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Astron. Astrophys. 357, 951-956 (2000)

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3. Spectral interpretation

RW Aur spectrum in the 2777-2823 Å band is shown in Fig. 1. The continuum flux was derived from a featureless part of the spectrum between 2810 Å and 2822 Å: [FORMULA]. We identified the two absorption features near 2781 Å as Fe II 2779.30 and 2783.69 lines of the [FORMULA] multiplet (uv 234). The gaussian fit of these lines is shown in the insert in the right corner, as well as the continuum level that appears here [FORMULA] above [FORMULA] value. The lines are blueshifted - the respective shift in km s-1 is shown in the insert. Both lines have a residual intensity near 0.5, i.e. they are strong enough, while the excitation energy of their low levels is [FORMULA] eV.

[FIGURE] Fig. 1. RW Aur spectrum in 2778-2822 Å spectral band with the line identification. The signed radial velocities are in km s-1. The dashed line shows the gaussian fit of the Mg II h line. The fits of the Fe II line (uv 234 multiplet) are shown in the insert at top right corner

We conclude therefore that the temperature and electron density is high in the line forming region. Therefore we connect the observed depression in the red wing of Fe II 2779.30 line with the two lines of the Mg I uv 6 multiplet: their low levels belong to the first excited term [FORMULA] from which downward transitions to the ground level [FORMULA] have low probabilities. The expected population of the [FORMULA] term levels is high enough to form subordinate absorption lines, and the relative populations of the [FORMULA] term levels are close to the ratio of their statistical weights. Then we conclude, from the uv6 multiplet gf-values, that the 2779.83 and 2781.42 Å lines should be the strongest within the multiplet.

One more absorption line was identified in the spectrum: Fe I 2788.10. This line should be the strongest of the [FORMULA] multiplet, assuming, as above, that the levels of the [FORMULA] term are populated proportionally to their statistical weights (Nave et al. 1994). The blue wing of the line is disturbed by an emission feature (see below), nevertheless the line center is blueshifted some tens of km s-1.

The Mg II h and k lines have profiles similar and look very unusual. We fitted the Mg II h 2802.71 line wings by a gaussian, with the center at [FORMULA] km s-1 relative to the rest frame - see Fig. 1. This value is in good agreement with the radial velocity of RW Aur found from optical spectra by Hartmann et al. (1986): [FORMULA] km s-1. Therefore this broad [FORMULA] km s-1) emission component originates at the star surface. The iron and probably the Mg I absorption lines are blueshifted relative to the star of [FORMULA] km s-1, so we assume they form in the stellar wind.

The narrow feature near the zero velocity position in the Mg II h and k lines is interstellar (IS), but obviously one more absorption component is present in the blue wing of the emission lines - apparently it forms in the stellar wind. Thus, the observed Mg II h and k line profiles are the result of the superposition of IS and wind absorption features onto the wide symmetrical stellar emission component - see also Imhoff & Appenzeller (1989). Unfortunately we cannot restore the blueshifted feature profile due to the lack of information from the "dead" diode. The answer crucially depends on the shape of the self-absorption feature in the central part of the stellar emission line.

Emission components of the Mg II h and k lines are so strong that their upper levels [FORMULA] have a large population. We identified the 2791 Å emission feature with the Mg II 2790.8 line of the uv 3 multiplet: its lower level just belongs to the [FORMULA] term. But stellar Mg II h and k lines can also pump the [FORMULA] term levels in stellar wind matter. We conclude therefore that the two emission features near 2790 Å are not two lines, but a wide stellar 2790.8 Å emission line with a wind absorption feature in the blue wing - similar to the Ca II 8542 line profile of the star in Fig. 1d of Muzerolle et al. (1998). Minima of the absorption features are shifted relative to the central wavelength by [FORMULA] and [FORMULA] km s-1 in the 2790.8 Å and Mg II k lines respectively. It is a good agreement, bearing in mind that the 2790.8 Å line is a subordinate one, so it has no IS absorption feature. Two other lines of the uv 3 multiplet (2797.9 and 2798.0 Å) fall in the red wing of the Mg II k line, so its profile is significantly disturbed from both sides. For this reason we could not successfully fit the Mg II k line profile with a gaussian.

We explain now why only the above mentioned lines of Fe II are present in the spectrum of RW Aur, while there are more than 50 lines of the ion in the [FORMULA] Å spectral band. Assuming that the emitting region is homogeneous and the local broadening of the Fe II lines is due to thermal motion, we write the optical depth in the center of a line as follows (Mihalas 1978):


where [FORMULA] and [FORMULA] - are the oscillator strength and wavelength (in Å) of the transition; A - is the ion atomic weight (56 for Fe); [FORMULA] - is the gas temperature in [FORMULA] K units; [FORMULA] - the particle density of the [FORMULA]th energy level [FORMULA] L - the extension of the absorbing region along the line of sight (in cm).

The lowest odd energy level in Fe II is [FORMULA], with [FORMULA] eV, so levels lower than this have no allowed radiative transitions between them. The populations of these levels, including that of the [FORMULA] term, are therefore likely governed by electron collisions. They probably have a Boltzmann (LTE) population at the wind electron temperature [FORMULA] (McMurry et al. 1999). Then the ratio of optical depths of the two Fe II lines is:


Adopting gf and E values from Nahar (1995) we have found from Eq. (2), that the optical depths of all other Fe II lines between 2777 Å and 2823 Å are 3 times less than that of the two observed lines of the uv 234 multiplet if [FORMULA] and 5 times less if [FORMULA]. It means that all other Fe II lines will be too weak to be observed if [FORMULA] K, so the wind electron temperature is somewhere below this value.

From atomic data of the Fe I by Nave et al. (1994) and under the same conditions (LTE population and [FORMULA] the Fe I 2788.93 line of the uv 44 multiplet should have an optical depth at least 2 times larger than other Fe I lines in the spectral band of Fig. 1. The line is weak enough to explain the absence of other neutral iron lines in the spectrum.

The RW Aur spectrum in the vicinity of the C IV 1550 doublet is presented in Fig. 2. To understand this very puzzling spectrum we compared the optical depth of Fe II lines within the spectral band of Fig. 2, with that of the Fe II 2783.69 line. There are many lines with [FORMULA] comparable or even larger than [FORMULA] so they should be strong enough. But the continuum is underexposed in the spectrum and we cannot observe most of the respective absorption features. On the other hand absorption lines, whose wavelengths are close to the C IV doublet ones, should superimpose onto the strong C IV emission lines and therefore can be seen. Information on Fe II 1546 Å [FORMULA]1552 Å lines are presented in Table 2. The last three columns show the ratio [FORMULA] for three values of [FORMULA] 0.6, 0.8 and 1.0, only lines with [FORMULA] were selected. Note that semiforbidden line gf-values were adopted by Smith et al. (1996) - they are less accurate than those of the allowed transitions (Nahar 1995). It follows from Table 2 that the 1548.67 Å and 1550.27 Å lines are expected to be strong and thus should disturb the profiles of the underlying C IV emission lines. The disturbing effect of other lines from Table 2 can be significant (see Fig. 2), but the S/N ratio of the spectrum is too poor to estimate the role of each line. We only note that the 1548.67 Å line becomes strong enough only if the wind temperature is below [FORMULA] K, in agreement with the absence of emission lines of C I uv 3 resonance multiplet near [FORMULA] Å.

[FIGURE] Fig. 2. RW Aur spectrum in 1532-1568 Å spectral band


Table 2. Relative optical depth of Fe II lines, which can disturb the profiles of C IV 1550 and Si IV 1400 lines

There are no strong Fe I lines between 1532 Å and 1568 Å but there are many lines of Si I , which originate from fine structure levels of the ground term. It may be that the Si I 1548.72 Å and 1551.23 Å lines disturb the profile of the C IV doublet, but we should know the relative abundance of the silicon atoms in the wind in order to quantitatively estimate the effect.

Additionally, the C IV emission line profiles are disturbed by lines of molecular hydrogen [FORMULA] Two molecular lines can be successfully identified in the spectrum: R(3) 1547.33 and P(5) 1562.39. Probably they originate close to the RW Aur H-H objects (Mundt & Eisl"offel 1998) due to [FORMULA] pumping by the stellar H [FORMULA] line. P(5) line looks double peaked: the red component is 130 km s-1 shifted relatively to the blue one. The latter is practically at stellar rest position: its heliocentric radial velocity is close to +13 km s-1. Possibly one is observing [FORMULA] emission from the RW Aur jet and counterjet discovered by Hirth et al. (1994). On the other hand these authors found that gas velocity varies in a very wide range both in the jet and counterjet.

The intensity ratio of R(3) and P(5) lines is [FORMULA] in spectra of T Tau (Brown et al. 1984), RU Lup (Lamzin 1999) and [FORMULA] Tau (McMurry et al. 1999). Apparently they originate from the same upper level [FORMULA] [FORMULA] [FORMULA] and have almost equal transition probabilities (Abgrall et al. 1993). We suppose that the strong emission peak near 1547.5 Å does not only represent the [FORMULA] R(3) line, but it is blended with the C IV 1548.20 line.

McMurry et al. (1999) observed many Fe II emission lines in [FORMULA] Tau spectra excited via pumping by the H [FORMULA] line. Probably 1534.84 Å [FORMULA] and 1539.05 Å [FORMULA] lines from their Table 2 are also present in the RW Aur spectrum. The pumping of these lines occurs from levels of the [FORMULA] term, whose excitation energy is [FORMULA] eV. The electron temperature in the region(s) where [FORMULA] lines form does not significantly exceed 2 000 K, which seems to be too low to produce a large enough population of the [FORMULA] term levels. Therefore the observed Fe II fluorescent lines originate in the warm wind along with the absorption lines. Profiles of the emission and absorption lines should differ significantly, because we observe emission from all regions of the wind above the accretion disk, while absorption only occurs along the line of sight. Furthermore the [FORMULA] ions cannot absorb molecular hydrogen emission because it forms far outside the warm wind.

Now we consider the RW Aur spectrum in the [FORMULA] Å spectral band (Fig. 3). The profiles of the Si IV 1393.8 and 1402.8 lines differ significantly, so we have checked if this can be due to the superposition of stellar wind absorption features. There are three Fe II lines near the Si IV doublet with relatively large optical depth (see Table 2), but the Ni II 1393.32 resonance line is much stronger. Iron and nickel are predominantly singly ionised in the wind regions where the Fe II absorption lines originate - see below. If the Ni II even levels below 4 eV have a LTE population (due to the same reason as the Fe II ones - see above) then [FORMULA] when [FORMULA] K, with [FORMULA] (Allen 1973) and [FORMULA] (Smith et al. 1996). Thus the Ni II 1393.32 wind absorption line almost completely cancels out the blue wing of the stellar Si IV 1393.8 emission line.

[FIGURE] Fig. 3. RW Aur spectrum in 1383-1419 Å spectral band

We identify an emission feature near 1399 Å as the [FORMULA] Å fluorescent P(2) line of molecular hydrogen. In contrast to the P(5) 1562 Å line, the P(2) 1399 Å line is not double peaked. It may be the result of the different pumping conditions of these lines (Lamzin 1999) or/and gas parameters in the jet and counterjet (Hirth et al. 1994). The R(0) 1393.72 Å line originates from the same upper level as the P(2) 1399 Å one - [FORMULA] [FORMULA] [FORMULA] - and has [FORMULA] times less transition probability (Abgrall et al. 1993). The contribution of the R(0) line to the emission in the blue wing of the Si IV 1394 line should be relatively small if R(0) and P(2) lines are optically thin.

There are two more lines of [FORMULA] which look much more strong: R(1) 1393.96 Å and P(3) 1402.65 Å. They originate from the common upper level [FORMULA] [FORMULA] [FORMULA] and the ratio of their transition probabilities is near 1.4. These lines give about 50% contribution to the emission in the 1393 Å and 1403 Å features, previously attributed only to Si IV lines.

We identify the bump in the blue wing of the Si IV 1403 line with the O IV ] 1401.16 line. It should be at least two times stronger than other lines of the O IV [FORMULA] multiplet in CTTS spectra (Lamzin & Gomez de Castro 1998). Probably an analogous [FORMULA] Å semiforbidden line of the S IV [FORMULA] multiplet is also present in the RW Aur spectrum.

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Online publication: June 5, 2000