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Nevertheless. the fact that large portions of these mass functions have been directly measured allows one to construct useful constraints.
Nevertheless, the fact that large portions of these mass functions have been directly measured allows one to construct useful constraints.
We [formulate these constraints within the context of a specific model of the Galactic bulee and disk.
We formulate these constraints within the context of a specific model of the Galactic bulge and disk.
The model of the disk is relatively secure. and uncertainties in it play an overall very small role in controversies about the total optical depth toward the bulge.
The model of the disk is relatively secure, and uncertainties in it play an overall very small role in controversies about the total optical depth toward the bulge.
That is whv most of the effort to explain the high optical depth has centered. on models of the bulee.
That is why most of the effort to explain the high optical depth has centered on models of the bulge.
The benchmark model (hat we adopt is therefore [ar Irom unique.
The benchmark model that we adopt is therefore far from unique.
We show. however. that it is possible to factor the stellar constraint into (wo terms. one representing the bulge model convolved with the observational strategy. ancl the other representing the star counts.
We show, however, that it is possible to factor the stellar constraint into two terms, one representing the bulge model convolved with the observational strategy, and the other representing the star counts.
In this way. our result can easily be applied to any model of the Galactic bulge.
In this way, our result can easily be applied to any model of the Galactic bulge.
We model the local vertical disk density profile in accord with the modelof (2002)..
We model the local vertical disk density profile in accord with the modelof \citet{zheng}.
To extend this model to (he whole Galactic disk. we assume that the column density of the disk has a scale length //=2.75 kpc. as measured by Zhengοἱal.(2002).
To extend this model to the whole Galactic disk, we assume that the column density of the disk has a scale length $H=2.75\,\kpc$ , as measured by \citet{zheng}.
. We account for the gradual flaring of the disk by rescalingthe scale heights in the [formula inproportion to the scale height derived by Kent(1992)..
We account for the gradual flaring of the disk by rescalingthe scale heights in the \citet{zheng} formula inproportion to the scale height derived by \citet{kent}.
We normalize the local stellar column density to X=36M.pe 7.
We normalize the local stellar column density to $\Sigma_0 = 36\,M_\odot\,\pc^{-2}$ .
This includes 28.42.pe7 in observable stars and white dwarls (Zhengetal.2002:Gould.Baheall.&Flynn1997). and another 8AL.pe7. which is a rough estimate of the column density of brown dwarls(BDs).
This includes $28\,M_\odot\,\pc^{-2}$ in observable stars and white dwarfs \citep{zheng,gbf} and another $8\,M_\odot\,\pc^{-2}$, which is a rough estimate of the column density of brown dwarfs(BDs).
Thedisk densityprofile in cxlindrical coordinates is (hen. where po=0.0493M.pe 7. 3= 0.565. hy= 2T0pe. hy=HOpe. 1=2.55kpe. lij =Skpec. and
Thedisk densityprofile in cylindrical coordinates is then, where $\rho_0=0.0493\,M_\odot\,\pc^{-3}$ , $\beta = 0.565$ , $h_1=270\,\pc$ , $h_2=440\,\pc$, $H=2.75\,\kpc$, $R_0=8\,\kpc$ , and
are to be reached.
are to be reached.
The chief mechanism which may be responsible for such fields is the excitation of streaming instabilities (SD by the same particles which are being accelerated (????)).
The chief mechanism which may be responsible for such fields is the excitation of streaming instabilities (SI) by the same particles which are being accelerated \cite{skilling75c,bell78a,lc83a,lc83b}) ).
The effect of magnetic field amplification on the maximum energy reachable at supernova remnant (SNR) shocks was investigated by ??., who reached the conclusion that cosmic rays could be accelerated up to energies of order ~10—10° GeV at the beginning of the Sedov phase.
The effect of magnetic field amplification on the maximum energy reachable at supernova remnant (SNR) shocks was investigated by \cite{lc83a,lc83b}, who reached the conclusion that cosmic rays could be accelerated up to energies of order $\sim 10^4-10^5$ GeV at the beginning of the Sedov phase.
This conclusion was primarily based on the assumption of Bohm diffusion and a saturation level for the induced turbulent Ποιά 6B/B~1.
This conclusion was primarily based on the assumption of Bohm diffusion and a saturation level for the induced turbulent field $\delta B/B\sim 1$.
On the other hand, recent observations of the X-ray surface brightness of the rims of SNRs have shown that 6B/B~100—1000 (see ? for a review of results), thereby renewing the interest in the mechanism of magnetic field amplification and in establishing its saturation level.
On the other hand, recent observations of the X-ray surface brightness of the rims of SNRs have shown that $\delta B/B\sim 100-1000$ (see \cite{volk} for a review of results), thereby renewing the interest in the mechanism of magnetic field amplification and in establishing its saturation level.
It is however worth to recall that the interpretation of the X-ray observations 1s not yet unique: the narrow rims observed in the X-ray synchrotron emission could be due to the damping of the downstream magnetic field ? rather than to severe synchrotron losses of very high energy electrons, although this interpretation has some serious shortcomings (see ? for à discussion).
It is however worth to recall that the interpretation of the X-ray observations is not yet unique: the narrow rims observed in the X-ray synchrotron emission could be due to the damping of the downstream magnetic field \cite{pohl} rather than to severe synchrotron losses of very high energy electrons, although this interpretation has some serious shortcomings (see \cite{morlino} for a discussion).
1n this context of excitement, due to the implications of these discoveries for the origin of cosmic rays, ? discussed the excitation of modes in a plasma treated in the MHD approximation and found that a new, purely growing, non-alfvénnic mode appears for high acceleration efficiencies.
In this context of excitement, due to the implications of these discoveries for the origin of cosmic rays, \cite{bell04} discussed the excitation of modes in a plasma treated in the MHD approximation and found that a new, purely growing, non-alfvénnic mode appears for high acceleration efficiencies.
The author predicted a saturation of this SI at the level 6B/B~M4Gyv,/0) where 77 is the cosmic ray pressure in units of the kinetic pressure pvr. v, is the shock speed and M4=v/v is the Alfvénnic Mach number.
The author predicted a saturation of this SI at the level $\delta B/B\sim M_A (\eta v_s/c)^{1/2}$ where $\eta$ is the cosmic ray pressure in units of the kinetic pressure $\rho v_s^2$, $v_s$ is the shock speed and $M_A=v_s/v_A$ is the Alfvénnic Mach number.
For comparison, standard SI for resonant wave-particle interactions leads to expect 05/B~M,j!
For comparison, standard SI for resonant wave-particle interactions leads to expect $\delta B/B\sim M_A^{1/2} \eta^{1/2}$.
For efficient acceleration 7)~1. and typically. for shocks in the interstellar medium, M,~107.
For efficient acceleration $\eta\sim 1$, and typically, for shocks in the interstellar medium, $M_A\sim 10^4$.
Therefore Bell's mode leads to 6B/B~300—1000 while the standard SI gives 6B/B~30.
Therefore Bell's mode leads to $\delta B/B\sim 300-1000$ while the standard SI gives $\delta B/B\sim 30$.
It is also useful to notice that the saturation level predicted by ? is basically independent of the value of the background field, since 6B7/87~(1/2)(,/c)Peg. where Pep 1s the cosmic ray pressure at the shock surface.
It is also useful to notice that the saturation level predicted by \cite{bell04} is basically independent of the value of the background field, since $\delta B^2/8\pi\sim (1/2)(v_s/c)P_{CR}$, where $P_{CR}$ is the cosmic ray pressure at the shock surface.
The resonant and non-resonant mode have different properties in other respects as well.
The resonant and non-resonant mode have different properties in other respects as well.
A key feature consists in the different wavelengths that are excited.
A key feature consists in the different wavelengths that are excited.
The resonant mode with the maximum growth rate has wavenumber & such that Απο=1. where Γεω is the Larmor radius of the particles that dominate the cosmic ray number density at the shock, namely. for typical spectra of astrophysical interest, the lowest energy cosmic rays at the shock.
The resonant mode with the maximum growth rate has wavenumber $k$ such that $\ko=1$, where $r_{L,0}$ is the Larmor radius of the particles that dominate the cosmic ray number density at the shock, namely, for typical spectra of astrophysical interest, the lowest energy cosmic rays at the shock.
At the shock location. the minimum momentum is the injection momentum, while at larger distances from the shock, the minimum momentum is determined by the diffusion properties upstream and is higher than the injection momentum, since higher energy particles diffuse farther upstream.
At the shock location, the minimum momentum is the injection momentum, while at larger distances from the shock, the minimum momentum is determined by the diffusion properties upstream and is higher than the injection momentum, since higher energy particles diffuse farther upstream.
When the
When the
mmap was made towards SAIN. with a grid spacing of10".
map was made towards SMM4, with a grid spacing of.
. The mapping was performed using the on-(he-lly mapping capability al the JCAIT.
The mapping was performed using the on-the-fly mapping capability at the JCMT.
Position switched observations were also made in the (transition towards (he central position of $MMA.
Position switched observations were also made in the transition towards the central position of SMM4.
The JCMT observations were done using the facility À31 single channel SIS receiver.
The JCMT observations were done using the facility A3i single channel SIS receiver.
The spectrometer backend used was the Dutch Autocorrelation Spectrometer (DAS) configured to operate with an effective resolution of 95 kllz and total bandwidth of 125 Mz.
The spectrometer backend used was the Dutch Autocorrelation Spectrometer (DAS) configured to operate with an effective resolution of 95 kHz and total bandwidth of 125 MHz.
In Figure 1 we present a plot of our spectral observations towards (he central position of SMALL.
In Figure \ref{plotall} we present a plot of our spectral observations towards the central position of SMM4.
The left panel shows the millimeter lines and the right panel shows the submillimeter lines.
The left panel shows the millimeter lines and the right panel shows the submillimeter lines.
In Table 2.. we list the centroid velocities with uncertainties of (he transitions shown in Figure 1.
In Table \ref{tbl-2}, we list the centroid velocities with uncertainties of the transitions shown in Figure \ref{plotall}.
, The centroid velocities were computed over a velocity interval (hat corresponded to the linewidth of the optically thin isotope.
The centroid velocities were computed over a velocity interval that corresponded to the linewidth of the optically thin isotope.
Also listed in the table are the beamwiclths of the observations repeated from Table 1..
Also listed in the table are the beamwidths of the observations repeated from Table \ref{tbl-1}.
several trends can be observed in the aad C'S spectra towards the central position that suggest that infall is occurring towards SMMA,
Several trends can be observed in the and CS spectra towards the central position that suggest that infall is occurring towards SMM4.
The noteworthy feature of most of the and CS lines (the main isotope) in Figure 1 is that they show the classic blue asvamnietric line profile. the signature expected for infall.
The noteworthy feature of most of the and CS lines (the main isotope) in Figure \ref{plotall} is that they show the classic blue asymmetric line profile, the signature expected for infall.
The millimeter transitions of CS and ddo not show identifiable double-peaked line profiles.
The millimeter transitions of CS and do not show identifiable double-peaked line profiles.
The sell-absorption dip in the line mav actually be down at the continuum level (there is some emission recdward of this dip: see Figurel)).
The self-absorption dip in the line may actually be down at the continuum level (there is some emission redward of this dip; see \ref{plotall}) ).
While the CS lline profile does not show a clear blue asvimnietry. its centroid velocity is clearly blueshiltec (see Table 2)).
While the CS line profile does not show a clear blue asymmetry, its centroid velocity is clearly blueshifted (see Table \ref{tbl-2}) ).
In contrast. the optically thin isotopic spectra (the aand CS (transitions) appear more gaussian and centered on the νεο of the object. again as expected in an infall interpretation (Naravananetal.1998:NaravananandWalker1993).
In contrast, the optically thin isotopic spectra (the and $^{34}$ S transitions) appear more gaussian and centered on the $_{LSR}$ of the object, again as expected in an infall interpretation \citep{nwb98,nar98}.
. In all cases. for molecular (racers that probe the infalling material. (he centroid velocity of the corresponding main. more oplically thick isotope would be expected to be lower
In all cases, for molecular tracers that probe the infalling material, the centroid velocity of the corresponding main, more optically thick isotope would be expected to be lower
quarki
diquark.
n the projectileis chosenin prelerence to thespin-clown valence quark bya se
We successfully analyzed the hyperon polarization, but could not explain the
n the projectileis chosenin prelerence to thespin-clown valence quark bya sea
We successfully analyzed the hyperon polarization, but could not explain the
Por=551.798+0.010 clays Tuus=2450479.11+0.07 Pou=1.312904+0.000003 das Tou=2450455.68uοτι—KὧτιLiVTμα Adopting these ephemerides the spectroscopic orbit (svstemic velocity. velocity amplitudes. eccenlricily, periastron passage and mass ratio) plus a Fourier series of order six (which approximates the pulsations of the Cepheid primary component) were fitted to the radi:l velocily data.
$ {\rm P_{orb}} = 551.798 \pm 0.010 $ days $ \hspace*{2cm} {\rm T}_{\rm 0, orb} = 2450479.11 \pm 0.07 $ $ {\rm P_{pul}} = 1.312904 \pm 0.000003 $ days $ \hspace*{2cm} {\rm T}_{\rm 0, pul} = 2450455.685 \pm 0.006 $ Adopting these ephemerides the spectroscopic orbit (systemic velocity, velocity amplitudes, eccentricity, periastron passage and mass ratio) plus a Fourier series of order six (which approximates the pulsations of the Cepheid primary component) were fitted to the radial velocity data.
The orbit solution and the pulsational radial velocity curves of the Cepheid component are shown in Figure 1.
The orbit solution and the pulsational radial velocity curves of the Cepheid component are shown in Figure 1.
Figure 2 shows the pulsational Z-band light curve of the Cepheid.
Figure 2 shows the pulsational $I$ -band light curve of the Cepheid.
Adopting the obtained spectroscopic mass ratio of 0.705 4 0.015. and an I-band third light contribution of 10κ... we modeled the OGLE-LMC-CEP-1812 system with the Wilson Deviunev code (Wilson and Devinnev 1971. Van Lamime and Wilson. 2007) in an ileralive wav removing the intrinsic brightness variations of the Cepheid component.
Adopting the obtained spectroscopic mass ratio of 0.705 $\pm$ 0.015, and an I-band third light contribution of 10, we modeled the OGLE-LMC-CEP-1812 system with the Wilson Devinney code (Wilson and Devinney 1971, Van Hamme and Wilson, 2007) in an iterative way removing the intrinsic brightness variations of the Cepheid component.
The parameters corresponding to our best model together with their uncertainties estimated from extensive Monte Carlo simulations are presented in Table 1.
The parameters corresponding to our best model together with their uncertainties estimated from extensive Monte Carlo simulations are presented in Table 1.
Unfortunately. at (he present moment without near infrared photometry. we are not in a position to derive Teg of the components.
Unfortunately, at the present moment without near infrared photometry, we are not in a position to derive ${\rm T}_{\rm eff}$ of the components.
Figure 3 presents the light curve of the eclipsing svstem with the intrinsic brightness variations of the Cepheid removed. together with our best fit WD model for the svstem.
Figure 3 presents the light curve of the eclipsing system with the intrinsic brightness variations of the Cepheid removed, together with our best fit WD model for the system.
Although the relative distance between components is very large the eclipses are total with the secondary. transiting over (he Cepheid disk during the primary eclipse.
Although the relative distance between components is very large the eclipses are total with the secondary transiting over the Cepheid disk during the primary eclipse.
It is worth mentioning that the presence of the third component only mareinally affects (he mass determination.
It is worth mentioning that the presence of the third component only marginally affects the mass determination.
Neglecting the third light leads to an inclination of 39.5 degrees.
Neglecting the third light leads to an inclination of 89.5 degrees.
The Cepheid in OGLE-LMC-CEP-1812 is a classical Cepheil. and not a Type II Cepheid.
The Cepheid in OGLE-LMC-CEP-1812 is a classical Cepheid, and not a Type II low-mass Cepheid.
The evidence comes [rom its mass of 3.74 4 0.06 Vey which agrees well wilh the predicted. pulsational mass [or a classical Cepheid of this short period (Bono et al.
The evidence comes from its mass of 3.74 $\pm$ 0.06 $M_{\bigodot}$ which agrees well with the predicted pulsational mass for a classical Cepheid of this short period (Bono et al.
2001). and from its radius (17.4 + 0.9 Rey) which again is in good agreement with the racius of a classical Cepheid predicted [rom period-racius relations (e.g. Gieren οἱ al.
2001), and from its radius (17.4 $\pm$ 0.9 $R_{\bigodot}$ ) which again is in good agreement with the radius of a classical Cepheid predicted from period-radius relations (e.g. Gieren et al.
1993).
1998).
The position of the Cepheid in the period-mean magnitude plane for LAIC Cepheids shown in Figure 4 also proves bevond any doubt that CEP-1312 is a classical Cephleid.
The position of the Cepheid in the period-mean magnitude plane for LMC Cepheids shown in Figure 4 also proves beyond any doubt that CEP-1812 is a classical Cepheid.
Dased on (the analvsis of the Fourier parameters Soszvuski et al. (
Based on the analysis of the Fourier parameters Soszynski et al. (
2008) firmly classified. (his star as a [fundamental mode pulsator.
2008) firmly classified this star as a fundamental mode pulsator.
The svstem OGLE-LMC-CEP-1812 is thus the second known double-lined eclipsing binary svstem with a classical huudamental mode. Cepheid component.
The system OGLE-LMC-CEP-1812 is thus the second known double-lined eclipsing binary system with a classical fundamental mode Cepheid component.
The secondary component is a less massive. smaller and cooler stable giant
The secondary component is a less massive, smaller and cooler stable giant
Each model parameter is restricted to a fixed range of values encompassing the plausible span for that parameter.
Each model parameter is restricted to a fixed range of values encompassing the plausible span for that parameter.
The range for each parameter is described below and listed in Table [Il.
The range for each parameter is described below and listed in Table \ref{tab:pars}.
The Levenberg-Marquardt minimization algorithm yields not only best-fit parameters, but also a full covariancematrix (6019) for all parameter pairs.
The Levenberg-Marquardt minimization algorithm yields not only best-fit parameters, but also a full covariancematrix $(C_{jk})$ for all parameter pairs.
Both the variance estimate for each parameter (o;= Cj) and the correlation matrix (pj;ΟνΟζ) provide additional indicators for the quality of the fit.
Both the variance estimate for each parameter $\sigma^2_j=C_{jj}$ ) and the correlation matrix $\rho_{jk}=C_{jk}/\sqrt{C_{jj}C_{kk}}$ ) provide additional indicators for the quality of the fit.
The variance estimates directly measure the constraint that the AIM minimization places on each model parameter for each object analyzed, and the correlation matrices can be used to evaluate and improve the model parametrization we use.
The variance estimates directly measure the constraint that the AIM minimization places on each model parameter for each object analyzed, and the correlation matrices can be used to evaluate and improve the model parametrization we use.
Models with many parameters (this implementation has twelve) can often have significant parameter degeneracies, making the correlation matrix obtained from the fitting a valuable tool for evaluating the specific model parametrization.
Models with many parameters (this implementation has twelve) can often have significant parameter degeneracies, making the correlation matrix obtained from the fitting a valuable tool for evaluating the specific model parametrization.
Large correlation coefficients (|p;,|~ 1) indicate significant degeneracies which can prevent the minimization algorithm from efficiently and accurately converging.
Large correlation coefficients $|\rho_{jk}|\sim1$ ) indicate significant degeneracies which can prevent the minimization algorithm from efficiently and accurately converging.
There is a lensing-parameter/shape-parameter degeneracy which is well known to the lensing community: the shear/ellipticity degeneracy utilized in standard weak lensing studies to measure shear.
There is a lensing-parameter/shape-parameter degeneracy which is well known to the lensing community: the shear/ellipticity degeneracy utilized in standard weak lensing studies to measure shear.
Our approach to mitigating this degeneracy, fixing the shear parameters during fitting, will be discussed in refsec:shearellip,, and the simulations described in refsec:aimtest confirm that there are no other strong degeneracies.
Our approach to mitigating this degeneracy, fixing the shear parameters during fitting, will be discussed in \\ref{sec:shearellip}, and the simulations described in \\ref{sec:aimtest} confirm that there are no other strong degeneracies.
Modest correlations ρε]~0.7 exist for some lensed images and parameter combinations, but even in these instances convergence is robust and accurate.
Modest correlations $|\rho_{jk}|\sim0.7$ exist for some lensed images and parameter combinations, but even in these instances convergence is robust and accurate.
The lensing transformation is characterizedby six variables: the three complex, reduced lensing fields g, Ψι, and V3 from Equation |.
The lensing transformation is characterizedby six variables: the three complex, reduced lensing fields $g$, $\Psi_1$ , and $\Psi_3$ from Equation \ref{eq:reducedflexionlensing}. .
The range allowed for each these parameters
The range allowed for each these parameters
energy to gravitational potential energy of Ξ0.16 and 0.77.
energy to gravitational potential energy of $\beta_{\rm r} = 0.16$ and $\Omega t_{\rm ff} = 0.77$ .
The internal energy of the cloud is set such that the ratio of thermal to gravitational energy is à.=0.26.
The internal energy of the cloud is set such that the ratio of thermal to gravitational energy is $\alpha = 0.26$.
This corresponds to an internal energy of 5.2310 ergs ¢! a sound speed in the cloud of 1.87 10l!em/s and. assuming a mean molecular weight of pj=2. an isothermal temperature of 8.4K. Given the 30:1] density ratio between the cloud and the surrounding medium and the condition of pressure equilibrium between the two. this means that the sound speed .in the external medium. .is c;wedi=10.2.10n ems. and αμαE(CAL/R)=3.1.
This corresponds to an internal energy of $5.23\times 10^{8}$ ergs $^{-1}$, a sound speed in the cloud of $1.87\times 10^{4}$ cm/s and, assuming a mean molecular weight of $\mu=2$, an isothermal temperature of 8.4K. Given the 30:1 density ratio between the cloud and the surrounding medium and the condition of pressure equilibrium between the two, this means that the sound speed in the external medium is $c_{\rm s,medium} = 10.2 \times 10^{4}$ cm/s, and $c_{\rm s,medium}^{2}/(GM/R) = 3.1$.
The initial density perturbations with amplitudes ο=0.1.0.2 were applied. retaining equal particle masses. by perturbing the particle positions from a uniform distribution according to the linearized continuity equation
The initial density perturbations with amplitudes $A=0.1, 0.2$ were applied, retaining equal particle masses, by perturbing the particle positions from a uniform distribution according to the linearized continuity equation
Lard X-ray NLSv1 look very heterogeneous in their broacl-band: properties. displaving a broad range of harcl X-ray (20-100 keV) photon indeces. Hatly. distributed. from. quite lat ( 1.3) to very steep (~ 3.6) slopes.
Hard X-ray NLSy1 look very heterogeneous in their broad-band properties, displaying a broad range of hard X-ray (20-100 keV) photon indeces, flatly distributed from quite flat $\sim$ 1.3) to very steep $\sim$ 3.6) slopes.
At these energies no clear separation between the DELSvI and. NLSv1 slopes distributions is found. although the hard. X-ray. selection may introduce a bias against faint steep spectrum sources.
At these energies no clear separation between the BLSy1 and NLSy1 slopes distributions is found, although the hard X-ray selection may introduce a bias against faint steep spectrum sources.
steeper photon indeces for NLSy1 are instead found. when considering the broad-band: spectra.
Steeper photon indeces for NLSy1 are instead found when considering the broad-band spectra.
In. only one source. Swift 2127.415654. a high energv cut-olf. is. measured at a relatively low value (12,4;orp 50 keV). and the rellection. fraction. constrained. to. R=1.0Os⊓⊽⊥⋡∩≱⊔↕⊔⊔∐⊔⋏∙≟ ∙⋅⋅⋪ previous measurements (Malizia et al.
In only one source, Swift J2127.4+5654, a high energy cut-off is measured at a relatively low value $_{cut-off}$ $\sim$ 50 keV), and the reflection fraction constrained to $^{+0.5}_{-0.4}$, confirming previous measurements (Malizia et al.
2008. Miniutti et. al.
2008, Miniutti et al.
2009).
2009).
Apparently. the hard. X-ray selection is not efficient in detecting strong soft. X-ray NLSv1 as. indeed. only one source. LGR. J19378-0617. shows a dominant ancl strongly variable soft. X-ray component.
Apparently, the hard X-ray selection is not efficient in detecting strong soft X-ray NLSy1 as, indeed, only one source, IGR J19378-0617, shows a dominant and strongly variable soft X-ray component.
“Phe presence. of fully or xwtiallv covering absorption is quite frequently measured in the spectra (in. four out of ten sources analyzed in his work).
The presence of fully or partially covering absorption is quite frequently measured in the spectra (in four out of ten sources analyzed in this work).
When the spectral quality is good (ic... with data). we almost always detect the presence of an Fe line (showing from narrow to moderately broad »ofiles).
When the spectral quality is good (i.e., with data), we almost always detect the presence of an Fe line (showing from narrow to moderately broad profiles).
A proper determination of the reflection strength is imitec by the non simultaneitv of the X-ray and hard X-ray observations.
A proper determination of the reflection strength is limited by the non simultaneity of the X-ray and hard X-ray observations.
This is particularly true in this class of sources where X-ray variability. (in Dux and spectrum) is strong. roth at short and long timescales. as here demonstrated ov the lisht. curves and in those sources where multiple observations were available.
This is particularly true in this class of sources where X-ray variability (in flux and spectrum) is strong, both at short and long timescales, as here demonstrated by the light curves and in those sources where multiple observations were available.
and BAT hard. X-rays average [luxes are consistent. with a maximum ιν variation of 30%..
and BAT hard X-rays average fluxes are consistent, with a maximum flux variation of $\sim$.
Overall. the average X-ray spectrum of NLSy1 looks more similar to the BLSy1 spectrum when the sources are selected at hard X-ray rather than soft N-ravs(e.g... AT).
Overall, the average X-ray spectrum of NLSy1 looks more similar to the BLSy1 spectrum when the sources are selected at hard X-ray rather than soft X-rays (e.g., ).
We estimate that the fraction of NLSv1 in the hard. X- sky is about of the tvpe AGN population. in agreement with the optically selected sample fraction.
We estimate that the fraction of NLSy1 in the hard X-ray sky is about of the type 1 AGN population, in agreement with the optically selected sample fraction.
The black hole mass istribution of NLSv1 selected in hard X-ravs peaks at M. and the Eclclington ratios distribution peaks at LO7. suggesting that hard X-ray NLSv occupy the lower tail of Edcington ratios distribution of NLSv1.
The black hole mass distribution of NLSy1 selected in hard X-rays peaks at $^{7}$ $_{\odot}$ and the Eddington ratios distribution peaks at $^{-2}$, suggesting that hard X-ray NLSy1 occupy the lower tail of Eddington ratios distribution of NLSy1.