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Astron. Astrophys. 318, 60-72 (1997)

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7. Discussion

7.1. The H [FORMULA] Emission Line

A significant result of these observations is that no obvious excess absorption feature has been detected for ER Vulpeculae which would indicate the presence of extended material surrounding the system. This confirms the results of Barden (1984), Newmark (1990), Lazaro & Arevalo (1994) and Eker, Hall & Anderson (1995). Mention should be made of the small excess absorption feature approximately 200 km s-1 blue-ward of the secondary star. This feature appears to be real but is so weak that analysis was impossible.

The observations have revealed a significant amount of excess emission in the core of the H [FORMULA] line for ER Vul. Analysis of the velocity of this component strongly suggests that the emission is arising predominantly on the secondary (G5 V) star although the variation in the FWHM indicates that there is a small amount of emission from the G0 V primary. The modulation of the excess emission with phase would indicate the appearance and disappearance of discrete active regions as different hemispheres of the stars are presented to the observer. No obvious variation of the emission has been detected for ER Vul over the phase of the observations. Although the phase coverage of the observations is small (0.14) and therefore inadequate to produce a significant variation, the velocity of the emission is highly suggestive that it is not associated with a transient phenomenon such as a flare or with localized individual active regions but arises due to a global phenomenon. Although it has been assumed that the velocity of the emission component is consistent with it arising on the secondary star there is clearly a consistent velocity shift of the emission by about 10 km s-1 blue-ward of the secondary position. It is interesting to note that in the solar case, observations in H [FORMULA] reveal systematic blue-shifts in the emission as a result of spicules (Beckers 1968; Beckers 1972). Spicules are rapidly changing thin filamentary features permeating the solar chromosphere at chromospheric temperatures (T [FORMULA] 104 K) which often extend upward into the hotter corona (T [FORMULA] 106 K). However these features cover only about 1-2% of the solar surface and consequently the disk-integrated H [FORMULA] solar spectrum does not display the systematic velocities associated with spicules. As will be demonstrated the plage coverage or filling factor for the ER Vul secondary appears to be very high (see below) so this observation may be of disk-integrated blue-shifts associated with active regions such as spicules.

In solar-type stars the radiation temperatures are sufficiently large or the electron temperatures are sufficiently low so that H [FORMULA] is dominated by photoionization; the Balmer lines are consequently coupled to the photospheric radiation field rather than the local line-formation region. For later stars or those with enhanced chromospheres the H [FORMULA] source function can become dominated by collisional processes as the electron density becomes higher and the photospheric radiation temperature becomes lower. In this case the H [FORMULA] core becomes filled-in and indicates the presence of a highly active chromosphere. Very high temperature chromospheres found in M-dwarfs actually drive the H [FORMULA] line into emission due to the dominance of collisional excitation in the line-forming region. The actual details of H [FORMULA] formation in late-type stars are extremely complex since many mechanisms and structures contribute to the line profile. However the observation of substantial excess emission in ER Vul is consistent with the general picture of an extremely active chromosphere (see for example Houdebine, Doyle & Koscielecki 1995).

For single stars the filling-in of the H [FORMULA] core has been well documented (Zarro & Rodgers 1983; Herbig 1985; Thatcher & Robinson 1993) and has also been observed in chromospherically active binaries (Strassmeier et al. 1990; Fernandez-Figueroa et al. 1994; Frasca & Catalano 1994; Eker, Hall & Anderson 1995; Montes et al. 1995). For such binaries the behaviour of the H [FORMULA] line is often found to be inconsistent with the photometric evolution and has been variously attributed to emission arising from star-spot regions or plages (Ramsey & Nations 1984; Newmark et al. 1990), from chromospheric network-like structures (Huenemoerder, Ramsey & Buzasi 1990) or from extended prominence-like material (Hall & Ramsey 1992b; Hall & Ramsey 1994). Unfortunately the present study does not encompass the variation of excess emission with phase since the primary goal was to detect the presence of extended material during the eclipse of the ER Vul system. However the analysis of the data has revealed evidence that the filling-in of the H [FORMULA] profile is due to plage regions. This is also confirmed by the appearance of the He I D3 absorption line in the spectrum (see below).

It is extremely difficult to quantify excess emission for a single star and most studies to date have involved surveys of large samples of active stars with a view to deciphering correlations between different chromospheric diagnostics and between these and other stellar parameters such as spectral type or rotational period (Young et al. 1989; Herbig 1985). Using subtracted spectra to derive physical parameters is dangerous since the technique is best suited to observing line variability in binaries across wide ranges in phase. However some simple calculations concerning the size of the emitting regions are possible as follows.

The continuum flux density in the H [FORMULA] region for ER Vul was calculated using the stellar atmosphere models of Kurucz (1979) for a G5 V star with solar abundances. The actual wavelength of the calculation was 6575Å and the effective temperature was taken as 5770 K. The resulting flux density is 2.14 106 erg s-1 cm-2 Å-1 which differs from the black-body flux at H [FORMULA] by only 3%. The corresponding luminosity at H [FORMULA] for ER Vul is 1.5 1029 erg s-1 Å-1 assuming a radius of 1.07 [FORMULA]. The mean equivalent width of the excess H [FORMULA] emission (0.596Å) was then converted to flux units using the relation;

[EQUATION]

where [FORMULA] is the H [FORMULA] continuum flux, [FORMULA] is the equivalent width of the excess emission and [FORMULA] was taken as 2Å, the width of the region used to define [FORMULA]. The resulting excess emission luminosity is 3.9 1029 erg s-1.

Fraquelli (1984) gives a relationship between the volume emissivity j, the electron density [FORMULA] and temperature of formation of the H [FORMULA] line T based on the assumption that the dominant emission mechanism is recombination. This mechanism was considered in detail by Burgess (1958). In case B of that analysis the assumption is made that the plasma is very opaque to Lyman line radiation and that the absorption from level 1 to level 2 in hydrogen is exactly balanced by the inverse spontaneous transition. The H [FORMULA] line radiation then results from cascades after electron capture and following absorption of Lyman line radiation. The approximation of Burgess (1958) is better than 1% for electron temperatures less than about 2.5 104 K. By comparing the emissivity variation in a non-LTE scaled VAL C model of solar plage regions for H [FORMULA] (Vernazza, Avrett & Loeser 1981), it has been shown (Andretta 1995, private communication) that within an order of magnitude the results of Fraquelli (1984) are valid. Hence the flux in the H [FORMULA] emission region is given by;

[EQUATION]

where V is the volume of the H [FORMULA] emitting region. The volume emissivity of H [FORMULA] is given by;

[EQUATION]

where [FORMULA] is the ionization energy of hydrogen. This equation is also based on the assumption that the proton and electron densities are equal ([FORMULA] = [FORMULA]) in this region of the chromosphere which is a reasonable assumption for a star of this spectral type. The same VAL C models show that the region of formation of the H [FORMULA] line core is at least a factor of three in electron density (1.0 1011 [FORMULA] [FORMULA] [FORMULA] 3.5 1010) and at least 50% in temperature (6000 [FORMULA] T [FORMULA] 10000). Taking the ionization energy of hydrogen as 2.18 10-11 erg then the corresponding range in the volume of the emission region is 1.6 1032 - 2.0 1033 cm3. Assuming the emission region is a homogeneous hemispherical shell then the chromospheric thickness for H [FORMULA] lies in the range [FORMULA] = 0.045 - 0.42 [FORMULA]. The chromospheric thickness reported for the Sun is of the order 0.004 [FORMULA] (Athay 1971) while Fraquelli (1984) shows by a similar method that the two components of HR 1099 (consisting of a G5 IV primary and a K1 IV secondary) have chromospheric thicknesses of 0.2 [FORMULA] and 0.06 [FORMULA] respectively. This may mean that the higher value is appropriate for ER Vul so that the main emission excess is formed at higher temperatures and lower electron densities than assumed by Fraquelli (1984). More detailed chromospheric modelling of the H [FORMULA] excess emission in active stars is required before any further definitive statements can be made.

7.2. The Ca II IRT emission lines

The observations of ER Vulpeculae have also revealed the presence of significant excess emission in the Ca II IRT line at [FORMULA] 8498. Unfortunately the remaining two lines in this triplet were just on the edge of the spectral coverage for these observations and were impossible to analyse. Analysis of the velocity of the excess emission has suggested that it arises almost equally on both components of the system (probably the secondary is slightly more active in these lines). The FWHM variation for ER Vul is also consistent with a dual emission peak. This is similar to observations of the BY Draconis-type variable DH Leonis studied by Newmark et al. (1990) in which the secondary was found to be more active in the Balmer lines while both stars showed almost equal Ca II IRT emission.

Enhanced emission cores in the Ca II H and K lines at [FORMULA] 3933.66 and [FORMULA] 3968.47 are the primary optical indicators of chromospheric activity. Their source functions are collisionally controlled and so these lines are sensitive probes of the electron density and temperature in the chromosphere. The H and K lines are favoured for chromospheric modelling since they are extremely important cooling mechanisms and their interpretation is relatively straight-forward. However less work has been done on the Ca II IRT although they are formed deeper in the atmosphere and are thus sensitive probes of the temperature minimum region and the temperature rise to the so-called Lyman plateau. Foing et al. (1989) observed the [FORMULA] 8498 and [FORMULA] 8542 IRT lines in a sample of stars from F9 to K4 and found that these lines correlate well with the H and K emission peaks. Linsky et al. (1979) also showed that filling-in of the IRT lines was a good indicator of stellar activity. The IRT lines are formed in the lower chromosphere of the Sun (Vernazza, Avrett & Loeser 1981) and in the temperature plateau in active chromosphere K2 dwarfs (Thatcher, Robinson & Rees 1991). In the quiet solar atmosphere these lines are simple absorption lines but as one goes to plages of brighter Ca II H and K emission the IRT line cores brighten and eventually develops self-reversed emission cores (Shine & Linsky 1972). On this basis it might then be expected that stars with Ca II H and K emissions comparable to those in solar plages will also exhibit IRT emission cores rather than simple filling-in. However this has not been found to be the case. Anderson (1974) surveyed 28 stars from F8 to M2 in the [FORMULA] 8498 line and Linsky et al. (1979) studied 49 stars from F9 to K3 at [FORMULA] 8542. Both studies revealed no distinct emission features even in the most active stars, although they did however display filling-in of the line cores. This behaviour is also displayed by the active components of ER Vul; on the evidence for plage regions on these stars we might expect to see emission cores in the Ca II IRT but instead see only filling-in. It has been suggested that such an emission core may be smeared out by rotation or large velocity fields and thus appear as filling-in (Linsky et al. 1979). However Thatcher & Robinson (1993) pointed out that rotational broadening could not account for the lack of emission cores in all cases. For ER Vul the rotation rate is probably not sufficient to smear the emission core beyond detection and so this star remains part of the Ca II IRT enigma. It can therefore be assumed that some additional broadening mechanism is at work in the plage emission from ER Vul. This is therefore another area where detailed chromospheric modelling is required.

Although the observation of excess emission in the IRT suggests the presence of non-radiative heating in the lower chromosphere for both components of ER Vul, and is consistent with plage emission, a reliable quantitative interpretation would be difficult without more detailed modelling.

7.3. The He I D3 absorption line

The observations have revealed obvious excess absorption in the He I D3 ([FORMULA] 5876) line for ER Vulpeculae. The use of He I D3 as an activity indicator has been largely ignored because it is extremely weak in normal stars and is generally blended with difficult water vapour lines when observed from low-to-medium altitude observatories. However observations of the stronger [FORMULA] 10830 line have been presented for large numbers of stars of different classes by Vaughan & Zirin (1968), Zirin (1976) and Zirin (1982) although observations of the D3 line are less common.

The triplet lines of He I at [FORMULA] 5876 and [FORMULA] 10830 appear in absorption in the solar spectrum; the weaker D3 absorption feature appears to be cospatial with plage regions and absent elsewhere. At [FORMULA] 10830 the absorption is strongest above active regions and very weak (with a tendency to be in emission) in coronal hole regions (McCabe & Mickey 1981). Landman (1981) studied high resolution spectra of He I D3 taken for solar plage regions. In stellar work Wolf & Heasley (1984) observed He I in 18 late-type stars and showed that the line depths, widths and the ratio [FORMULA] 10830/ [FORMULA] 5876 in dwarf stars was similar to that measured in solar plages. However they found that the line ratio in giants was much larger than in either solar plages or active dwarfs and suggested that the conditions under which He I forms may be very different in highly luminous stars. The factors which control the formation of the He I triplet lines, in particular D3 and [FORMULA] 10830, are however not well understood. The basic problem is that the transition region models based on UV and EUV resonance lines cannot account for the observed intensity levels in the quiet Sun He I resonance lines (Jordan 1975). Zirin (1975) suggested that the He I triplet levels are populated by radiative recombinations following photoionization of He atoms by coronal far-UV and X-ray line and continuum radiation. Giampapa et al. (1978) argued that He I D3 line they found in AD Leonis was not excited by recombination following photoionization since such an assumption would also require unrealistically high X-ray luminosity. Instead they suggest that He I [FORMULA] 5876 is excited by collisions from the ground state in the hotter (T [FORMULA] 8000 K) region of the stellar chromosphere. Recently Andretta, Giampapa & Jones (1995) suggest on the basis of their non-LTE calculations that at least some of the He I spectral features in the Sun and late-type stars could be linked to regions of enhanced UV and X-ray emission but that He I formed by photospheric ionization and decay and collisional processes seemed to be present even for the case of an inactive corona.

Despite the confusion as to the details of He I formation in the Sun and stars it is now almost certain that active regions on the stellar surface are the dominant area contributing to the observed flux profiles. The equivalent width of D3 is controlled by the temperature-density profile in the middle chromosphere and the fraction of disk covered by plages. The observation of He I D3 absorption in ER Vul is therefore very suggestive of large areas of plage-like plasma in the chromospheres of one or both components in this binary. The equivalent widths of the absorption remain fairly constant with a mean value of 78 mÅ. The velocity measurements of the excess absorption are however confused for all measurements and it can only be assumed that this is due to the weakness of the line and noise effecting the gaussian fitting. It is a reasonable assumption that the majority of the absorption is occurring on the secondary component of the binary since this appears to be the more active star with a higher probability of plage regions. The FWHM variation however clearly implies that absorption is occurring for both stars.

Andretta & Giampapa (1995) presented a computational approach to the He I triplet lines in dwarf F and G stars. They show that the lines can be utilised to infer the fractional area coverage or filling factor of active regions on stellar surfaces if the intrinsic absorption strengths are known. They compute He I profiles using a grid of model chromospheres superposed on late-type dwarf photospheric models. They estimate that in both F and G type stars a conservative value for the maximum absorption equivalent width ([FORMULA]) in He I D3 is [FORMULA] 100-150 mÅ. This calculation ignores the quiet component and is therefore appropriate for measurements on the subtracted spectra. A lower limit to the filling factor is then given by;

[EQUATION]

where [FORMULA] is the observed He I equivalent width. Using this relationship the filling factor for the secondary component of ER Vul is 0.53. This is extremely high even if we apportion this value between the two blended components. Unfortunately there are no measurements of the He I [FORMULA] 10830 line for this system which could be used to verify this result. However Andretta & Giampapa (1995) report a filling factor of between 0.5 and 0.6 for the active G8 V star [FORMULA] Boo A while O'Neal, Neff & Saar (1994) deduced a filling factor of between 0.41 and 0.52 for II Peg. So although ER Vul has a high filling factor this is not inconsistent with other observations of highly active stars (Wolf & Heasley 1984).

Andretta & Giampapa (1995) also computed the theoretical dependency of the [FORMULA] 5876 equivalent width with the excess emission equivalent width in the core of the H [FORMULA] line from plage regions, and these predictions showed good agreement with numerous observations of solar plage regions. The results for ER Vul show that the He I D3 and H [FORMULA] measurements are again entirely consistent with both observational and theoretical considerations of active plage regions in the solar analogy. In conclusion the He I D3 observations for ER Vul have demonstrated the existence of very large filling factors presumably associated with the secondary component. ER Vul appears to be highly active with regards to plage regions and the existence of such slab-like structures is entirely consistent with the lack of detection of extended material around the system.

7.4. The Mg I b and Na I D1 & D2 lines

Strong neutral metal absorption lines are of special interest for stellar activity studies since they are formed in the lower chromosphere and the region of temperature minimum for solar type stars. Basri, Wilcots & Stout (1989) concluded in their study of 29 main-sequence stars that the Mg I b lines were good diagnostics of photospheric activity. The Na I lines are collision dominated due to the small photoionization cross-section and relatively large collisional cross-section for Na I (Johnson 1964). This means that the source functions of the Na I lines are functions of electron temperature and pressure and so are good indicators of changes in the lower chromosphere.

The observations have revealed a significant amount of excess emission in the Mg I b lines from ER Vulpeculae. Measurements have been made on the lines at [FORMULA] 5172.68 and [FORMULA] 5183.61 while the line at [FORMULA] 5167.33, although present, was too weak to give reliable measurements. The velocities of the two measured lines show that the emission is again predominantly from the secondary component although the FWHM variation is consistent with a small amount of emission from the primary. In contrast, no excess emission has been detected in the Na I D1 ([FORMULA] 5889.95) and D2 ([FORMULA] 5895.92) lines in the observations of ER Vul. This is significant since both these neutral metals form in the lower chromosphere. Furthermore the ionization energies of Mg I and Na I are 7.6 eV and 5.1 eV respectively indicating a disparity between the observed emission and that expected from simple energetics. We regard it unlikely that crowded atmospheric absorption could account for a difference in EW between the standard and target stars. An explanation may lie in the existence of a complex chromospheric structure and/or opacity effects.

Further theoretical work is required before any definitive statements about the lack of excess Na I emission in ER Vul can be made. Current work underway by Andretta, Doyle & Byrne (1996) using a more complete set of opacities will provide a better insight into the use of these lines in studies of stellar activity.

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© European Southern Observatory (ESO) 1997

Online publication: July 8, 1998
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