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The same aiglow emission was used as wavelength reference (Oliva Origlia 1992)) and the spectra were corrected for instrumental and atmospheric transmission using spectra of featureless early O stars and flux calibrated using measurements of photometric standard stars.
The same airglow emission was used as wavelength reference (Oliva Origlia \cite{oliva92}) ) and the spectra were corrected for instrumental and atmospheric transmission using spectra of featureless early O stars and flux calibrated using measurements of photometric standard stars.
The quasi-complete spectra of the nucleus and of knot C are shown in Figs 2.. 3. respectively. and the derived line fluxes are summarized in Table 1..
The quasi-complete spectra of the nucleus and of knot C are shown in Figs \ref{spec_nuc}, \ref{spec_knc} respectively, and the derived line fluxes are summarized in Table \ref{tab_obs1}.
Dilution by a stellar continuum is particularly strong in the nuclear spectrum where the equivalent width of OLIL[AS5007 is only 50 (cf.
Dilution by a stellar continuum is particularly strong in the nuclear spectrum where the equivalent width of 5007 is only 50 (cf.
Table 2)) and a factor of ~10 lower than found in typical Seyfert 2s.
Table \ref{tab_obs2}) ) and a factor of $\sim$ 10 lower than found in typical Seyfert 2's.
The stellar contribution is normally estimated and subtracted using either off-nuclear spectra extracted from the same 2D long slit frames. or a suitable combination of spectra of non-active galaxies used as templates (e.g. Ho 1996.. Ho et al. 1997)).
The stellar contribution is normally estimated and subtracted using either off–nuclear spectra extracted from the same 2D long slit frames, or a suitable combination of spectra of non–active galaxies used as templates (e.g. Ho \cite{ho96}, Ho et al. \cite{ho97}) ).
However. neither of the methods proved particularly useful because line emission contaminates the stellar emission all along the slit. and we could not find any template which accurately reproduces the prominent stellar absorption features typical of quite young stellar populations.
However, neither of the methods proved particularly useful because line emission contaminates the stellar emission all along the slit, and we could not find any template which accurately reproduces the prominent stellar absorption features typical of quite young stellar populations.
The fluxes of weak lines (<5% of the continuum) in the nucleus are therefore uncertain and. in a few cases. quite different than those reported in O94.. the largest discrepancy being for |NI] which ts a factor of 2 fainter here.
The fluxes of weak lines $<$ of the continuum) in the nucleus are therefore uncertain and, in a few cases, quite different than those reported in \cite{O94}, the largest discrepancy being for [NI] which is a factor of 2 fainter here.
The spectrum of knot C has a much more favourable line/econtinuum ratio and shows many faint lines which are particularly useful for the modelling described in Sect. 4..
The spectrum of knot C has a much more favourable continuum ratio and shows many faint lines which are particularly useful for the modelling described in Sect. \ref{photion_model}.
The spatial variation of the most important lines is visualized in Figs. 4..
The spatial variation of the most important lines is visualized in Figs. \ref{velcont},
5 which show contour plots of the continuum subtracted long slit spectra and selected spectral sections of the various knots respectively.
\ref{spec_all} which show contour plots of the continuum subtracted long slit spectra and selected spectral sections of the various knots respectively.
The fluxes are summarized in Table 2. together with the extinctions which were derived from hydrogen recombination lines assuming standard case-B ratios (Hummer Storey 1987)).
The fluxes are summarized in Table \ref{tab_obs2} together with the extinctions which were derived from hydrogen recombination lines assuming standard case-B ratios (Hummer Storey \cite{hummer}) ).
A remarkable result is the large variations of the typical line diagnostic ratios0.[OI]/Ha.. παπα which are plotted in Fig.
A remarkable result is the large variations of the typical line diagnostic ratios, and which are plotted in Fig.
6 and range from values typical of high excitation Seyferts (nucleus. knots A. B. C. D). to low excitation LINERs (knots H. 1) and normal HII regions (knots E. L).
\ref{knotdiag} and range from values typical of high excitation Seyferts (nucleus, knots A, B, C, D), to low excitation LINERs (knots H, I) and normal HII regions (knots E, L).
Another interesting result is the steep extinction gradient between the regions outside (knots C. D. H. D and those close to the galactic disk (nucleus and knots A. E. L).
Another interesting result is the steep extinction gradient between the regions outside (knots C, D, H, I) and those close to the galactic disk (nucleus and knots A, E, L).
However. a comparison between the mmap (Moorwood Oliva 1994)). the images (M94)) and the observed Bro flux from the whole galaxy (M96)). do not show evidence of more obscured (ly— 30) tonized regions such as those observed in NGC4945 and other starburst galaxies (e.g. Moorwood Oliva 1988)).
However, a comparison between the map (Moorwood Oliva \cite{invited94}) ), the images \cite{M94}) ) and the observed $\alpha$ flux from the whole galaxy \cite{M96}) ), do not show evidence of more obscured $\AV\!\sim\!10\!-\!30$ ) ionized regions such as those observed in NGC4945 and other starburst galaxies (e.g. Moorwood Oliva \cite{moorwood88}) ).
Nevertheless. these data cannot exclude the presence of deeply embedded ionized gas which ts obscured even at 4/m (1.8. ο Οι).
Nevertheless, these data cannot exclude the presence of deeply embedded ionized gas which is obscured even at $\mu$ m (i.e. $>$ 50 mag).
Particularly interesting is the variation of the line ratios between the adjacent knots C and D. The ratio is à factor of2 larger in knot D than in C. but this most probably reflects variations of the iron gas phase abundance (see also Sect. 4.5)).
Particularly interesting is the variation of the line ratios between the adjacent knots C and D. The ratio is a factor of 2 larger in knot D than in C, but this most probably reflects variations of the iron gas phase abundance (see also Sect. \ref{iron}) ).
Much more puzzling is the spatial variation. of the low excitation lines [OI]. [SH]. [NIE] which drop by a factor1.8. while the high excitation lines Hell. [OHI]. [Nell]. [ArH] together with the [SII] density sensitive ratio and [OIII/[OII] vary by much smaller amounts (cf.
Much more puzzling is the spatial variation of the low excitation lines [OI], [SII], [NII] which drop by a factor1.8, while the high excitation lines HeII, [OIII], [NeIII], [ArIII] together with the [SII] density sensitive ratio and [OIII]/[OII] vary by much smaller amounts (cf.
Table 2)).
Table \ref{tab_obs2}) ).
This
This
Sununarizing. metallicity increases the C—74 and Af—Tj colors in such a wav that objects move along the reddening line on (he color-color diagram. which is parallel to the SSP tracks of varving metallicitv.
Summarizing, metallicity increases the $C-T_1$ and $M-T_1$ colors in such a way that objects move along the reddening line on the color-color diagram, which is parallel to the SSP tracks of varying metallicity.
Svstematically redder C—71 at a given metallicity indicate older ages. redder horizontal branches. or the influence of emission-line objects on the integrated colors.
Systematically redder $C-T_1$ at a given metallicity indicate older ages, redder horizontal branches, or the influence of emission-line objects on the integrated colors.
This is why cores of disrupted dwarl galaxies. containing multiple stellar populations. mav have bluer (V—734) ancl redder C—2) colors.
This is why cores of disrupted dwarf galaxies, containing multiple stellar populations, may have bluer $(M-T_1)$ and redder $C-T_1$ colors.
The GCs of Gi and G3 seem (o be located roughly in the same region of the color-color diagram. but the scatter in G3 is large. presumably. because the photometric uncertainties on these fainter objects are larger. and precludes any definite interpretation.
The GCs of G1 and G3 seem to be located roughly in the same region of the color-color diagram, but the scatter in G3 is large, presumably because the photometric uncertainties on these fainter objects are larger, and precludes any definite interpretation.
The GCs of G2 and a subset of LSB dwarl galaxies stand out in Fig.8: most of them lie on a track of vounger age and/or of higher metallicity than the GCs of G1 and the Galactic GCs,
The GCs of G2 and a subset of LSB dwarf galaxies stand out in Fig.8: most of them lie on a track of younger age and/or of higher metallicity than the GCs of G1 and the Galactic GCs.
Cellone Forte (1996) interpret the "deviating branch” of LSB dwail galaxies as caused by a mixture of stellar populations. including vounger components.
Cellone Forte (1996) interpret the "deviating branch" of LSB dwarf galaxies as caused by a mixture of stellar populations, including younger components.
We adopt (his interpretation for G2 and argue that (he GCs of this group. which are (he most massive GCs. have several generations of stars.
We adopt this interpretation for G2 and argue that the GCs of this group, which are the most massive GCs, have several generations of stars.
This property is shared by a growing number of massive GCs in our own Galaxy (Piotto. 2009 and references therein).
This property is shared by a growing number of massive GCs in our own Galaxy (Piotto, 2009 and references therein).
However. (hese galactic GCs are metal-poor. while G2 is composed mostly of metalrich GCs.
However, these galactic GCs are metal-poor, while G2 is composed mostly of metal-rich GCs.
Furthermore. galactic GCs appear to be intermediate between G2 and GI in Fig &: their average mass is 5.2 + 0.6 in loe(AL/AL. ). using the mass estimates of MeLaughlin Van Der Marel (2005). thus closer to GI than to G2.
Furthermore, galactic GCs appear to be intermediate between G2 and G1 in Fig 8; their average mass is 5.2 $\pm$ 0.6 in $_\odot$ ), using the mass estimates of McLaughlin Van Der Marel (2005), thus closer to G1 than to G2.
In other words. the GCs in G2 bear little resemblance to the galactie GCs. and their large mass presumably allowed for multiple generations of stars more like what occurs in galaxies. thanks (to their large potential well which retained the metals lost to the stars.
In other words, the GCs in G2 bear little resemblance to the galactic GCs, and their large mass presumably allowed for multiple generations of stars more like what occurs in galaxies, thanks to their large potential well which retained the metals lost to the stars.
The kinematic properties of the different groups may provide clues to their origin.
The kinematic properties of the different groups may provide clues to their origin.
G1 rotates in the same wav as (he majority of GCs and PNe of the galaxy (Woodley et al.
G1 rotates in the same way as the majority of GCs and PNe of the galaxy (Woodley et al.
2007
2007
a clear view of the region around these systems (Fie 1). it is not possible to sav whether we have discovered au intermediate redshift compact group. or a subcomponcut of a larger structure.
a clear view of the region around these systems (Fig 4), it is not possible to say whether we have discovered an intermediate redshift compact group, or a subcomponent of a larger structure.
We would like to thank the staff of W. ML Necks Observatory for their kind support for our observations.
We would like to thank the staff of W. M. Keck Observatory for their kind support for our observations.
We would also like to thank the referee for useful comments.
We would also like to thank the referee for useful comments.
YO ijs a JSPS Research Fellow.
YO is a JSPS Research Fellow.
WDV acknowledges partial support in the form of a fellowship from the Beatrice Watson Parrent Foundation.
WDV acknowledges partial support in the form of a fellowship from the Beatrice Watson Parrent Foundation.
JET thanks D. M. Whittle aud T. XN. Thuan for useful conversations.
JEH thanks D. M. Whittle and T. X. Thuan for useful conversations.
This work was partly supported bv the Japanese Ministry of Education. Science. aud Culture (Nos.
This work was partly supported by the Japanese Ministry of Education, Science, and Culture (Nos.
10011052. and 10301013).
10044052, and 10304013).
Some of the data prescuted herein were obtained at the WAL Weel Observatory. which is operated as a scientific partucrship among the California Institute of Technology. the University of Califormia and the National Aeronautics and Space Adiuinistratiou.
Some of the data presented herein were obtained at the W.M. Keck Observatory, which is operated as a scientific partnership among the California Institute of Technology, the University of California and the National Aeronautics and Space Administration.
The Observatory was made possible bv the eeuerous financial support of the WAL Iseck Foundation.
The Observatory was made possible by the generous financial support of the W.M. Keck Foundation.
This research has made extensive use of the NASA/TPAC Extragalactic Database. (NED). which is operated by the Jet Propulsion Laboratory. California lustitute of Technology. under coutract with the National Acronauties aud Space Adiiuistration.
This research has made extensive use of the NASA/IPAC Extragalactic Database (NED), which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration.
The Digitized Sky Surveys were produced at the Space Telescope Scieuce Iustitute under U.S. Coveriuent eraut NAC W-2166.
The Digitized Sky Surveys were produced at the Space Telescope Science Institute under U.S. Government grant NAG W-2166.
The iuages of these surveys are based on photographic data obtained using the Oschiu Schinidt Telescope ou Palomar Mountain and the UI& Schinidt Telescope.
The images of these surveys are based on photographic data obtained using the Oschin Schmidt Telescope on Palomar Mountain and the UK Schmidt Telescope.
The plates were processed into the present compressed digital form with the permission of these institutions.
The plates were processed into the present compressed digital form with the permission of these institutions.
Astrophysical disks are common because the specilic angular momentum of the matter inside them is well-conservecl.
Astrophysical disks are common because the specific angular momentum of the matter inside them is well-conserved.
They evolve because angular momentum conservation is weakly compromised. either because of diffusion of angular momentum within the disk or because of direct application of external torques.
They evolve because angular momentum conservation is weakly compromised, either because of diffusion of angular momentum within the disk or because of direct application of external torques.
In astrophysical disks composed of a well-ionized plasma it is likely that some. perhaps most. of the evolution is driven by diffusion of angular momentum within the disk.
In astrophysical disks composed of a well-ionized plasma it is likely that some, perhaps most, of the evolution is driven by diffusion of angular momentum within the disk.
This view is cerlainly consistent with observations of steadily accreting cataclysmic variable svstems like UX Ursa Majoris (Baptistaetal.19983:Daptista2004).. whose radial surface-briehtness profile is consistent wilh steady. accretion-Llow models in which the bulk of the accretion enerev is dissipated within the disk.
This view is certainly consistent with observations of steadily accreting cataclysmic variable systems like UX Ursa Majoris \citep{bap98,bap04}, whose radial surface-brightness profile is consistent with steady accretion-flow models in which the bulk of the accretion energy is dissipated within the disk.
Although deuterium is a trace element (D/H ~10™). the study of deuterium fractionation in molecules is à very vivid. field. of research.
Although deuterium is a trace element (D/H $\sim$ $^{-5}$ ), the study of deuterium fractionation in molecules is a very vivid field of research.
The main astrochemical interest derives from the fact that the understanding of fractionation processes may ultimately help to understand the formation of the related hydrogenated species.
The main astrochemical interest derives from the fact that the understanding of fractionation processes may ultimately help to understand the formation of the related hydrogenated species.
As an example. the first detection of a doubly-deuterated molecule. D:CO. by Was interpreted as the hint that grain chemistry may be at work for anynthesizing formaldehyde.
As an example, the first detection of a doubly-deuterated molecule, $_2$ CO, by was interpreted as the hint that grain chemistry may be at work for synthesizing formaldehyde.
However. the impact of deuterium fractionation studies rapidly outpassed the mere astrochemical nterest.
However, the impact of deuterium fractionation studies rapidly outpassed the mere astrochemical interest.
Since deuterated molecules are formed preferentially at low temperatures and high densities. they became valuable probes for studying. e.g.. the kinematics of cold cloud cores harbouring the earliest stages of star formation(?).
Since deuterated molecules are formed preferentially at low temperatures and high densities, they became valuable probes for studying, e.g., the kinematics of cold cloud cores harbouring the earliest stages of star formation.
. The basic process transferring deuterium from the HD molecule to other molecules is through 10n-molecule reactions involving Hy. CHy. and CH;therein).
The basic process transferring deuterium from the HD molecule to other molecules is through ion-molecule reactions involving $_3^+$, $_3^+$, and $_2$ $_2^+$.
. The route involving Hy ts the dominant one at very low temperatures (7:220 KK). whereastheimportanceofCHy takes over at slightly higher temperatures (20; 7 1440 Koduetoitshugherendothermicity
The route involving $_3^+$ is the dominant one at very low temperatures $T$ $<$ K), whereas the importance of $_3^+$ takes over at slightly higher temperatures $<$ $T$ $<$ K) due to its higher endothermicity.
Yo??)..H3D is thus a very good probe of very cold objects. and has been extensively used in the last years to probe gas in which CO is strongly depleted. i.e.. frozen out on dust grain surfaces?).
$_2$ $^+$ is thus a very good probe of very cold objects, and has been extensively used in the last years to probe gas in which CO is strongly depleted, i.e., frozen out on dust grain surfaces.
. The deuteration level is also à decisive argument in the study of the evolutionary state of early phases of star formation.
The deuteration level is also a decisive argument in the study of the evolutionary state of early phases of star formation.
The deuterium fractionation ts believed to decrease once the central protostar starts to heat its envelope?).
The deuterium fractionation is believed to decrease once the central protostar starts to heat its envelope.
. For instance. partly based on molecular deuteration. concluded that the dense core Cha-MMS1 inthe Chamaeleon I molecular cloud is likely at the stage of the first hydrostatic core when H» has not been dissociated yet.
For instance, partly based on molecular deuteration, concluded that the dense core Cha-MMS1 inthe Chamaeleon I molecular cloud is likely at the stage of the first hydrostatic core when $_2$ has not been dissociated yet.
The study of several isotopologues of methanol toward the low-mass protostar IRAS16293-2422 showed that grain chemistry models require a very high atomic. D/H ratio accreting on the grains in order to explain the high CHD:;OH and CD;OH abundances.
The study of several isotopologues of methanol toward the low-mass protostar IRAS16293-2422 showed that grain chemistry models require a very high atomic D/H ratio accreting on the grains in order to explain the high $_2$ OH and $_3$ OH abundances.
At the time. gas-phase models could not reproduce such a high atomic D/H ratio. but a decisive step forward was done with the inclusion of D)H™ and D; in the reaction networks(?).
At the time, gas-phase models could not reproduce such a high atomic D/H ratio, but a decisive step forward was done with the inclusion of $_2$ $^+$ and $_3^+$ in the reaction networks.
. Since then. D.:H was searched for through observation of its. ground para state line.
Since then, $_2$ $^+$ was searched for through observation of its ground para state line.
The frequency of this line was initially published to be 691.660440 GHz and was later revised to 691.660483 GHz?).
The frequency of this line was initially published to be 691.660440 GHz and was later revised to 691.660483 GHz.
. These astronomical searches are extremely difficult because of the high frequeney of the molecular line. for which excellent instruments and weather conditions are required.
These astronomical searches are extremely difficult because of the high frequency of the molecular line, for which excellent instruments and weather conditions are required.
Using the Caltech Submillimeter Observatory mm telescope (CSO). detected a ~3.30 line (in peak) toward the Ophiuchus dense core 16293E (4.4c in integrated intensity).
Using the Caltech Submillimeter Observatory m telescope (CSO), detected a $\sim$ $\sigma$ line (in peak) toward the Ophiuchus dense core 16293E $\sigma$ in integrated intensity).
This detection is to date the only published one. and its low signal-to-noise ratio as well as the velocity shift between H:D and D>H™ lafterremeasurementofo and certainly require a confirmation of the detection by improving the observation sensitivity and frequency calibration precision (which was ~ kkmss~! in the observations of ?)).
This detection is to date the only published one, and its low signal-to-noise ratio as well as the velocity shift between $_2$ $^+$ and $_2$ $^+$ after remeasurement of and certainly require a confirmation of the detection by improving the observation sensitivity and frequency calibration precision (which was $\sim$ $^{-1}$ in the observations of ).
We present here the first secure detection of D4H toward the prestellar core H-MMI in L1688. the main molecular cloud in the Ophiuchus star forming region.
We present here the first secure detection of $_2$ $^+$ toward the prestellar core H-MM1 in L1688, the main molecular cloud in the Ophiuchus star forming region.
The L1688 cloud (the main cloud in the Ophiuchus dark cloud) has been mapped both in molecular tracers and in continuum emission?).
The L1688 cloud (the main cloud in the Ophiuchus dark cloud) has been mapped both in molecular tracers and in continuum emission.
. These early mapping studies revealed different areas of high column density. named Oph-A to Oph-G. Using the Submillimeter Common User Bolometer Array (SCUBA) at the I5mm James-Clerck-Maxwell Telescope (JCMT). covered a larger region. leading to the detection of two new dense cores. that they named Oph-H and Oph-l. following the same nomenclature.
These early mapping studies revealed different areas of high column density, named Oph-A to Oph-G. Using the Submillimeter Common User Bolometer Array (SCUBA) at the m James-Clerck-Maxwell Telescope (JCMT), covered a larger region, leading to the detection of two new dense cores, that they named Oph-H and Oph-I, following the same nomenclature.
Oph-H contains a single continuum peak. H-
Oph-H contains a single continuum peak, H-MM1.
The distance of the Ophiuchus region has been recently accurately determined to ppc (?).
The distance of the Ophiuchus region has been recently accurately determined to pc .
. We present in Fig.
We present in Fig.
| a part of the SCUBA map of the Ophiuchus region. as available from the SCUBA Legacy
\ref{rho_oph} a part of the SCUBA map of the Ophiuchus region, as available from the SCUBA Legacy
42 charge coupled devices (CCDs).
42 charge coupled devices (CCDs).
The boresight of the telescope remains constant Lor (he mission. but the spacecraft rolls 90 degrees every three months.
The boresight of the telescope remains constant for the mission, but the spacecraft rolls 90 degrees every three months.
Due to restrictions in memory and bandwidth. only the pixels associated with the target stars are sent to the eround for processing.
Due to restrictions in memory and bandwidth, only the pixels associated with the target stars are sent to the ground for processing.
Target stars are defined by software. aud. pixels not associated with Largels are saved only infrequently.
Target stars are defined by software, and pixels not associated with targets are saved only infrequently.
The basic integration time is 6.02 seconds. ancl (he lone cadence (LC) sequence (Jenkinsetal.2010). co-adds (he target pixels For 29.4 minutes.
The basic integration time is 6.02 seconds, and the long cadence (LC) sequence \citep{jen10} co-adds the target pixels for 29.4 minutes.
The co-adcded pixel data are sent to the ground every month lor processing and analvsis.
The co-added pixel data are sent to the ground every month for processing and analysis.
This strategy enables the extremely high signal-to-noise (SNR) observations needed to achieve the planetary detection mission.
This strategy enables the extremely high signal-to-noise (SNR) observations needed to achieve the planetary detection mission.
To obtain (he large field of view. the image sampling is verv coarse compared to other astrometric assets.
To obtain the large field of view, the image sampling is very coarse compared to other astrometric assets.
The Pixel Response Function (PRE) is described in great. detail by Drysonetal.(2010).. but a quick summary is as follows.
The Pixel Response Function (PRF) is described in great detail by \citet{bry10}, but a quick summary is as follows.
The images contain three components. a sharp spike in Che middle (hat comes from optical diffraction (about 0.1 arcsecond). a wider component with a characteristic size of 5 or 6 arcseconds set primarily by mechanical alignment tolerances of the CCD mosaic. and a much broader scattering prolile.
The images contain three components, a sharp spike in the middle that comes from optical diffraction (about 0.1 arcsecond), a wider component with a characteristic size of 5 or 6 arcseconds set primarily by mechanical alignment tolerances of the CCD mosaic, and a much broader scattering profile.
The intermediate component of (he image prolile produces (he majority of the astrometric signal. and it contains about of the light.
The intermediate component of the image profile produces the majority of the astrometric signal, and it contains about of the light.
The 43 days of data available so [ar have shown a remarkable astrometric precision. but the demonstration of astrometric accuracy is still a work in progress.
The 43 days of data available so far have shown a remarkable astrometric precision, but the demonstration of astrometric accuracy is still a work in progress.
Even if the centroiding process was [fully understood. the short interval of available data precludes the lifting of the degeneracies between effects of proper motion. parallax. ancl velocity. aberration.
Even if the centroiding process was fully understood, the short interval of available data precludes the lifting of the degeneracies between effects of proper motion, parallax, and velocity aberration.
No measured astrometric parameters
No measured astrometric parameters
Iu the case when the disk is optically thick iu continua. he disk outflow may be driven not oulv by the radiation rou the stellar surface. but also by the stellar radiation reprocessed w the disk (Cawleyetal.1999).
In the case when the disk is optically thick in continuum, the disk outflow may be driven not only by the radiation from the stellar surface, but also by the stellar radiation reprocessed by the disk \citep{abla}.
.. To estimate he optical depth of the disk. let us assume Lyedrogen aud relimm to be ionized iun the disk.
To estimate the optical depth of the disk, let us assume hydrogen and helium to be ionized in the disk.
Iu this case a siguificaut vat of the disk optical depth originates due to the light scattering on free electrons (for wavelengths lower than hat corresponding to the Balmer or Lyman jump also )ounud-free transitions nav contribute).
In this case a significant part of the disk optical depth originates due to the light scattering on free electrons (for wavelengths lower than that corresponding to the Balmer or Lyman jump also bound-free transitions may contribute).
The transverse optical depth is then roughly given by 7—[παρα= πα. where αι 1 the Thomson scattering cross-section per unit of mmass.
The transverse optical depth is then roughly given by $\tau=\int\kappa_\text{e}\rho\,\de z=\kappa_\text{e}\Sigma$ , where $\kappa_\text{e}$ is the Thomson scattering cross-section per unit of mass.
The disk is optically thick iu the vertical direction (7> 1) if the mass-loss rate is larger than For a given mass-loss rate the disk is optically thick close to the star. while becomine optically thin at larger distances.
The disk is optically thick in the vertical direction $\tau>1$ ) if the mass-loss rate is larger than For a given mass-loss rate the disk is optically thick close to the star, while becoming optically thin at larger distances.
For example. for a typical disk iass-loss rate required by the evolutionary calculations 102M,/vear+ (e.c.Ekstroimetal.2008b) the disk is optically thick even at large distances from the star rkzmLOR. for subsonic radial velocities.
For example, for a typical disk mass-loss rate required by the evolutionary calculations $10^{-5}\,\text{M}_\Sun/\text{year}^{-1}$ \citep[e.g.][]{eznula} the disk is optically thick even at large distances from the star $r\approx10^3\text{R}_\Sun$ for subsonic radial velocities.
Cousequeuthy. in realistic situations the disk is likely to be optically thick. at least close to the star. reseiubliug the “pscudophotosphere” of Be stars (e.g.Ixoubsk*etal.
Consequently, in realistic situations the disk is likely to be optically thick, at least close to the star, resembling the "pseudophotosphere" of Be stars \citep[e.g.][]{kouha}.
1997).. Contrary to very dense tot star winds (where the radiative flux comes from regions below the plotosphere). here we expect that the wiud from the optically thick disk starts to accelerate above the point where the disk optical depth is unity.
Contrary to very dense hot star winds (where the radiative flux comes from regions below the photosphere), here we expect that the wind from the optically thick disk starts to accelerate above the point where the disk optical depth is unity.
Numerical results show that the height of this poiut is comparable o the disk scale height Lf for moderate disk iiass-loss rates M.<10%Mfear1
Numerical results show that the height of this point is comparable to the disk scale height $H$ for moderate disk mass-loss rates $\dot M\lesssim 10^{-5}\,\text{M}_\Sun/\text{year}^{-1}$.