3 Spectroscopic features analysed

3.1 Chromospheric activity indicators

The echelle spectra analysed in this paper allow us to study the behaviour of the different optical chromospheric activity indicators formed at different atmospheric heights: Na  I D₁, D₂ and Mg  I b triplet (upper photosphere and lower chromosphere), Ca  II IRT lines (lower chromosphere), H$\alpha$, H$\beta$, Ca  II H & K (middle chromosphere) and He  I D₃ (upper chromosphere). The chromospheric contribution in these features has been determined using the spectral subtraction technique described in detail by Montes et al. (1995a,b,c), (see Papers I and II). The synthesized spectrum was constructed using the program STARMOD developed at Penn State (Barden 1985). The inactive stars used as reference stars in the spectral subtraction were observed during the same observing run as the active stars or were taken from our libraries of late-type stars (see Montes 1998). We have determined the excess emission equivalent width (EW) (measured in the subtracted spectra) and converted to absolute chromospheric flux at the stellar surface. We have estimated the errors in the measured EW taking into account the typical internal precisions of STARMOD (0.5 - 2 km s^-1 in velocity shifts, $\pm$5 km s^-1 in $V\sin{\it i}$, and 5% in intensity weights), the rms obtained in the fit between observed and synthesized spectra in the regions outside the chromospheric features (typically in the range 0.01-0.03) and the standard deviations resulting in the EW measurements. The estimated errors are in the range 10-20%. For low active stars errors are larger and we have considered as a clear detection of excess emission or absorption in the chromospheric lines only when these features in the difference spectrum are larger than 3 $\sigma$.

Table 3 gives the H$\alpha$ line parameters, measured in the observed and subtracted spectra of the sample. The Col. 3 gives the orbital phase ($\varphi$) for each spectrum. In the Col. 4, H and C mean emission from the hot and cool components, respectively, and T means that at these phases the spectral features cannot be deblended. The Col. 5 gives the relative contribution of the hot and cool components to the total continuum ($S_{\rm H}$ and $S_{\rm C}$), respectively. The Col. 6 describes the observed H$\alpha$ profile, i.e. whether the line is in absorption (A), in emission (E) or totally filled-in by emission (F). The Cols. 7, 8 and 9 give the following parameters measured in the observed spectrum: the full width at half maximum ( $W_{\rm obs}$), the residual intensity ($R_{\rm c}$ = $\frac {F_{\rm c}}{F_{\rm cont}}$) and the equivalent width (EW). The last four columns give the following parameters measured in the subtracted spectrum: the full width at half maximum ( $W_{\rm sub}$), the peak emission intensity (I), the excess H$\alpha$ emission equivalent width ( $EW({\rm H}\alpha$)), corrected for the contribution of the components to the total continuum and the logarithm of the absolute flux at the stellar surface (log$F_{\rm S}$(H$\alpha$)) obtained with the calibration of Hall (1996) as a function of (V-R).

In Table 4 we list the parameters ( I, FWHM, EW) of the broad and narrow components used in the two Gaussian-component fit to the H$\alpha$ subtracted emission profile. We have performed this fit in the stars that show broad wings. See the comments for each individual star in Sect. 4 and the interpretation of these components given in Sect. 5.

Table 5 gives the H$\beta$ line parameters, measured in the observed and subtracted spectra, as in the case of the H$\alpha$ line. In this table we also give the ratio of excess emission EW in the H$\alpha$ and H$\beta$ lines, $\frac{EW({\rm H}\alpha)}{EW({\rm H}\beta)}$, and the ratio of excess emission $\frac{E_{{\rm H}\alpha}}{E_{{\rm H}\beta}}$with the correction:

$$\frac{E_{{\rm H}\alpha}}{E_{{\rm H}\beta}} = \frac{EW({\rm H}\alpha)}{EW({\rm H}\beta)}*0.2444*2.512^{(B-R)}$$

given by Hall & Ramsey (1992) that takes into account the absolute flux density in these lines and the color difference in the components. We have used this ratio as a diagnostic for discriminating between the presence of plage-like and prominence-like material at the stellar surface, following the theoretical modelling by Buzasi (1989) who found that low $E_{\rm H\alpha}/E_{\rm H\beta}$ ( $\approx ~1-2$) can be achieved both in plages and prominences viewed against the disk, but that high ratios ( $\approx ~3-15$) can only be achieved in extended regions viewed off the limb. The study of chromospherically active binaries by Hall & Ramsey (1992) has demonstrated the presence of large amounts of extended, prominence-like material in these stars.

We also analyse the possible filling-in of the core of the Na  I D₁ and D₂ lines as other chromospheric activity indicator as well as the behaviour of the He  I D₃ line, which can be in absorption, filled-in due to frequent low-level flaring or in emission due to flares (see Papers I and II; Saar et al. 1997; Montes et al. 1996b, 1999).

Table 6 gives the Ca  II H & K and H$\epsilon$ lines parameters, measured in the observed and subtracted spectra. In Cols. 5 and 6 we list the EW for the K and H lines, obtained by reconstruction of the absorption line profile (described by Fernández-Figueroa et al. 1994, hereafter FFMCC). In Cols. 7, 8 and 9 we give the EW for the K, H and H$\epsilon$ lines using the spectral subtraction technique (explained by Montes et al. 1995c, 1996a) and corrected for the contribution of the components to the total continuum. In Cols. 10, 11 and 12 we list the corresponding logarithm of the surface flux obtained by means of the linear relationship between the absolute surface flux at 3950 Å$\$ (in erg cm^-2 s^-1 Å^-1) and the colour index (V-R) by Pasquini et al. (1988).

Figure 1: H$\alpha$, H$\beta$, Ca  II H & K, and Ca  II IRT spectra of UX Ari. Observed and synthetic spectra in the left panel and subtracted spectra in the right panel

Figure 2: H$\alpha$, H$\beta$, Ca  II H & K, and Ca  II IRT spectra of 12 Cam. Observed and synthetic spectra in the left panel and subtracted spectra in the right panel

Figure 3: H$\alpha$, H$\beta$, Ca  II H & K, and Ca  II IRT spectra of V1149 Ori. Observed and synthetic spectra in the left panel and subtracted spectra in the right panel

Figure 4: H$\alpha$ and H$\beta$ spectra of OU Gem. Observed and synthetic spectra in the left panel and subtracted spectra in the right panel

Figure 5: Ca  II H & K and Ca  II IRT spectra of OU Gem. Observed and synthetic spectra in the left panel and subtracted spectra in the right panel

Figure 6: H$\alpha$, H$\beta$, Ca  II H & K, and Ca  II IRT spectra of $\sigma$ Gem. Observed and synthetic spectra in the left panel and subtracted spectra in the right panel

Figure 7: H$\alpha$ observed (left panel) and subtracted (right panel) spectra of BF Lyn

Figure 8: H$\alpha$ and Ca  II IRT $\lambda$8542 EW of the hot (left panel) and cool (right panel) components of BF Lyn for the McD98 run versus the orbital phase

Figure 9: H$\beta$ observed (left panel) and subtracted (right panel) spectra of BF Lyn

Figure 10: Ca  II H & K observed (left panel) and subtracted (right panel) spectra of BF Lyn. We have also plotted the Gaussian fit to the subtracted spectrum used to deblend the contribution of the hot (H) and cool (C) components (short- and long-dashed lines)

Figure 11: Ca  II IRT $\lambda$8542 and $\lambda$8662 observed (left panel) and subtracted (right panel) spectra of BF Lyn

Figure 12: H$\alpha$, H$\beta$, Ca  II H & K, and Ca  II IRT spectra of IL Hya. Observed and synthetic spectra in the left panel and subtracted spectra in the right panel

Figure 13: H$\alpha$, H$\beta$, Ca  II H & K, and Ca  II IRT spectra of FG UMa. Observed and synthetic spectra in the left panel and subtracted spectra in the right panel

Figure 14: H$\alpha$, H$\beta$, Ca  II H & K, and Ca  II IRT spectra of LR Hya. Observed and synthetic spectra in the left panel and subtracted spectra in the right panel

Table 7 gives the Ca  II IRT lines parameters, measured in the observed spectra by reconstruction of the absorption line profile and using the spectral subtraction. The columns of this table have the same meaning as in Table 6 for the Ca  II H & K lines. The absolute fluxes at the stellar surface have been obtained using the calibration of Hall (1996) as a function of (V-R). For the observing runs in which the $\lambda$8542 and $\lambda$8498 lines are included we also give the ratio of excess emission EW, E₈₅₄₂/E₈₄₉₈, which is also an indicator of the type of chromospheric structure that produces the observed emission. In solar plages, values of $E_{8542}/E_{8498} \approx ~1.5-3$ are measured, while in solar prominences the values are $\approx$ 9, the limit of an optically thin emitting plasma (Chester 1991). However, the observations of active stars (Chester et al. 1994; Arévalo & Lázaro 1999) indicate that these lines exhibit markedly different behaviour. The E₈₅₄₂/E₈₄₉₈ ratios found in these stars are smaller (closer to the optically thick value of one) than solar plages. These values indicate that the Ca  II IRT emission arises predominantly in chromospheric plages.

3.2 The Li  I $\lambda$6707.8 line

The resonance doublet of Li  I at $\lambda$6708 Å is an important diagnostic of age in late-type stars since it is destroyed easily by thermonuclear reactions in the stellar interior. It is well-known that a large number of chromospherically active binaries shows Li  I abundances higher than other stars of the same mass and evolutionary stage (Barrado et al. 1997, 1998; Paper I, Montes & Ramsey 1998). This line is only included in our echelle spectra in the McD98 and INT99 observing runs. In Fig. 25 we have plotted representative spectra of OU Gem, BF Lyn and HU Vir in this spectral region. A K1III reference star with some photospheric lines identified has been also plotted in order to better identify the expected position of the Li  I line. It was only possible to measure the equivalent width of the Li  I absorption line in the SB1 system HU Vir. In the case of the SB2 systems OU Gem and BF Lyn the possible small absorption Li  I of one or both components is blended with photospheric lines of the other component.