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. It therefore seems unlikely that IRSLOSW is a Cepheid vari;le. | It therefore seems unlikely that IRS16SW is a Cepheid variable. |
Nonetheless. the shape of the IRSIGSW light curve is reminiscent of pulsating variable stars (Le. a sharper rise to μπακπαπα lieht followed by a slower decline). | Nonetheless, the shape of the IRS16SW light curve is reminiscent of pulsating variable stars (i.e. a sharper rise to maximum light followed by a slower decline). |
One possibility is that IRSIGSW isa 3 Cephei variable. | One possibility is that IRS16SW is a $\beta$ Cephei variable. |
These are Ligh mass stars (6-30 NL. ) that pulsate due to the &-mechanuisui caused by ion absorption peaked at 77:210? Ix. that excites fuucdiunental mode oscillatious 2001). | These are high mass stars (6-30 $_\odot$ ) that pulsate due to the $\kappa$ -mechanism caused by ion absorption peaked at $\approx$ $\times$ $^5$ K that excites fundamental mode oscillations \citep[see][]{dx01}. |
. Ixnown Cephlei variables are early B-type stars (except for HD 31656: Pigulski&lxolaczkowski(1998):Fullertonetal. (1991))). | Known $\beta$ Cephei variables are early B-type stars (except for HD 34656; \citet{pk98,fullerton91}) ). |
This μιαν be due to the very short lengtli of time higher mass stars spend in the appropriate part of the Hertzsprune-Russell Diagrain or because the pulsatious in higher mass stars are of very small amplitude and difficult to detect (Dziembowki&Pamyatanukl1993:DengXiongὃν--2001). | This may be due to the very short length of time higher mass stars spend in the appropriate part of the Hertzsprung-Russell Diagram or because the pulsations in higher mass stars are of very small amplitude and difficult to detect \citep{dp93,dx01}. |
. Furthermore. observatious of B-type 2 (ορμοί variables suggestMO that the photometric amplitucle of the variation decreases with wavelenethOm (Heynelericksetal.199D). | Furthermore, observations of B-type $\beta$ Cephei variables suggest that the photometric amplitude of the variation decreases with wavelength \citep{heynderickx94}. |
. Also. the periods of 3 Cephei variables are typically «1 day. | Also, the periods of $\beta$ Cephei variables are typically $<$ 1 day. |
Both these are not particularly consistent. with observatious of IRSIGSW. although there are uo observations ol this class of variables in the iufrared or any mocleling predictiug their behaviour. | Both these are not particularly consistent with observations of IRS16SW, although there are no observations of this class of variables in the infrared or any modeling predicting their behaviour. |
We are not aware of any other well-observed class of pulsating stars that have characteristics slinilar to IRSI6SW. | We are not aware of any other well-observed class of pulsating stars that have characteristics similar to IRS16SW. |
The Ser C complex (C359.ba15-(05) is located 2307 for 75 pe) in projection to the South of the Ser A complex. aloug the ic plane. | The Sgr C complex (G359.45-0.05) is located $\sim$ (or 75 pc) in projection to the South of the Sgr A complex, along the Galactic plane. |
Figure 31 shows the 1.1 GHz continuum image of the ain compouents of the Ser C Spinmnσον: the SerC HIT region. the NTF. an HII regious kuown as FIR 1. Source C and L2 (Liszt&ker(1995).. Roy(2003).. aud Ocenwad&Fazio(1981))). | Figure 34 shows the 1.4 GHz continuum image of the main components of the Sgr C complex: the SgrC HII region, the NTF, and HII regions known as FIR 4, Source C and D \citet{LS95}, \citet{R03}, and \citet{OF84}) ). |
The NTF has been labeled “Part Av alle "Part B after Roy (20()3). | The NTF has been labeled “Part A” and “Part B” after Roy (2003). |
There are sjuiliarities between Ser C and the Badio Arc: the Ser C HII region may coutain more tlau 250 solar tnasses of ioulzed gas with massive stars present. incluist£& some late-tyye O stars. | There are similiarities between Sgr C and the Radio Arc: the Sgr C HII region may contain more than 250 solar masses of ionized gas with massive stars present, including some late-type O stars. |
Tus HII regiMI ls ASSOClated with a molectlar cloud which has tle saiJe velocity (~—6: Skms LT | This HII region is associated with a molecular cloud which has the same velocity $\sim$$-$ 65 ). |
ie molecular ‘loud associated with Ser C appears to be orgauized in a shell-like distribution surrotudiug he HII region: the massive stars are thought to have blown a cavi vin the ¢istribution of gas (Lizst 5jxker 1995). | The molecular cloud associated with Sgr C appears to be organized in a shell-like distribution surrounding the HII region; the massive stars are thought to have blown a cavity in the distribution of gas (Lizst Spiker 1995). |
Finally.“ the Ser C NTF is also thought to be physically relaed to thil5 5ar forijatlon 'eglon. | Finally, the Sgr C NTF is also thought to be physically related to this star formation region. |
Fiewes 35-37 show the correspoucliig inteeratec| HI abso‘ptiou profiles for each of these regioS Part À sud B of the Ser C NTE. the SerC HIT 'egion. FIRI aud Sotrees C' aud D. Most of tje profiles lave a complicated velocity stΠοιο Wwith absoὪjohn at many velocities. | Figures 35-37 show the corresponding integrated HI absorption profiles for each of these regions: Part A and B of the Sgr C NTF, the SgrC HII region, FIR4, and Sources C and D. Most of the profiles have a complicated velocity structure with absorption at many velocities. |
As mentioned above. 1iere are molecular clouds in tlis region with velocities of —65 and —100 kin noted by Liszt&Spiker(1995). aud Okaetal.(2001). | As mentioned above, there are molecular clouds in this region with velocities of $-$ 65 and $-$ 100, noted by \citet{LS95}
and \citet{O01}. |
. "ngu'e 38 shows contours of 1.1 GHz radio COLWUui euission overlaid ou greyscale representing the HI optical depth at velocities of —62 and —10() | Figure 38 shows contours of 1.4 GHz radio continuum emission overlaid on greyscale representing the HI optical depth at velocities of $-$ 65 and $-$ 100. |
L2 Figure 39 shows the same HI optical depth as in Figure 38 overlaid ou coutours of CO (J—]-0) emission from the survey of Osa et al. ( | Figure 39 shows the same HI optical depth as in Figure 38 overlaid on contours of CO (J=1-0) emission from the survey of Oka et al. ( |
1998) at velocities of —65 aud — LOC0 | 1998) at velocities of $-$ 65 and $-$ 100. |
Figure 35 shows the HI absorption profile toward Part A and Part B of the Ser C NTF. | Figure 35 shows the HI absorption profile toward Part A and Part B of the Sgr C NTF. |
These proliles are siular with HI absorption near —G60-65 towa«€ both regions of the NTE. but clearly strouger toward Part B of tie NTF. | These profiles are similar with HI absorption near $-$ 60-65 toward both regions of the NTF, but clearly stronger toward Part B of the NTF. |
This dilference in the HI absorption by the —65 cloud across tve Ser C NTF issriking and may indicate that lie atomic gas in the cirection of Part A is em»edded in or located on the backside of the —65 molecular cloud. | This difference in the HI absorption by the $-$ 65 cloud across the Sgr C NTF is striking and may indicate that the atomic gas in the direction of Part A is embedded in or located on the backside of the $-$ 65 molecular cloud. |
In addition. there is HlaSOLPion near — LOOla in both profiles. | In addition, there is HI absorption near $-$ 100 in both profiles. |
Iu Par . "mthe HI absorption occurs at —136 aud —1t(X)fs: both com»onelis nay be associated wil1 —10t) inolecular cloud if there is a velocity gradieut. | In Part B, the HI absorption occurs at $-$ 136 and $-$ 109; both components may be associated with the $-$ 100 molecular cloud if there is a velocity gradient. |
Figures 35 ai a38 (right) illustate that the HI absorption by the —10t) cloud. is strouger toward Part A of the NTF (τμ.ο) than toward Part B of the NTF (7y,;=6).1[5). | Figures 35 and 38 (right) illustrate that the HI absorption by the $-$ 100 cloud is stronger toward Part A of the NTF $\tau$$_{HI}$ =0.15) than toward Part B of the NTF $\tau$$_{HI}$ =0.45). |
In addition. Figure:39 (right) shows tha the distribution of molecular gas traced in the CO emission near —10(hans peaks near Part A of the NTF and is nearly absent near Part B of the Ser C NTF. | In addition, Figure 39 (right) shows that the distribution of molecular gas traced in the CO emission near $-$ 100 peaks near Part A of the NTF and is nearly absent near Part B of the Sgr C NTF. |
This correlation of the strouger HI absorption at —100kin with the peak of molecular eas suggests that tlie atomic gas Is pivsically associated with the molecular cloud at this location | This correlation of the stronger HI absorption at $-$ 100 with the peak of molecular gas suggests that the atomic gas is physically associated with the molecular cloud at this location |
featuring such a low-mass star could survive the supernova explosion. | featuring such a low-mass star could survive the supernova explosion. |
Combining the constraints from the optical/IR flux limits to the ones derived from the timing properties of the source (?),, we may test the hypothesis of the existence of a disk formed of supernova fallback material, surrounding the NS. | Combining the constraints from the optical/IR flux limits to the ones derived from the timing properties of the source , we may test the hypothesis of the existence of a disk formed of supernova fallback material, surrounding the NS. |
Such a debris disk has been invoked as a possible explanation of several puzzling properties of the X-ray source. | Such a debris disk has been invoked as a possible explanation of several puzzling properties of the X-ray source. |
In order to compute the expected optical/IR. flux from such a putative disk, we consider a geometrically thin, optically thick fallback disk, locally emitting as a blackbody. | In order to compute the expected optical/IR flux from such a putative disk, we consider a geometrically thin, optically thick fallback disk, locally emitting as a blackbody. |
Following a standard approach, we included two contributions (?):: (i) viscous dissipation in the disk, yielding a temperature profile T(r)οςr-?/4(?); (i) reprocessing of X-rays emitted by the central object. | Following a standard approach, we included two contributions : (i) viscous dissipation in the disk, yielding a temperature profile $T(r)\propto r^{-3/4}$; (ii) reprocessing of X-rays emitted by the central object. |
Assuming Lx=2.2x10?? erg s! for a distance of 2.2 kpc and a standard disk structure yields T(r)=1100K(1—9)?"(Rg/r)/7(?), where η is the X-ray albedo of the disk. | Assuming $L_X=2.2\times10^{33}$ erg $^{-1}$ for a distance of 2.2 kpc and a standard disk structure yields $T(r)=1100K(1-\eta)^{2/7}(R_{\odot}/r)^{3/7}$, where $\eta$ is the X-ray albedo of the disk. |
The total flux is obtained by integrating the emissivity over the disk surface and taking into account the inclination of the disk with respect to the line of sight, as well as the source distance. | The total flux is obtained by integrating the emissivity over the disk surface and taking into account the inclination of the disk with respect to the line of sight, as well as the source distance. |
The inner radius of the disk is set at the distance from the star where the magnetic pressure of the rotating dipole of the neutron star disrupts the disk itself. | The inner radius of the disk is set at the distance from the star where the magnetic pressure of the rotating dipole of the neutron star disrupts the disk itself. |
Taking into account the role of viscosity in the disk, such a radius is ry~0.5R,; where Ry ids the so-called Alfven radius?). | Taking into account the role of viscosity in the disk, such a radius is $r_M\sim0.5R_M$ where $R_M$ is the so-called Alfven radius. |
. The outer radius of the disk may be constrained by the flux limits. | The outer radius of the disk may be constrained by the flux limits. |
We assumed a disk inclination 4=60°. | We assumed a disk inclination $i=60^{\circ}$. |
The contribution of the reprocessed X-ray flux critically depends on the poorly known value for the X-ray albedo η. | The contribution of the reprocessed X-ray flux critically depends on the poorly known value for the X-ray albedo $\eta$. |
Although a value ηΞ0.5 has been assumed in several investigations??),, a much larger value η=0.97 was evaluated by?,, based on the detection of a disk around AXP 0142+61. | Although a value $\eta=0.5$ has been assumed in several investigations, a much larger value $\eta=0.97$ was evaluated by, based on the detection of a disk around AXP 0142+61. |
Thus, we used 7=0.97 (yielding a much lower disk luminosity) in our study. | Thus, we used $\eta=0.97$ (yielding a much lower disk luminosity) in our study. |
In the above assumptions, we computed the expected flux in a specific filter band as a function of the NS magnetic field and disk accretion rate m, for a set of values of the disk outer radius. | In the above assumptions, we computed the expected flux in a specific filter band as a function of the NS magnetic field and disk accretion rate $\dot{m}$, for a set of values of the disk outer radius. |
We repeated such exercise for all of the bands of our optical/infrared dataset. | We repeated such exercise for all of the bands of our optical/infrared dataset. |
For completeness, we also used the Spitzer 4.5um and 8.4um bands studied by?. | For completeness, we also used the Spitzer $4.5\mu m$ and $8.4\mu m$ bands studied by. |
. The most constraining limit turned out to be those in the FF555W and F814W bands reported here. | The most constraining limit turned out to be those in the F555W and F814W bands reported here. |
In the allowed range of parameters, the model is rather insensitive to the value of the disk outer radius. | In the allowed range of parameters, the model is rather insensitive to the value of the disk outer radius. |
The flux contributions are almost negligible in all bands for Τουι larger than ~3x1013 cm. | The flux contributions are almost negligible in all bands for $r_{out}$ larger than $\sim3\times10^{12}$ cm. |
Results are plotted in Figure 3 where a red line marks the m-B field region allowed by the combination of the F555W and F814W flux limits. | Results are plotted in Figure \ref{limits} where a red line marks the $\dot{m}$ -B field region allowed by the combination of the F555W and F814W flux limits. |
The region ruled out by llimits is colored in yellow. | The region ruled out by limits is colored in yellow. |
For comparison, we also marked (in orange) the region formerly ruled out by the less constraining Spitzer/IRAC 4.5jm limit. | For comparison, we also marked (in orange) the region formerly ruled out by the less constraining Spitzer/IRAC $4.5\mu m$ limit. |
Flux limits in other bands are less constraining than the Spitzer one. | Flux limits in other bands are less constraining than the Spitzer one. |
In order to use the timing constraintsf07, wenotethattheinteractiono ftheputativediskwitht. with respect to the neutron star light cylinder as well as to the corotation radius(?). | In order to use the timing constraints, we note that the interaction of the putative disk with the rotating neutron star magnetosphere should yield different regimes of angular momentum transfer, according to the relative positions of $r_M$ with respect to the neutron star light cylinder as well as to the corotation radius. |
. We evaluated the expected neutron star P as a function of the NS magnetic field and disk m in such different regimes (the so-called ejector, propeller and accretor regimes) using standard relations??). | We evaluated the expected neutron star $\dot{P}$ as a function of the NS magnetic field and disk $\dot{m}$ in such different regimes (the so-called ejector, propeller and accretor regimes) using standard relations. |
. This allows to identify a rm-B field region allowed by the existing limit on P, which is overplotted in Figure 3.. | This allows to identify a $\dot{m}$ -B field region allowed by the existing limit on $\dot{P}$, which is overplotted in Figure \ref{limits}. |
The line marking the P limit has been drawn arbitrarily (in dashed style) in the small region connecting the propeller regime to the ejector regime (where standard relations for the propeller torque are not valid), to visualize the reduced efficiency of the propeller effect as the magnetospheric radius approaches the light cylinder. | The line marking the $\dot{P}$ limit has been drawn arbitrarily (in dashed style) in the small region connecting the propeller regime to the ejector regime (where standard relations for the propeller torque are not valid), to visualize the reduced efficiency of the propeller effect as the magnetospheric radius approaches the light cylinder. |
The region ruled out by X-ray timing is colored in grey. | The region ruled out by X-ray timing is colored in grey. |
A region in the accretor regime yielding a luminosity larger than 2x10?? erg s! (ie. the total X-ray luminosity of assuming a distance of 2 kpc) is also ruled out. | A region in the accretor regime yielding a luminosity larger than $2\times10^{33}$ erg $^{-1}$ (i.e. the total X-ray luminosity of assuming a distance of 2 kpc) is also ruled out. |
It is marked in red. | It is marked in red. |
The region allowed by both the P limit and the optical/infrared flux limits is coloured in green in Figure 3.. | The region allowed by both the $\dot{P}$ limit and the optical/infrared flux limits is coloured in green in Figure \ref{limits}. |
We can infer some interesting indications about the possible role of a fallback disk. | We can infer some interesting indications about the possible role of a fallback disk. |
In the absence of any model describing a quantitative relation between the accretion rate and the phenomenology of the spectral features, we will focus here on the issue of the luminosity of the hot spot. | In the absence of any model describing a quantitative relation between the accretion rate and the phenomenology of the spectral features, we will focus here on the issue of the luminosity of the hot spot. |
The maximum allowed disk Πι is slightly | The maximum allowed disk $\dot{m}$ is slightly |
it is important to investigate the clleet of binary stars on the derived LOSVD. | it is important to investigate the effect of binary stars on the derived LOSVD. |
Since. Monte. Carlo calculations at present are unable to vield. the full LOSVD or even higher order moments such as the kurtosis without excessive computational effort. they cannot. be emploved. in. this context. | Since Monte Carlo calculations at present are unable to yield the full LOSVD or even higher order moments such as the kurtosis without excessive computational effort, they cannot be employed in this context. |
Llence. we give analvtical expressions for the velocity moments of the binary stars and for the LOSVD itself. (approximating all binaries to consist. of stars on circular orbits). | Hence, we give analytical expressions for the velocity moments of the binary stars and for the LOSVD itself (approximating all binaries to consist of stars on circular orbits). |
We first give a number of definitions of important quantities and clarify some notations. | We first give a number of definitions of important quantities and clarify some notations. |
We focus our attention on a single binary system. | We focus our attention on a single binary system. |
One star. the “primary”. has a mass A and the other. the “secondary” has a mass m. | One star, the “primary”, has a mass $M$ and the other, the “secondary” has a mass $m$. |
E orbital parameters pertain to the primary's orbit The specific binary LOSVD gives the probability of finding he primary star with mass À/ with a linc-of-sight velocity in he interval Py—Ae,/2.e,|/2]. given that it revolves around a secondary star with mass m on an elliptical orbit. with orbital parameters e. e. c anc 7. | All orbital parameters pertain to the primary's orbit: The specific binary LOSVD gives the probability of finding the primary star with mass $M$ with a line-of-sight velocity in the interval $[v_p-\Delta v_p/2,v_p+\Delta v_p/2]$, given that it revolves around a secondary star with mass $m$ on an elliptical orbit with orbital parameters $a$, $e$, $\omega$ and $i$. |
Intuitively. his probability is proportional to the fraction of its period during which the primary has a velocity in this interval. | Intuitively, this probability is proportional to the fraction of its period during which the primary has a velocity in this interval. |
Mathematicallv. this translates into the following expression or the specific binary LOSVD: withορ the line-of-sight velocity of the primary at phase angle © on its orbit Llere. Llence. if follows that The specific binary LOSVD is a function solely of e. | Mathematically, this translates into the following expression for the specific binary LOSVD: with$v_p$ the line-of-sight velocity of the primary at phase angle $\phi$ on its orbit Here, Hence, if follows that and consequently The specific binary LOSVD is a function solely of $v_p$ . |
Lo | To |
motion of Io in the Jovian magnetospheric plasma. | motion of Io in the Jovian magnetospheric plasma. |
In the ultra-relativistic wind of a pulsar. it can be approximated very simply by (See Eq. ( | In the ultra-relativistic wind of a pulsar, it can be approximated very simply by (See Eq. ( |
67-69) of (MH1) for more details.) | 67-69) of (MH1) for more details.) |
Then. adopting a simplified geometry. it is possible to estimate the total electric current. | Then, adopting a simplified geometry, it is possible to estimate the total electric current. |
Neubauer(1980) gives useful expressions for the total current 7 flowing along an Alfvénn wing. | \citet{Neubauer_1980} gives useful expressions for the total current $I$ flowing along an Alfvénn wing. |
Writing Rp for the body's radius. he gets: The electric field £;. set along the body. is caused by its ionosphere or surface internal resistance. | Writing $R_P$ for the body's radius, he gets: The electric field $E_i$, set along the body, is caused by its ionosphere or surface internal resistance. |
The Joule dissipation is maximum when £;=Eo/2. | The Joule dissipation is maximum when $E_i = E_0/2$. |
In our estimations. we shall use Neubauer's values for /. | In our estimations, we shall use Neubauer's values for $I$. |
The above theory. because of the involved symmetries. describes mainly what happens in space. far enough from the body. | The above theory, because of the involved symmetries, describes mainly what happens in space, far enough from the body. |
At closer distances. the plasma suffers compressive motions and compressive MHD waves certainly have a non-negligible influence on the system. | At closer distances, the plasma suffers compressive motions and compressive MHD waves certainly have a non-negligible influence on the system. |
These waves propagate quasi-isotropically. | These waves propagate quasi-isotropically. |
Their amplitude decrases as the inverse of the distance to the body and they contribute to deflect the wind around it. | Their amplitude decrases as the inverse of the distance to the body and they contribute to deflect the wind around it. |
Nevetheless. without entering into these consideration. we can still make a few inferences based on the theory of the Alfvénn wings. as presented in the previous section. | Nevetheless, without entering into these consideration, we can still make a few inferences based on the theory of the Alfvénn wings, as presented in the previous section. |
As Neubauer(1980).. we can assume that the current associated to the wing is closed in the vicinity of the body (see Fig. 5)). | As \citet{Neubauer_1980}, we can assume that the current associated to the wing is closed in the vicinity of the body (see Fig. \ref{fig_inducteur_unipolaire}) ), |
through its surface or its tonosphere. | through its surface or its ionosphere. |
We can estimate (roughly) what force the wind exerts on it. | We can estimate (roughly) what force the wind exerts on it. |
The two Alfvénn wings carry a current that. in the two branches flowing along the body (perpendicular to the plane of Fig 5)). generates a force density jxB. | The two Alfvénn wings carry a current that, in the two branches flowing along the body (perpendicular to the plane of Fig \ref{fig_inducteur_unipolaire}) ), generates a force density $\ve{j} \times \ve{B}$. |
The two current systems flowing on each side of the body exert this force density in the same radial direction. | The two current systems flowing on each side of the body exert this force density in the same radial direction. |
We may expect the body to orbit near the equatorial plane of the pulsar. | We may expect the body to orbit near the equatorial plane of the pulsar. |
In this plane. at such a distance. the magnetic field direction is almost azimuthal. being perpendicular to the wind flow velocity. | In this plane, at such a distance, the magnetic field direction is almost azimuthal, being perpendicular to the wind flow velocity. |
The sign of Bi depends on whether the magnetic moment of the neutron star is parallel or antiparallel to the rotation axis. | The sign of $B_0^\phi$ depends on whether the magnetic moment of the neutron star is parallel or antiparallel to the rotation axis. |
Nevertheless. the force density jxB always has the same direction as the wind velocity. | Nevertheless, the force density $\ve{j} \times \ve{B}$ always has the same direction as the wind velocity. |
At first order. considering Eq. (4)). | At first order, considering Eq. \ref{eq_unipolar_pour_application_numerique}) ), |
Ey=O,V/r and the force is expressed explicitely as a function of the distance r from the pulsar to the body as The power £, dissipated by Joule effect along the ionosphere or in the body is maximized when the internal load matches the external one. that is. still according to (1980).. when E;=£o/2. | $E_0=\Omega_* \Psi /r$ and the force is expressed explicitely as a function of the distance $r$ from the pulsar to the body as The power $\dot E_{J}$ dissipated by Joule effect along the ionosphere or in the body is maximized when the internal load matches the external one, that is, still according to \citet{Neubauer_1980}, when $E_i=E_0/2$. |
In that case the force is On the night side of the body. this force tends to wipe out the ionosphere (if there is one). but on the day side. on the contrary. it pushes the tonosphere towards it. | In that case the force is On the night side of the body, this force tends to wipe out the ionosphere (if there is one), but on the day side, on the contrary, it pushes the ionosphere towards it. |
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