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The period ratio between different classes of QPP have no obvious trend.
The period ratio between different classes of QPP have no obvious trend.
This fact may imply that there is no originated link between different classes of QPP, even if they are occurred simultaneous in the same frequency range.
This fact may imply that there is no originated link between different classes of QPP, even if they are occurred simultaneous in the same frequency range.
Actually, it is possible that the short periodic QPP (e.g. fast-VSP, etc) is a small quasi-periodic perturbation which superposed on the longer periodic QPPs (e.g. VLP, etc), and the latter may dominate the whole evolution of the flaring processes.
Actually, it is possible that the short periodic QPP (e.g. fast-VSP, etc) is a small quasi-periodic perturbation which superposed on the longer periodic QPPs (e.g. VLP, etc), and the latter may dominate the whole evolution of the flaring processes.
However, so far, because of the lack of imaging observations with spatial resolutions in the corresponding
However, so far, because of the lack of imaging observations with spatial resolutions in the corresponding
with three Gaussians. each characterised by centroid and a width parameter taken as the standard deviation of the Gaussian. c.
with three Gaussians, each characterised by centroid and a width parameter taken as the standard deviation of the Gaussian, $\sigma$.
Two were narrow. having c of only a few A.
Two were narrow, having $\sigma$ of only a few $\AA$.
The third component was much broader. with σ up to 20 A.
The third component was much broader, with $\sigma$ up to 20 $\AA$.
This third component was identified with an origin in the wind above the accretion disk. and it is the subject of the present paper.
This third component was identified with an origin in the wind above the accretion disk, and it is the subject of the present paper.
Fig.l] of Blundell. Bowler Schmidtobreick (2008) displays the centroids of these fitted Gaussians as a function of time. in the same sequence as the montage in Fig.2 of Schmidtobreick Blundell (2006).
Fig.1 of Blundell, Bowler Schmidtobreick (2008) displays the centroids of these fitted Gaussians as a function of time, in the same sequence as the montage in Fig.2 of Schmidtobreick Blundell (2006).
Because of their importance. | have reproduced these data in Fig.1.
Because of their importance, I have reproduced these data in Fig.1.
This figure shows that the two narrow Gaussian components of Ha have centroids that scarcely move over more than two orbits. and the broad component oscillates in Doppler shift with a period of 13 days.
This figure shows that the two narrow Gaussian components of $\alpha$ have centroids that scarcely move over more than two orbits, and the broad component oscillates in Doppler shift with a period of 13 days.
It is most redshifted close to orbital phase 0. at primary eclipse.
It is most redshifted close to orbital phase 0, at primary eclipse.
This might suggest that this component ts formed in an aceretion stream rather than in the wind.
This might suggest that this component is formed in an accretion stream rather than in the wind.
The two narrow components were identified with the inner rim of a cireumbinary disk in Blundell. Bowler Schmidtobreick (2008).
The two narrow components were identified with the inner rim of a circumbinary disk in Blundell, Bowler Schmidtobreick (2008).
A very similar plot has been obtained from a set of Ha spectra taken between 2004 and 2008 at the Purple Mountain Observatory.
A very similar plot has been obtained from a set of $\alpha$ spectra taken between 2004 and 2008 at the Purple Mountain Observatory.
The original data are to be found in Fig.2 of Li Yan (2010) and the motions of fitted centroids as a function of orbital phase are displayed in Fig.2 of Bowler (2011a).
The original data are to be found in Fig.2 of Li Yan (2010) and the motions of fitted centroids as a function of orbital phase are displayed in Fig.2 of Bowler (2011a).
The same features are apparent.
The same features are apparent.
In Fig.2 I show the variation of Doppler speed for the centre of the He broad component (Blundell. Bowler Schmidtobreick 2008) and above it the same for the broader component of the He I 6678 lline. obtained from my fits reported in Bowler (201Ib).
In Fig.2 I show the variation of Doppler speed for the centre of the $\alpha$ broad component (Blundell, Bowler Schmidtobreick 2008) and above it the same for the broader component of the He I 6678 line, obtained from my fits reported in Bowler (2011b).
These broad components have centroids that oscillate with a day period and with velocity amplitude ~110 km s7!.
These broad components have centroids that oscillate with a 13-day period and with velocity amplitude $\sim$ 110 km $^{-1}$.
They are most redshifted a little before orbital phase O and most blueshifted a little before 0.5.
They are most redshifted a little before orbital phase 0 and most blueshifted a little before 0.5.
Lines produced in an atmosphere co-moving with the compact object would be most redshifted at orbital phase 0.75 and most blueshifted at 0.25 - the broad lines associated with the wind lag the motion of the compact object by ~ 0.2 of a period and have a reduced orbital Doppler amplitude.
Lines produced in an atmosphere co-moving with the compact object would be most redshifted at orbital phase 0.75 and most blueshifted at 0.25 - the broad lines associated with the wind lag the motion of the compact object by $\sim$ 0.2 of a period and have a reduced orbital Doppler amplitude.
The identification of the broad component of the Haw emission line with an origin in the wind from the disk is discussed in Blundell. Bowler Schmidtobreick (2008).
The identification of the broad component of the $\alpha$ emission line with an origin in the wind from the disk is discussed in Blundell, Bowler Schmidtobreick (2008).
In brief. the width parameter ο drops smoothly from approximately 20 nnear JD +245 (precession phase approximately 0; accretion disk wide open) to 10 nnear JD +274 as the accretion disk comes closer to edge on.
In brief, the width parameter $\sigma$ drops smoothly from approximately 20 near JD +245 (precession phase approximately 0; accretion disk wide open) to 10 near JD +274 as the accretion disk comes closer to edge on.
In addition. this measure of the line of sight wind speed follows the nodding motion of the accretion disk. as inferred from the Doppler shifts of the relativistic jets (Blundell. Bowler Schmidtobreick 2007).
In addition, this measure of the line of sight wind speed follows the nodding motion of the accretion disk, as inferred from the Doppler shifts of the relativistic jets (Blundell, Bowler Schmidtobreick 2007).
The synthesis of absorption line studies of the wind (Fabrika 1997, 2004) shows that the line of sight expansion speed varies as approximately the square of the cosine of the angle between the jet axis and the line of sight. reaching 1600 km s! for 60°.
The synthesis of absorption line studies of the wind (Fabrika 1997, 2004) shows that the line of sight expansion speed varies as approximately the square of the cosine of the angle between the jet axis and the line of sight, reaching 1600 km $^{-1}$ for $^{\circ}$.
Thus these observations have established that the source of this wind line ts rooted 1n the accretion disk.
Thus these observations have established that the source of this wind line is rooted in the accretion disk.
In Blundell. Bowler and Schmidtobreick (2008) the authors cautiously observed that despite this. the motion of the centroid of the wind should not be taken as a measure of the orbital velocity of the compact object.
In Blundell, Bowler and Schmidtobreick (2008) the authors cautiously observed that despite this, the motion of the centroid of the wind should not be taken as a measure of the orbital velocity of the compact object.
Ifà given parcel of wind continues emitting Ha over a period of several days. the data displayed in Fig.2 are reconciled with an origin in the tilting and nodding aceretion disk. according to the scenario sketched in the introduction.
If a given parcel of wind continues emitting $\alpha$ over a period of several days, the data displayed in Fig.2 are reconciled with an origin in the tilting and nodding accretion disk, according to the scenario sketched in the introduction.
Absorption line studies have shown that the wind from SS 433 is slow in the plane of the accretion disk. but that the speed increases rapidly as the angle to the jet axis decreases.
Absorption line studies have shown that the wind from SS 433 is slow in the plane of the accretion disk, but that the speed increases rapidly as the angle to the jet axis decreases.
The results are summarised in Fabrika (1997. 2004) in the form where V,. is the velocity (in km s) of the gas flowing out from the disk as a function of the polar angle a.
The results are summarised in Fabrika (1997, 2004) in the form where $V_w$ is the velocity (in km $^{-1}$ ) of the gas flowing out from the disk as a function of the polar angle $\alpha$.
The width c of the broad H« component in the data of Blundell. Bowler Schmidtobreick (2008) has a value of ~20 eearly on. when the angle y between the Jet axis and the line of
The width $\sigma$ of the broad $\alpha$ component in the data of Blundell, Bowler Schmidtobreick (2008) has a value of $\sim$ 20 early on, when the angle $\chi$ between the jet axis and the line of
+).
$^{-1}$ ).
lOcQsPECTRUAL shows the results of a spectra svnthesis of the continuum emission.
\\ref{FIG:QSPECTRUM} shows the results of a spectral synthesis of the continuum emission.
We find that the accretion discs light has a power-law index of 0.60250.3 ane contributes 37-18 percent to the observed continuum ligh atAA.
We find that the accretion disc's light has a power-law index of $\pm$ 0.3 and contributes $\pm$ 13 percent to the observed continuum light at.
From the residual of the fit it can be seen that a power-law model for the cise predicts more fux short-ware ofSOOOA.. so a power-Iaw. function does not best describe the disc's light.
From the residual of the fit it can be seen that a power-law model for the disc predicts more flux short-ward of, so a power-law function does not best describe the disc's light.
LE we moclel the spectrum of the disc with a blackbody function. we obtain a sienificantly better fi (at the 99 percent level). where the dise contributes 5843:16 percent to the continuum [ux at aancd the disc has a blackbody temperature of 4600+ Why and a radius of 0.502£0.
If we model the spectrum of the disc with a blackbody function, we obtain a significantly better fit (at the 99 percent level), where the disc contributes $\pm$ 16 percent to the continuum flux at and the disc has a blackbody temperature of $\pm$ K and a radius of $\pm$.
061... Using a WS star in order to match the blue end of the spectrum does not give a better fit (at the 99.99 percent confidence level.)
Using a K3 star in order to match the blue end of the spectrum does not give a better fit (at the 99.99 percent confidence level.)
ALL the uncertainties quoted are 1-0 and have been rescaled so that the 42 ofthe fit is 1.
All the uncertainties quoted are $\sigma$ and have been rescaled so that the $\chi_{\nu}^{2}$ of the fit is 1.
Although our determination of the fractional contribution of the accretion disc's light to that observed is consistent with the findings of Oke(1077).. 7? and MeClintocketal.(1995). the form of power-law description of the clise’s light is not.
Although our determination of the fractional contribution of the accretion disc's light to that observed is consistent with the findings of \citet{Oke77}, \citet{MR86} and \citet{MHR95}, the form of power-law description of the disc's light is not.
However. it should. be noted that the non-variable accretion disc light is à composite of light from the steady-state accretion disc plus the excess light from star-spots or a late-superhump. which most probably have very dilferent spectral shapes.
However, it should be noted that the non-variable accretion disc light is a composite of light from the steady-state accretion disc plus the excess light from star-spots or a late-superhump, which most probably have very different spectral shapes.
For the case where the excess light is due to a late-superhump. the late-superhump modulation may have its origin in the outer regions of the disc in the changing streame-cdisc interaction (Rolle. Laswell Pattterson 2001). the spectrum. of the [ate-superhump is expected. to. be different compared to the steady-state disc.
For the case where the excess light is due to a late-superhump, the late-superhump modulation may have its origin in the outer regions of the disc in the changing stream-disc interaction (Rolfe, Haswell Pattterson 2001), the spectrum of the late-superhump is expected to be different compared to the steady-state disc.
Although the tically heated streame-disc impact region will be hotter than the rest of the outer disc. it is not obvious how much hotter this will be compared. to. the inner regions of the disc. where viscous stresses due to dillerential rotation are much higher.
Although the tidally heated stream-disc impact region will be hotter than the rest of the outer disc, it is not obvious how much hotter this will be compared to the inner regions of the disc, where viscous stresses due to differential rotation are much higher.
Thus it is. cillicult to estimate the spectrum. of the late-superhump without detailed computations.
Thus it is difficult to estimate the spectrum of the late-superhump without detailed computations.
Therefore. it is no surprise that the discs light and spectrum are observed to change with time. as it just reflects the variable behaviour of the accretion disc.
Therefore, it is no surprise that the disc's light and spectrum are observed to change with time, as it just reflects the variable behaviour of the accretion disc.
Assuming that the light produced by the Lares is simply added) t0. the quiescent spectrum. which contains the secondary star and light from the accretion disc (assumed to be non-variable). the dillerence between Hare-state spectra and the quict-state spectra. vields an estimate. for. the spectrum of the Dares.
Assuming that the light produced by the flares is simply added to the quiescent spectrum which contains the secondary star and light from the accretion disc (assumed to be non-variable), the difference between flare-state spectra and the quiet-state spectra yields an estimate for the spectrum of the flares.
Phe actual portions of the lighteurves used to determine the Dare-state and (quiet-state spectra are marked in retIG:PLAIUEPOS and were selected after subtracting the secondary star's ellipsoidal modulation.
The actual portions of the lightcurves used to determine the flare-state and quiet-state spectra are marked in \\ref{FIG:FLAREPOS} and were selected after subtracting the secondary star's ellipsoidal modulation.
retICESPECTTRUSM shows the resulting spectrum. which has been binned for clarity.
\\ref{FIG:FSPECTRUM} shows the resulting spectrum, which has been binned for clarity.
We compared the average Hare spectrum taken during the beginning and end of the night.
We compared the average flare spectrum taken during the beginning and end of the night.
Fhey. were found to be the same to within 1.5 percent.
They were found to be the same to within 1.5 percent.
Before the subtraction we shift the spectra into the rest. frame of the secondary. star. so that the features from the secondary are removed cleanly.
Before the subtraction we shift the spectra into the rest frame of the secondary star, so that the features from the secondary are removed cleanly.
The I-band mag varies on a timescale of a few hundred cays with an amplitude of 0.3 mag (LeibowitzLemar&Orio908).
The R-band mag varies on a timescale of a few hundred days with an amplitude of 0.3 mag \citep{Leibowitz98}.
. However. since this timescale is much longer than he orbital period it is safe to assume that the lighteurve of he superhump of excess light does not change shape over he length of the orbital period.
However, since this timescale is much longer than the orbital period it is safe to assume that the lightcurve of the superhump of excess light does not change shape over the length of the orbital period.
The Hare spectrum has a relatively fat continuum sugeesting that a high temperature model is needed. to it the data.
The flare spectrum has a relatively flat continuum suggesting that a high temperature model is needed to fit the data.
We attempt to fit the [lare spectrum. with iWerent models.
We attempt to fit the flare spectrum with different models.
A fit using a blackbody (423.93). gives v temperature of 5900+200 Why and a racius 0.10 ((90 percent confidence).
A fit using a blackbody $\chi^{2}_{\nu}$ =3.93) gives a temperature of $\pm$ K and a radius $\pm$ (90 percent confidence).
A power-law fit (4523.85);E S of w form £\xA has an index of -0.6020.20. or pDUET (00 percent. confidence)
A power-law fit $\chi^{2}_{\nu}$ =3.85) of the form $F_\lambda \propto \lambda^\beta$ has an index of $\pm$ 0.20, or $F_\nu \propto \nu^{-1.40\pm0.20}$ (90 percent confidence).
A more realistic model is ᾱ continuum emission μαvectrum of an LEE slab of hydrogen.
A more realistic model is a continuum emission spectrum of an LTE slab of hydrogen.
Using the svnthetic shotometry SYNPLIIOT package/STSDAS) to compute 16 L'TI models we fit the continuum regions to estimate 10 temperature. racius and barvon column density of the 4ab.
Using the synthetic photometry SYNPHOT package) to compute the LTE models we fit the continuum regions to estimate the temperature, radius and baryon column density of the slab.
We find that the X7 surface has a broad. minimum at vo of ~4.
We find that the $\chi^2$ surface has a broad minimum at $\chi^{2}_{\nu}$ of $\sim$ 4.
The requirement that the Hare region is smaller iun the area of the accretion disc Rp: Marshοἱal. 1994)) adds a weak constraint. only ruling out temperatures ess than 25000 Why. We constrain the temperature and equivalent: radius of the optically thin LYLE slab to lie in the range Kis and 0.0320.044 ((99 percent confidence) for a barvon column density in the range 107. 1074 cem7 (see rotέως ΕΠ}.
The requirement that the flare region is smaller than the area of the accretion disc $R_{\rm L1}$; \citealt{Marsh94}) ) adds a weak constraint, only ruling out temperatures less than $\sim$ K. We constrain the temperature and equivalent radius of the optically thin LTE slab to lie in the range K and 0.032–0.044 (99 percent confidence) for a baryon column density in the range $^{20}$ $^{24}$ $^{-3}$ (see \\ref{FIG:CHI}) ).
Phe emission. covers 0.050.08 percent (99 percent confidence) of the accretion clise’s projected surface area (q—0.067. /—41 and Αιξ Marshetal.1994 and Celineetal. 2001).
The emission covers 0.05–0.08 percent (99 percent confidence) of the accretion disc's projected surface area $q$ =0.067, $i$ $^{\circ}$ and $M_{\rm 1}$; \citealt{Marsh94} and \citealt{Gelino01}) ).
Tight constraintsΤΙM. on the parameters of interest cannot be placed because of the correlations oween the temperature ancl column. density: ᾱ- high emperature and low column density model gives the same Voas a low temperature and high. column density. model.
Tight constraints on the parameters of interest cannot be placed because of the correlations between the temperature and column density; a high temperature and low column density model gives the same $\chi^2$ as a low temperature and high column density model.
Although this can only be done by using the Balmer emission ine fluxes. for this is very dillicult as it is clear that the Balmer emission ines are contaminated by the emission from the bright-spot (see retDOPPLEDS )).
Although this can only be done by using the Balmer emission line fluxes, for this is very difficult as it is clear that the Balmer emission lines are contaminated by the emission from the bright-spot (see \\ref{DOPPLER}) ).
AAqr is an unusual cataclysmic variable in that it exhibits. flaring behaviour which can be described. in the ramework of a magnetic propeller that throws out eas out of the binary svstem (Wynn.Wing&Lorne1997).
Aqr is an unusual cataclysmic variable in that it exhibits flaring behaviour which can be described in the framework of a magnetic propeller that throws out gas out of the binary system \citep{Wynn97}.
. The laves are thought to arise from collisions between high-density regions in the material expelled from the system alter interaction with the rapidly rotating magnetosphere of he white cwarl (Pearson.Horne&Skidmore 2003)..
The flares are thought to arise from collisions between high-density regions in the material expelled from the system after interaction with the rapidly rotating magnetosphere of the white dwarf \citep{Pearson03}. .
The spectrum of the Dares in AAqr can be described by a optically thin gas with a temperature of Why (Welsh.Horne&Oke 1993)..
The spectrum of the flares in Aqr can be described by a optically thin gas with a temperature of K \citep{Welsh93}. .
The spectrum of the [lares seen in aare cleseribecl by optically thin. eas with a similar
The spectrum of the flares seen in are described by optically thin gas with a similar
in the intermediate separation. and about two zones in the simall separation case.
in the intermediate separation, and about two zones in the small separation case.
At the convective core boundary location the zone size is about 0.00Ll of the surface radius aud a radial zone contains about 0.1 ML. when integrated over all angles.
At the convective core boundary location the zone size is about 0.004 of the surface radius and a radial zone contains about 0.1 $\Msun$ when integrated over all angles.
From these cousideratious it :j»pears that the interior quautities are not much affected by the presence of the companion.
From these considerations it appears that the interior quantities are not much affected by the presence of the companion.
We vow turn to the surface configuration.
We now turn to the surface configuration.
e We shall frame our discussion in terms of the elongation of the model.
We shall frame our discussion in terms of the elongation of the model.
The elongation is defined as the ratio of the surface radius in the direction opposite the binary companion to the surface radius in the direction of the binary companion.
The elongation is defined as the ratio of the surface radius in the direction opposite the binary companion to the surface radius in the direction of the binary companion.
When the ratio is above about 0.9. we find that the Roche equipotential fits the surface shape as accurately as we can measure with our discrete zoning (approximately 0.1 per cent).
When the ratio is above about 0.9, we find that the Roche equipotential fits the surface shape as accurately as we can measure with our discrete zoning (approximately 0.4 per cent).
As the elongation approaches 0.5 the two shapes are slightly differeut. mostly iu the direction between the two components of the binary system.
As the elongation approaches 0.8 the two shapes are slightly different, mostly in the direction between the two components of the binary system.
This is shown in Fig. 13..
This is shown in Fig. \ref{fig3},
a plot of the surface shape for the largeste separatiou model just before the egravitational contraction phase beeius.
a plot of the surface shape for the largest separation model just before the gravitational contraction phase begins.
e The difference in surface shape between the ROTORC aud Roche models becomes more significant as Roche lobe overllow is approached.
The difference in surface shape between the ROTORC and Roche models becomes more significant as Roche lobe overflow is approached.
A comparison for a model very close to Roche lobe overflow (elongation of about 0.69) for the intermeciate separation case is shown in Fig. 11...
A comparison for a model very close to Roche lobe overflow (elongation of about 0.69) for the intermediate separation case is shown in Fig. \ref{fig4}.
There is clearly a clillerence in the two potentials iu the direction between the two binary componeuts.
There is clearly a difference in the two potentials in the direction between the two binary components.
To be fair the Roche potential contours are chaneine quite rapidly here iu the direction of looking more like the ROTORC contours for slightly larger fractional raclil.
To be fair the Roche potential contours are changing quite rapidly here in the direction of looking more like the ROTORC contours for slightly larger fractional radii.
Clearly one factor in the differiug potentials is the possibility that the self gravity of the primary at the model surface is not that given by the point source potential.
Clearly one factor in the differing potentials is the possibility that the self gravity of the primary at the model surface is not that given by the point source potential.
The self gravitational potential ol the primary ou a spherical surface whose radius is given by the largest surface radius in the model is shown in Fig. 15..
The self gravitational potential of the primary on a spherical surface whose radius is given by the largest surface radius in the model is shown in Fig. \ref{fig5}.
Clearly the magnitude of the potential[un is largest iu the direction between the wo components. and the amplitude variation is a little more than one percent.
Clearly the magnitude of the potential is largest in the direction between the two components, and the amplitude variation is a little more than one percent.
Another way of viewing this is to examine what we refer to as the "column mass”.
Another way of viewing this is to examine what we refer to as the “column mass".
This is deliued by integrating the adial density distribution at a given augle over a spherical volume element: Le.. it is the interior uass distribution tlie moclel would have if this racial distribution was spherically syiuinetric.
This is defined by integrating the radial density distribution at a given angle over a spherical volume element; i.e., it is the interior mass distribution the model would have if this radial distribution was spherically symmetric.
We show this column mass for three aueles versus radial zone number in Fig. 16..
We show this column mass for three angles versus radial zone number in Fig. \ref{fig6},
where it is evident hat there is little dillerence iu the mass distribution interior to approximately 5M...
where it is evident that there is little difference in the mass distribution interior to approximately $\Msun$.
Closer to the uodel surface there is more mass concentrated in the direction toward the companion. aud the columa Wass ruouotonically decreases golug away [rom this direction.
Closer to the model surface there is more mass concentrated in the direction toward the companion, and the column mass monotonically decreases going away from this direction.
One interesting feature related to the shape of the surface is the rate of change in that shape as the time of Roche lobe overflow is approached.
One interesting feature related to the shape of the surface is the rate of change in that shape as the time of Roche lobe overflow is approached.
We compare the expansion rate of the surface iu the direction of the companion for the largest separation model with the surface expansion of a spherical model in Fig. 17..
We compare the expansion rate of the surface in the direction of the companion for the largest separation model with the surface expansion of a spherical model in Fig. \ref{fig7}.
Clearly. the primary surface in the direction of the secondary. is expaudiug at a much larger rate. one that is fairly close to tlie expausion experienced iu the early phases of hydrogeu-shell burning.
Clearly, the primary surface in the direction of the secondary is expanding at a much larger rate, one that is fairly close to the expansion experienced in the early phases of hydrogen-shell burning.
The expansion rate iucreases vet faster as Roche lobe overflow is approached. as is evideuced in Fig. 18..
The expansion rate increases yet faster as Roche lobe overflow is approached, as is evidenced in Fig. \ref{fig8}.
Here we illustrate the expausion rate of the surface of the primary iu the direction of the secondary in the intermediate separation case.
Here we illustrate the expansion rate of the surface of the primary in the direction of the secondary in the intermediate separation case.
The expansion rate is sufficiently rapid that it is possible that the stellar surface does uot remain an equipoteutial during
The expansion rate is sufficiently rapid that it is possible that the stellar surface does not remain an equipotential during
Ch.
\citep{rey2002}.
This cluster is dominated by three ceutral OB stars that constitute the primary excitation sources of the IIIT region (?7)..
This cluster is dominated by three central OB stars that constitute the primary excitation sources of the HII region \citep{zei1978,smi1985}.
The stars in this cluser appear to be heavily obscured along our liue of sight. with ely~LO magnitudes of visual extinction. (?)..
The stars in this cluster appear to be heavily obscured along our line of sight, with $A_{V} \sim 10$ magnitudes of visual extinction \citep{shu1999}.
Wl10 IRS2a has been suggested as he dominant source of ionizing radiation iu this region (?7). and milluueter observations have shown evidence for cireunustellar dust shells around IRSla. 2a. aud 3a 1}.
W40 IRS2a has been suggested as the dominant source of ionizing radiation in this region \citep{smi1985}, and millimeter observations have shown evidence for circumstellar dust shells around IRS1a, 2a, and 3a \citep{smi1985,val1994}.
From the collective results of these studies we can see that the W10 ΠΠ reeion has broken through the surrounding molecular cloud iu a direction hat lies at an angle to our liue of sight.
From the collective results of these studies we can see that the W40 HII region has broken through the surrounding molecular cloud in a direction that lies at an angle to our line of sight.