4. Discussion

4.1. The excess H emission

The inspection of the subtracted H spectra shows that in some stars the emission line profile has very broad wings, and is not well matched using a single-Gaussian fit. These profiles have therefore been fitted using two Gaussian components: a narrow component having a FWHM of 45- 90 km s^-1 and a broad component with a FWHM ranging from 133 to 470 km s^-1. In Table 4 (click here) we list the parameters (I, FWHM, EW) of the broad and narrow components. As can be seen in this table the average contribution of the broad component to the total EW of the line ranges from 35% in AR Psc to 77% in V711 Tau. We have observed this behaviour in the chromospheric H line only in the most active systems of the sample, the stars with intense H emission above the continuum (AR Psc, XX Tri, UX Ari, V711 Tau (Fig. 6 (click here)), II Peg (Fig. 19 (click here))) and also in systems with important excess H emission (MM Her, V815 Her, HK Lac). Furthermore, a revision of the H spectra of a large sample of chromospherically active binaries previously analysed by us (Montes et al. 1995a,b) indicates that the very active systems DM UMa and VV Mon also have broad components and that some systems also studied in this paper (V711 Tau, V815 Her, HK Lac) exhibit this behaviour at different epochs.

This parameterization of the subtracted H profile using a narrow and a broad component have been only reported until now for the RS CVn system DM UMa by Hatzes (1995) who suggests that the broad component could result from large-scale motions or winds in the chromosphere. Similar broad components have also been found in several transition region lines of the dM0e star AU Mic and the RS CVn systems Capella and V711 Tau using high-resolution UV observations obtained with the HST's GHRS (Linsky & Wood 1994; Linsky et al. 1995; Wood et al. 1996; Dempsey et al. 1996b,c; Robinson et al. 1996). The broad components in the transition region lines are interpreted by Linsky & Wood (1994) as arising from microflaring, because these broad profiles are reminiscent of the broad profiles observed in solar transition region explosive events, which are thought to be associated with emerging magnetic flux regions where field reconnection occurs.

The microflares are frequent, short-duration, energetically weak disturbances, i.e. they are the low-energy extension of flares, and therefore have large-scale motions associated that could explain the broad wings observed in these lines. The microflaring activity could occur not only in the transition region but also in the chromosphere of very active stars as indicates the detection of broad components in the chromospheric Mg II h & k lines of V711 Tau (Wood et al. 1996) but not in the chromospheric lines of the less active star Capella (Linsky et al. 1995).

Our detection of broad wings in the chromospheric H line of the most active systems of our sample allows us to conclude that microflaring occurs also in the chromosphere and that it is much more important in extremely active stars. Furthermore, within the group of stars that present this phenomenon a correlation between the contribution of the broad components and the degree of stellar activity seems to be present, as can be seen in Fig. 21 (click here) left-panel where we have plotted for each star the EW of the broad component versus the total excess H EW. This correlation was also noted by Wood et al. (1996) when compared the dominant broad component of V711 Tau with the smaller broad component of the less active stars AU Mic and Capella.

On the other hand, when we plot the FWHM of the broad component versus the total H EW (Fig. 21 (click here) right-panel) a general trend is not observed. However, in this case the relation appears for individual star, i.e. there is an increase in the FWHM when the activity level increases, which is consistent with the hypothesis that microflaring is responsible for the broad component emission.

We have also found that the larger changes in the excess H emission in the stars analyzed appear to occur predominantly in the broad component, as have been already noted in the case of DM UMa by Hatzes (1995). An extreme case of this behaviour is the strong change in the broad component that occurs during the flares detected in UX Ari and II Peg (see Figs. 20 (click here) and 21 (click here)). Since large scale mass motions do occur in solar flares, a large flare in these two systems may explain the increase of the H emission and the very broad wings observed in these spectra.

 

(Å) (Å) (%) (Å) (Å) (%)

0.443 0.132 4.473 0.625 44.5 0.534 1.358 0.772 55.0

0.519 0.108 3.362 0.388 32.0 0.587 1.319 0.823 68.0

0.524 0.107 3.652 0.417 33.0 0.603 1.316 0.845 67.0

0.438 0.297 2.917 0.922 65.6 0.344 1.316 0.483 34.4

0.576 0.145 5.153 0.778 50.8 0.468 1.512 0.753 49.2

0.736 0.405 5.414 2.229 67.8 0.580 1.723 1.069 32.4

0.261 0.425 4.563 2.046 81.1 0.344 1.254 0.459 18.2

0.280 0.417 4.406 1.947 80.4 0.344 1.265 0.463 19.1

0.606 0.318 5.176 1.707 64.6 0.509 1.635 0.887 33.6

0.641 0.402 4.063 1.735 72.8 0.406 1.486 0.642 27.0

V711 Tau (1992) 0.130 0.546 4.022 2.303 88.3 0.234 1.226 0.305 11.7

V711 Tau (1986) 0.200 0.570 4.489 2.660 87.8 0.251 1.381 0.369 12.2

0.260 0.576 4.513 2.570 85.2 0.311 1.486 0.448 14.8

V711 Tau (1988) 0.880 0.362 5.063 1.821 70.5 0.410 1.743 0.761 29.5

0.630 0.044 3.612 0.167 30.9 0.274 1.283 0.374 69.1

0.745 0.064 3.222 0.221 44.7 0.217 1.180 0.273 55.3

0.520 0.134 3.577 0.509 58.2 0.267 1.288 0.366 41.8

0.099 0.113 4.775 0.561 51.2 0.334 1.467 0.521 47.6

V815 Her (1989) 0.520 0.092 4.420 0.416 60.1 0.200 1.294 0.276 39.9

0.110 0.095 5.286 0.515 58.1 0.329 1.065 0.373 42.0

0.149 0.089 5.951 0.522 56.2 0.339 1.129 0.408 43.9

HK Lac (1989) 0.100 0.142 5.654 0.726 58.8 0.346 1.384 0.509 41.2

0.587 0.247 3.352 0.882 41.1 0.913 1.301 1.264 58.9

0.735 0.271 3.880 1.118 47.3 0.856 1.366 1.244 52.7

0.749 0.210 5.000 1.119 45.0 0.898 1.430 1.368 55.0

0.760 0.245 7.984 2.077 56.3 0.952 1.594 1.616 43.7

0.874 0.223 3.697 0.876 44.6 0.750 1.359 1.086 55.3

0.890 0.197 4.049 0.849 43.1 0.761 1.386 1.121 56.8

0.907 0.242 3.563 0.918 46.8 0.743 1.319 1.042 53.1

0.530 0.246 8.607 1.135 38.0 0.918 1.893 1.850 62.0

4.1.1. The He I D₃ line

The He I D₃ 5876 and He I 10830 triplets are known to be activity indicators in the Sun and late type stars (Zirin 1988; Shcherbakov et al. 1996). In the Sun, the He I D₃ line appears like an absorption feature cospatial with plages (Landman 1981), and it almost disappears when we look at the solar disk. This feature is also seen in absorption in surges, eruptive prominences, and weaker flares, whereas in emission in more intense flares (Zirin 1988).

The He I D₃ line is formed at middle chromosphere, and its correlation with X-ray flux and Ca  II H & K lines suggests that the fractional area of the stellar disk covered by plages may be a key factor in the formation of D₃ (Danks & Lambert 1985). Moreover it has been observed a slight rotational modulation in Cet (Lambert & O'Brien 1983) and ¹ Ori (Danks & Lambert 1985).

Historically, there have been basically two models to explain the line formation of He I D₃: (i) Zirin (1975) suggested that He I triplet levels were populated by over-photoionisation of the He I atoms by EUV and X-ray radiation, and subsequent radiative recombinations and cascade. (ii) Wolff et al. (1985) argued that collisional excitation and ionization in the chromosphere contributed also to the He I D₃ formation, and not only the EUV and X-ray radiation from the corona. However, the most recent models (Andretta & Giampapa 1995; Lanzafame & Byrne 1995) seem to indicate that the primary mechanism in the formation of the He I triplets is the collisional excitation and ionization (followed by recombination cascade) by electron impact.

The He I D₃ line usually appears, in stars, in absorption, but sometimes is in emission. There are two possible reasons: (i) Temperature and/or electronic density conditions are higher than ordinary, like may occur in flares (Zirin 1988; Andretta & Giampapa 1995; Lanzafame & Byrne 1995). (ii) As it has been seen in He I 10830, depending on the position of the emitting region in the disk or off the limb, the He I D₃ line would appear in absorption or emission. Since the He I 10830 is formed in emitting regions located at some distance from the stellar photosphere, when the emitting region is seen in projection against the stellar disk, He I 10830 line appears in absorption, and when the emitting region is observed off the stellar limb, the line is in emission (Simon et al. 1982; Wolff & Heasley 1984). These conclusions could extend to the case of He I D₃, since it is produced at the same region that He I 10830.

The He I D₃ line has been studied only in some chromospherically active binaries as II Peg (Huenemoerder & Ramsey 1987; Huenemoerder et al. 1990), DM UMa (Hatzes 1995), ER Vul (Gunn & Doyle 1997) and GK Hya (Gunn et al. 1997). The observation of emission in the He I D₃ line supports the detection of flare like events as in the case of II Peg (Huenemoerder & Ramsey 1987) and the weak-lined T Tauri star V410 Tau (Welty & Ramsey 1997).

In our spectra the He I D₃ line has been found in emission only during the flares of UX Ari and II Peg. We wish to emphasize that the detection of He I D₃ in emission in the RS CVn systems seems to occur at orbital phases near to the quadrature. In our observations we have detected He I D₃ in emission at orbital phase 0.74 in UX Ari (Montes et al. 1996b) and at 0.76 in II Peg. This line has been also observed in emission at orbital phases 0.22, 0.26, 0.77 in II Peg by Huenemoerder & Ramsey (1987) and Huenemoerder et al. (1990). Probably we are observing a flare off the limb, i.e. when the plage regions are near the limb (the active regions are preferably in the opposite faces of the stars), which is the most favourable situation to see an off the limb flare. But we cannot distinguish whether the emission is only due to the existence of the flare, or it is favoured by the relative position on the star.

The application of the spectral subtraction to our sample reveals that the He I D₃ line appears as an absorption feature more frequently in giants than in dwarfs. Three out of five giants observed show clear absorptions (BD Cet, V1149 Ori and HK Lac) and two of them exhibit a possible absorption (AY Cet and XX Tri), while among IV and V luminosity class stars there are only two plain absorptions. Various authors seem to point out a more frequent presence of He I 10830 and 5876 triplets in giants and supergiants than in dwarfs (Simon et al. 1982; Zirin 1982; Wolff & Heasley 1984).

Zirin (1982) observed He I 10830 usually in absorption, but sometimes it appears in emission, especially in giant and supergiants, with a P Cygni form, and he attributes it to a mass-ejection phenomenon (see also O'Brien & Lambert 1986). Simon et al. (1982) saw that none of the single red giants, in their sample, having strong 10830 absorption or emission has prominent transition region emission lines or soft X-ray emission, and they proposed a scattering process-like responsible for the 10830 line formation. Smith (1983) attributed a larger intensity in 10830 line for giants and supergiants to the most efficient ionization by EUV and X-ray radiation in atmospheres of coronally active giants. Other authors say that 10830 line is sometimes produced by the propagation of acoustic shock waves, or that He I 10830 transition represents a wind diagnostic. Some of the above proposed mechanisms could also be applied to the He I D₃ line.

4.1.2. The Na I D₁ and D₂ lines

The Na I D₁ and D₂ lines are collisionally-controlled in the atmospheres of late-type stars and are formed in the lower chromosphere. So, the detection of filled - in absorption in the D₁ and D₂ lines may provide information about chromospheric emission. (see the recent models of these lines for M dwarfs stars by Andretta et al. 1997).

In the Sun, Barrado et al. (1995) and Barrado (1996) have found changes in the EW of Na I lines in spectra taken at different regions over the solar surface, and a relation with the filled-in absorption H that might indicate that there is a non-radiative effect in the formation of these lines.

In other stars the D₁ and D₂ lines have been observed in emission or as a filled-in in very active red dwarf flare stars (Pettersen et al. 1984; Pettersen 1989). However, no systematic study of these lines has been performed in stars with different levels of activity, and in chromospherically active binaries only the negative and uncertain detection of filled-in in the few active binaries ER Vul and GK Hya, respectively, has been reported in the recent studies of Gunn & Doyle (1997) and Gunn et al. (1997).

The application of the spectral subtraction technique in these lines is more difficult that in the H line, because their wings are very sensitive to the effective temperature, mainly in latter spectral types. Therefore, small differences in spectral type, not appreciated in the H line, produce significant changes of the subtracted spectra in the wings of the Na I lines. Moreover, in this spectral region there is a large number of telluric lines, and in the spectra of some stars interstellar Na I could be present. However, the distances of the majority of the stars is lower than 50 pc and the effect of the interstellar Na I is negligible.

In spite of this problems, some conclusions can be drawn. In the chromospherically active binaries analysed here, the spectral subtraction reveals that the core of the Na I D₁ and D₂ lines are filled-in by chromospheric emission in the more active star of the sample (the star with H emission above the continuum, and with larger excess H emission EW). The stars with only a small or without excess H emission as BD Cet, AY Cet, V1149 Ori and KT Peg do not exhibit excess emission in the Na I lines. Moreover, the excess D₁ and D₂ emissions obtained are larger in the systems with larger excess H emission, and also increase in the flares observed in UX Ari and II Peg. In short, we can conclude that the filled-in of the core of the Na I D₁ and D₂ lines could be used as a chromospheric activity indicator.