4 Chromospheric activity indicators

The echelle spectra analysed in this work allow us to study the behaviour of the different indicators from the Ca  II H & K to the Ca  II IRT lines, which are formed at different atmospheric heights. The chromospheric contribution in these features has been determined by using the spectral subtraction technique described in detail by Montes et al. (1995a) and Papers I, II, and III. The synthesized spectrum was constructed using the program STARMOD. Taken into account the stellar parameters derived in Sect. 3 we have used only a K5V primary component without any contribution from a secondary component. The inactive K5V stars used as reference stars are HD 154363 for the first observing run, and 61 Cyg A for the other three runs.

In Table 6 we give the excess emission equivalent width (EW) (measured in the subtracted spectra) for the Ca  II H & K, H$\epsilon$, H$\delta$, H$\gamma$, H$\beta$, H$\alpha$, and Ca  II IRT ($\lambda$8498, $\lambda$8542, $\lambda$8662) lines. When the emission features from both components can be deblended, we give the EW for the hot and cool (H/C) components. The uncertainties in the measured EW were estimated taking into account: a) the typical internal precisions of STARMOD (0.5-2 km s^-1 in velocity shifts, and $\pm$5 km s^-1 in $v\sin{i}$), b) the rms obtained in the fit between observed and synthesized spectra in the spectral regions outside the chromospheric features (typically in the range 0.01-0.03) and c) the standard deviations resulting in the EW measurements. The final estimated errors are in the range 10-20%.

The measured EWs given in Table 6 have been corrected for the relative contribution of each component to the total continuum determined by means of the radii and temperatures assumed in Sect. 3. For instance, in the H$\alpha$ line region the relative contributions are $S_{\rm H}=0.94$ for the hot component and $S_{\rm C}=0.06$ for the cool component, and the corrected EWs for the hot and cool components are obtained multiplying by a factor $1/S_{\rm H}$ and $1/S_{\rm C}$, respectively. Finally, these corrected EWs have been converted to absolute surface fluxes by using the empirical stellar flux scales calibrated by Hall (1996) as a function of the star colour index. In our case, we have used the B-V index and the corresponding coefficients for Ca  II H & K, H$\alpha$ and Ca  II IRT, using the same as Ca  II H & K for H$\epsilon$, and derived the H$\delta$, H$\gamma$and H$\beta$ fluxes by making an interpolation between the values of Ca  II H & K and H$\alpha$. The logarithm of the obtained absolute flux at the stellar surface (log $F_{\rm S}$) for the different chromospheric activity indicators is given in Table 7.

In Figs. 2 and 5 we have plotted for each observation in the H$\alpha$ and Ca  II IRT $\lambda$8498, $\lambda$8542 line region the observed spectrum (solid-line) and the synthesized spectrum (dashed-line) in the left panel, and the subtracted spectrum (dotted line), in the right panel. The observing run and the orbital phase ($\varphi$) of each spectrum are also given in these figures. The observed spectra in the Ca  II H & K line region are plotted in Fig. 4, and representative subtracted spectra in the H$\beta$, H$\gamma$ and H$\delta$ line regions are plotted in Fig. 3.

4.1 The H $\mathsf{\alpha}$ line

The H$\alpha$ line region is included in our spectra in the four observing runs. In all cases we have detected, in the observed spectra (see Fig. 2 left panel), strong H$\alpha$emission above the continuum coming from the primary component and a small H$\alpha$ emission coming from the secondary component. In all the spectra, except two which are very close to conjunction, we were able to deblend the emission coming from both components by using a two-Gaussian fit to the subtracted spectra (see Fig. 2 right panel).

The H$\alpha$ emission of the primary exhibits a central self-absorption similar to that observed in many M active stars (Stauffer & Hartmann 1986) and some K active stars like the dK5e binary V833 Tau (Montes et al. 1995b) and the K4V single star V834 Tau (Montes et al. 2001c). This self-absorption feature is a consequence of the line formation process in the chromosphere of very active stars (Houdebine & Doyle 1994).

The H$\alpha$ emission of BK Psc is persistent during the period of time covered by our observations (from 1999 to 2001). In addition, strong H$\alpha$ emission above the continuum from the primary component was also detected in previous spectra of this system taken in 1992 with EW(H$\alpha$) = 1.0 Å (Jeffries et al. 1995) and EW(H$\alpha$) = 1.1 Å (Mason et al. 1995). These EWs are lower than the EWs determined by us because these authors determined the EWs in the observed spectra and our EWs have been measured in the subtracted spectra, after eliminating the photospheric contribution. This persistent H$\alpha$ emission detected in BK Psc indicates that it is a very chromospherically active binary system similar to some RS CVn systems like V711 Tau, UX Ari, HU Vir, and DM UMa, and some BY Dra systems like BY Dra itself, and YY Gem, which always show H$\alpha$ emission above the continuum.

The detection of H$\alpha$ emission from the cool secondary component (M3V) of BK Psc indicates that this star has a very strong H$\alpha$ emission, since its photospheric contribution to the observed continuum is practically negligible. Strong H$\alpha$ emission is typical of the group of M type stars called dMe stars, some of which also show a scaled-up version of solar flares and are known as flare stars of UV Cet type stars. These latter stars are characterized by dramatic increases in the Hydrogen Balmer emission lines. However, this seems not to be the case for the secondary component of BK Psc since the H$\alpha$ emission we have detected has a similar intensity in the four observing runs.

The H$\alpha$ emission of the primary component shows small variations with the orbital phase, for instance, in the first run the EW changes from 1.7 to 1.1 Å. Seasonal variations are also detected, with larger EW(H$\alpha$) in 1999 than in 2001.

Figure 3: Subtracted spectra in the region of the H$\beta$, H$\gamma$, and H$\delta$ lines. Clear excess emission from the primary (P) and secondary (S) components is detected.

   

Table 6: EW of the different chromospheric activity indicators of BK Psc.

		EW(Å) in the subtracted spectrum
Obs.	$\varphi$$^{\star}$	CaII							CaII IRT
		K3	H3	H$\epsilon$3	H$\delta$	H$\gamma$	H$\beta$	H$\alpha$	$\lambda$8498	$\lambda$8542	$\lambda$8662
2.2 m 99	0.06	6.14/0.43	7.03/0.68	1.73	0.52	0.58	0.731	1.70/0.14	0.55	0.70	0.63
2.2 m 99	0.99	*	*	*	0.29	0.49	0.661	1.661	0.49	0.65	0.61
2.2 m 99	0.91	4.54/0.38	3.78/0.26	1.31	0.52/0.17	0.31/0.10	0.64/0.17	1.32/0.35	0.49	0.69	0.83
INT 00	0.42	-	-	-	*	0.17+	0.581	1.51/0.12	0.67	1.62	1.30
INT 00	0.88	-	-	-	*	0.27	0.50/0.05	1.32/0.24	0.59	1.36	1.09
INT 00	0.35	-	-	-	*	0.28	0.51/0.01	1.48/0.31	0.69	1.24	1.35
INT 00	0.81	-	-	-	*	*	0.54/0.08	1.37/0.32	0.64	1.36	1.42
NOT 00	0.52	2.04	-	-	*	0.098+	0.731	1.791	0.55	0.76	-
NOT 00	0.95	3.56/0.38	-	-	*	*	0.602	1.46/0.05	0.59	0.90	-
NOT 00	0.43	2.15	-	-	*	*	0.572	1.49/0.16	0.59	0.75	-
NOT 00	0.88	*	-	-	*	0.246+	0.55/1.13	1.32/0.31	0.53	0.89	-
2.2 m 01	0.17	1.88/0.37	2.31/4	1.014	0.50/0.35	0.50/0.13	0.45/0.04	1.18/0.32	0.51	0.65	0.48
2.2 m 01	0.58	3.57/0.31+	3.17/0.14+	0.91	0.27/0.03	0.36/0.12	0.52/0.11	1.09/0.37	0.46	0.64	0.51

$^{\star}$ Orbital phase calculated with the orbital period, $P_{\rm orb}$ and date of conjunction, $T_{\rm conj}$ determined in this paper.
* Data not measured due the very low S/N.
+ Data measured with low S/N.
¹ Data for the primary and secondary components not deblended due to the orbital phase of the observation.
² For these lines we can observe the two components but they could not be deblended.
³ These values have been measured at the observed spectra.
⁴ The H$\epsilon$ line of the primary component is blended with the Ca  II H line of the secondary.

   

Table 7: Absolute surface flux of the different chromospheric activity indicators of BK Psc.

		$\log F_{\rm S}$ in the subtracted spectrum
Obs.	$\varphi$$^{\star}$	CaII							CaII IRT
		K3	H3	H$\epsilon$3	H$\delta$	H$\gamma$	H$\beta$	H$\alpha$	$\lambda$8498	$\lambda$8542	$\lambda$8662
2.2 m 99	0.06	6.75/5.48	6.81/5.69	6.20	5.70	5.68	5.701	6.55/6.14	6.16	6.26	6.22
2.2 m 99	0.99	*	*	*	5.44	5.71	5.931	6.541	6.11	6.23	6.20
2.2 m 99	0.91	6.60/5.42	6.52/5.26	6.06	5.70/5.09	5.51/4,91	5.91/5.24	6.44/6.53	6.11	6.26	6.39
INT 00	0.42	-	-	-	*	5.25+	5.871	6.50/6.07	6.24	6.63	6.53
INT 00	0.88	-	-	-	*	5.45	5.80/4.69	6.44/6.37	6.19	6.55	6.45
INT 00	0.35	-	-	-	*	5.47	5.81/3.86	6.50/6.48	6.26	6.51	6.55
INT 00	0.81	-	-	-	*	*	5.84/4.91	6.46/6.50	6.22	6.55	6.57
NOT 00	0.52	6.26	-	-	*	5.01+	5.971	6.551	6.16	6.30	-
NOT 00	0.95	6.50/5.42	-	-	*	*	5.882	6.48/5.71	6.19	6.37	-
NOT 00	0.43	6.28	-	-	*	*	5.862	6.49/6.20	6.19	6.29	-
NOT 00	0.88	*	-	-	*	5.41+	5.85/6.07	6.44/6.48	6.14	6.37	-
2.2 m 01	0.17	6.22/5.41	6.31/4	5.954	5.68/5.41	5.72/5.02	5.76/4.62	6.39/6.50	6.13	6.23	6.10
2.2 m 01	0.58	6.50/5.34+	6.44/4.99+	5.91	5.41/4.34	5.58/4.98	5.82/5.05	6.39/6.56	6.08	6.22	6.13

Notes as in Table 6.

4.2 The H $\mathsf{\beta}$, H $\mathsf{\gamma}$ and H $\mathsf{\delta}$ lines

The other three Balmer lines included in all our spectra (H$\beta$, H$\gamma$ and H$\delta$) also show evidence of chromospheric activity. After applying the spectral subtraction, clear excess emission from both components is detected (see Fig. 3). When the S/N is high enough we have deblended the emission coming from both components by using a two-Gaussian fit to the subtracted spectra (see Table 6). The three lines show small seasonal and orbital phase variations with the same trend that the H$\alpha$ line.

We have also measured 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:

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

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 obtained for the primary component $\frac{E_{\rm H\alpha}}{E_{\rm H\beta}}$ in the range of 3 to 4 in all our spectra. These values indicate, according to Buzasi (1989) and Hall & Ramsey (1992), the presence of prominence-like material at the stellar surface.

Figure 4: Observed spectra in the region of the Ca  II H & K and H$\epsilon$ lines. The position of these lines for the primary (P) and secondary (S) components are marked with short vertical lines.

4.3 The Ca  II H & K and H $\mathsf{\epsilon}$ lines

The Ca  II H & K line region is included in the spectra of the FOCES 1999 and 2001 observing runs. Only the Ca  II K line is included in the NOT 2000 run, and this spectral region is not covered in the MUSICOS 2000 run. In all the spectra strong emission in the Ca  II H & K lines and a clear emission in the H$\epsilon$ line coming from the primary component is observed (see Fig. 4). The Ca  II H & K emission lines from the secondary component are also detected with relative wavelength shifts with respect to the primary in agreement with the wavelength shifts calculated in the H$\alpha$ and the other Balmer lines.

In our spectra the Ca  II H & K line spectral region is located at the end of the echellogram, where the efficiency of the spectrograph and the CCD decrease very rapidly, and therefore the S/N ratio obtained is very low, and the normalization of the spectra is very difficult. For these reasons we have not applied the spectral subtraction in this spectral region, and we have plotted in Fig. 4 only the observed spectra. The EWs measured in these spectra are in agreement with the strong Ca  II K emission (EW= 2.7 Å) reported by Mason et al. (1995).

Variations of the Ca  II H & K emissions with the orbital phase and from one epoch to another are observed for the primary component, and as in the case of the H$\alpha$ line the level of chromospheric activity is higher in the first observing run than in the last one.

4.4 The Ca  II IRT lines ( $\mathsf{\lambda}$8498, $\mathsf{\lambda}$8542, $\mathsf{\lambda}$8662)

The three lines of the Ca  II infrared triplet (IRT) are included in our echelle spectra, except $\lambda$8662 which is not included in the NOT 2000 observing run. In all the observed spectra of BK Psc a clear emission reversal is observed in the core of the Ca II IRT absorption lines (see Fig. 5 left panel). After applying the spectral subtraction technique, only the excess emission arising from the primary component is detected (see Fig. 5 right panel). Unlike the other chromospheric activity indicators, no evidence of excess emission is detected from the cooler secondary component, even when in this red spectral region its relative contribution is slightly larger than in the blue one. In addition, the orbital phase and seasonal variations detected in the Ca II IRT excess emission are more smooth than in the other activity indicators.

Figure 5: Observed and subtracted spectra, as in Fig. 2, in the region of the Ca  II IRT (8498, 8542 Å) lines.

We have calculated the ratio of excess emission EW, $\frac{E_{8542}}{E_{8498}}$, which is also an indicator of the type of chromospheric structure that produces the observed emission. In solar plages, values of $\frac{E_{8542}}{E_{8498}}$ $\approx$ 1.5-3 are measured, while in solar prominences the values are $\approx$9, the limit of an optically thin emitting plasma (Chester 1991). The small values of the $\frac{E_{8542}}{E_{8498}}$ ratio we have found for the primary component in all our spectra (ranging from 1.3 to 2.4) indicate that the Ca  II IRT emission of this star arises from plage-like regions, in contrast to the Balmer lines that seem to come from prominences. This markedly different behaviour of the Ca  II IRT emission has also been found in other chromospherically active binaries (see Paper III and references therein).