (see ??;; see also ?)), where m and m are the mass of the primary, Μι, and the total mass, M4+M», in units of Mo respectively, m2,9.1 is the mass of the secondary in units of 0.1Mo, P3, is the orbital period in units of 5 h (that is appropriate for because Porn5.57 h), n is the index of the mass-radius relation of the secondary RgοςM3, and where the term Tg~20 takes into account the effect of the magnetic braking.
Under the further assumption of conservative mass transfer, orbital evolution calculations show that the orbital period derivative should be (see \citet{Disalvo_08, Verbunt_93}; see also \citet{Rappaport_87}) ), where $m_1$ and $m$ are the mass of the primary, $M_1$, and the total mass, $M_1 + M_2$ , in units of $M_\odot$ respectively, $m_{2, \; 0.1}$ is the mass of the secondary in units of $0.1\; M_\odot$, $P_{\rm 5h}$ is the orbital period in units of 5 h (that is appropriate for because $P_{\rm orb} = 5.57$ h), $n$ is the index of the mass-radius relation of the secondary $R_2 \propto M_2^{\rm n}$, and where the term ${\rm T_{MB}} \sim 20$ takes into account the effect of the magnetic braking.
In line with ?,, ?,, and ? (see? for a review), we can parametrise this term as where f is a dimensionlessmiomi parameterPB, of order of unity: preferred values are f=0.73 (?) or f=1.78 (?),, and ko.277 is the radius of gyration of the star k in units of 0.277, which is the appropriate value for a 1Meo ZAMS star (?).
In line with \citet{Verbunt_81}, , \citet{Verbunt_93}, and \citet{King_88} (see \citet{Tauris_01} for a review), we can parametrise this term as where $f$ is a dimensionless parameter of order of unity: preferred values are $f = 0.73$ \citep{Skumanich_72} or $f = 1.78$ \citep{Smith_79}, and $k_{0.277}$ is the radius of gyration of the star $k$ in units of $0.277$, which is the appropriate value for a $1 \; M_\odot$ ZAMS star \citep{Claret_89}.
Note that the expression for the MB term given in ? is recovered from the above adopting f=1 and k for a 1Me ZAMS star.
Note that the expression for the MB term given in \citet{Verbunt_93} is recovered from the above adopting $f = 1$ and $k$ for a $1 \; M_\odot$ ZAMS star.
Actually, ? discussed three different expression for the angular momentum losses due to MB, Jyp, namely that proposed by Skumanich (?),, that proposed by ?,, and that proposed by ?,, respectively.
Actually, \citet{Tauris_01} discussed three different expression for the angular momentum losses due to MB, $\dot J_{\rm MB}$, namely that proposed by Skumanich \citep{Verbunt_81}, that proposed by \citet{Stepien_95}, and that proposed by \citet{Rappaport_83}, respectively.
However |Jstepien|<0.1|Jgicumanich| with |JRappaport| somewhat between them.
However $|\dot J_{\rm Stepien}| \le 0.1 |\dot J_{\rm Skumanich}|$ with $|\dot J_{\rm Rappaport}|$ somewhat between them.
Because we found that to describe a quite large Jw is required, we decided to adopt 1822-371|Jgkumanich| which resulted in the term Typ~20 adopted above.
Because we found that to describe a quite large $\dot J_{\rm MB}$ is required, we decided to adopt $|\dot J_{\rm Skumanich}|$ which resulted in the term ${\rm T_{MB}} \sim 20$ adopted above.
'The orbital period derivative we measured cannot be explained by a conservative scenario however.
The orbital period derivative we measured cannot be explained by a conservative scenario however.
A positive orbital period derivative certainly indicates a mass-radius index n«1/3; this is indeed a quite general result, which does not depend on the details of the angular momentum losses (see also Eq.
A positive orbital period derivative certainly indicates a mass-radius index $n < 1/3$; this is indeed a quite general result, which does not depend on the details of the angular momentum losses (see also Eq.
4 below).
\ref{eq:dotM} below).
However, the orbital period derivative we measured, Py,=1.50(7)x10-1? s/s, is about three orders of magnitude larger than what is expected even including the (strongest) MB term!
However, the orbital period derivative we measured, $\dot P_{\rm orb} = 1.50(7) \times 10^{-10}$ s/s, is about three orders of magnitude larger than what is expected even including the (strongest) MB term!
This discrepancy is embarrassingly large suggesting that the conservative evolutionary scenario cannot be applied in this case.
This discrepancy is embarrassingly large suggesting that the conservative evolutionary scenario cannot be applied in this case.
À similar conclusion was reached by ?,, who give an improved ephemeris for this source based onnew optical eclipse measures; these authors also note that an extremely high mass accretion rate onto the neutron star, corresponding to about four times the Eddington limit, would be required to explainthe observed large orbital period derivative, and conclude that much of the
A similar conclusion was reached by \citet{Bayless_09}, , who give an improved ephemeris for this source based onnew optical eclipse measures; these authors also note that an extremely high mass accretion rate onto the neutron star, corresponding to about four times the Eddington limit, would be required to explainthe observed large orbital period derivative, and conclude that much of the
We have reported on a long observation of Ark 120, one of the few known examples of nearly absorption-free active nuclei.
We have reported on a long observation of Ark 120, one of the few known examples of nearly absorption-free active nuclei.
This ‘bare’ BLS1 galaxy is the optimal candidate to investigate the nature of the ubiquitous soft excess observed in the X-ray spectra of unobscured AGN.
This `bare' BLS1 galaxy is the optimal candidate to investigate the nature of the ubiquitous soft excess observed in the X-ray spectra of unobscured AGN.
Its broadband 0.5-40 keV spectrum reveals a significant curvature with respect to the main power-law component, as excess emission is found at both low («2 keV) and high (>10 keV) energies.
Its broadband 0.5–40 keV spectrum reveals a significant curvature with respect to the main power-law component, as excess emission is found at both low $<2$ keV) and high $>10$ keV) energies.
Also, a complex iron K-shell feature is clearly detected at ~5-7 keV, showing a structured profile in which can be identified two (three, tentatively) narrow gaussian cores originating roughly at the BLR scale, and a relativistically skewed line arising from the disc at r=10rg.
Also, a complex iron K-shell feature is clearly detected at $\sim$ 5–7 keV, showing a structured profile in which can be identified two (three, tentatively) narrow gaussian cores originating roughly at the BLR scale, and a relativistically skewed line arising from the disc at $r \ga 10 \ r_\rmn{g}$.
The latter, even if not as prominent as in the most impressive cases, cannot be interpreted through a blend of unresolved lines or a system of multiple absorbers.
The latter, even if not as prominent as in the most impressive cases, cannot be interpreted through a blend of unresolved lines or a system of multiple absorbers.
Hence relativistic effects have to be included in the broadband analysis.
Hence relativistic effects have to be included in the broadband analysis.
We have shown that all the spectral complexity can be successfully explained in terms of a self-consistent reflection model, allowing for both a warm/blurred and a cold/distant reflection component.
We have shown that all the spectral complexity can be successfully explained in terms of a self-consistent reflection model, allowing for both a warm/blurred and a cold/distant reflection component.
In this picture, the reprocessing of the primary X-ray radiation takes place in two distinct regions, that can be plainly identified with the inner accretion disc (rS10rg) and the far-off obscuring torus.
In this picture, the reprocessing of the primary X-ray radiation takes place in two distinct regions, that can be plainly identified with the inner accretion disc $(r \la 10 \ r_\rmn{g})$ and the far-off obscuring torus.
Depending on the exact geometry, however, additional physical locations (e.g. a belt of orbiting clouds) may contribute, since fluorescent iron emission suggests the existence of a wide range of ionization states.
Depending on the exact geometry, however, additional physical locations (e.g. a belt of orbiting clouds) may contribute, since fluorescent iron emission suggests the existence of a wide range of ionization states.