4. Chromospheric activity

The 8 hypergiants that have been observed show that the Ca II line profiles have a complex structure. Emission components appear within the absorption profile, perhaps similar to the ones that are observed in the Ca II H and K profiles of stars known to have chromospheric activity. However, the rest of the stars, i.e., the dwarfs, giants and supergiants show only absorption profiles in the Ca II triplet lines. Anderson (1974) observed the 8498 line in 28 stars of spectral types F8-M2 at a resolution of 0.28 Å and found no distinct emission in any star. Only one star, Dra showed a central residual intensity (CRI) larger than in the other stars, possibly due to filling in by chromospheric emission. This star is the only one with a very active chromosphere in Anderson's list (intensity class 3 according to Wilson 1976). Linsky et al. (1979) also observed 8542 profiles at a resolution of 0.14 Å in 49 stars as a probe for chromospheric activity. Their resolution is twice as high as that of Anderson and is sufficient to resolve emission features within the absorption profiles of the Ca II triplet in stars if they are present. No evidence of emission was found even for the most active chromosphere star.

 

Star Spectral I(CaII K) CRI EQW type

HR 2269 K3 Ib - - 0.32 0.25 0.30 1.59 3.40 3.36

3 Cet K3 Ib - <17 0.18 0.12 0.20 2.11 4.66 3.63

HR 2269 K3 Ib - - 0.32 0.25 0.30 1 59 3.40 3.36

K3 Ib 4 0.17 0.12 0.08 2.52 4.59 3.96

56 Peg G8 Ib 4 <17 0.45 0.35 0.24 1.50 3.43 2.85

K1.5 Ib 3 <17 0.16 0.11 0.10 2.60 5.17 4.51

56 Peg G8 Ib 4 <17 0.45 0.35 0.24 1.50 3.43 2.85

G8 Ib 4 <17 0.19 0.15 0.11 2.62 4.71 4.02

G2 Ib-IIa 3 13 0.37 0.27 0.31 1.83 4.40 3.43

G0 Ib - 0.28 0.15 0.15 1.76 4.66 4.35

M3 III 4 - 0.33 0.20 0.17 1.59 3.37 3.05

M3 III 5 - 0.28 0.14 0.18 1.85 3.89 3.36

K4 III 3 <19 0.30 0.17 0.16 1.49 2.94 2.62

K5 III 3 <17 0.25 0.12 0.13 1.88 4.20 3.22

K2 III 3 <17 0.33 0.23 0.18 1.32 3.01 2.84

K2 III 2 - 0.28 0.15 0.17 1.57 3.46 2.94

K1 III - 22 0.65 0.51 0.56 1.31 3.23 2.09

K0 IIIb 1 <17 0.37 0.17 0.23 1.37 3.53 2.64

G8 III 5 <19 0.73 0.57 0.58 0.86 2.58 1.71

K0 III 2 - 0.37 0.20 0.23 1.20 2.85 2.52

G5 III - 19 0.5 0.33 0.36 1.15 3.04 2.35

o UMa G5 III 1 17 0.38 0.21 0.12 1.30 3.27 2.21

G2 IV 3 0 0.58 0.50 0.43 1.06 2.36 2.01

G5 IV 0 20 0.45 0.22 0.29 1.10 3.48 2.28

G2 IV 3 0 0.58 0.50 0.43 1.06 2.36 2.01

G0 IV 0 0.47 0.33 0.25 1.18 2.28 2.07

K2 V 4 <17 0.58 0.44 0.47 1.16 2.96 2.06

K1 V 2 - 0.50 0.32 0.33 1.32 3.11 2.38

G8 V - 2 0.54 0.40 0.43 0.97 2.31 1.80

82 Eri G8 V - - 0.48 0.31 0.33 1.07 3.04 2.37

G0 V 2 6 0.66 0.60 0.46 0.90 1.66 1.58

G0 V 1 6 0.57 0.37 0.37 0.90 2.29 2.02

G0 V 2 6 0.66 0.60 0.46 0.90 1.66 1.58

F9 V 0 3 0.50 0.35 0.33 1.10 1.88 1.83

HR 5317 F7 V - 30 0.58 0.37 0.32 1.17 3.40 2.17

F8 V - 54 0.36 0.21 0.21 1.78 4.27 2.84

 

However, it was found that when the observed spectra of stars are arranged into groups with similar spectral type and luminosity class, within each group there are stars for which the line profiles differ in shape near the line center from the rest of the stars in that group, in the sense that CRI is higher for these stars compared to that of the rest. Some of the examples cited by Linsky et al. (1979) are 111 Tau- Vir, Boo A- Cet, And- Crt, Ori- Com, Gem- Cet, Aur- Crv, Dra- Leo, 9 Peg- Gem; the first one in each pair having a higher CRI than the other. These are perhaps the stars whose absorption profiles show filling in of the 8542 cores by unresolved chromospheric emission. According to Wilson (1976), most of these are the active chromosphere stars with an intensity index higher than 3. Following Linsky et al. (1979), we have also grouped all the observed stars of similar spectral type, luminosity and also metallicity (within [Fe/H] = 0.2) and have found about 13 pairs that are tabulated in Table 3 (click here). The first one is with the larger CRI than the other and is therefore suspected to be chromospherically more active. Listed in the Cols. 5, 6 and 7 of Table 3 (click here), CRI is systematically higher in all the 3 lines in the first star of each pair for all pairs. For 4 of the chromospherically active stars, two examples each of above pairs are cited. The intensity index according to Wilson & Bappu (1957) and Wilson (1976) in the more active stars is 3 or higher except for Ori as seen in Col. 3 of Table 3 (click here). None of these stars are fast rotators. The quantity taken from the Bright Star Catalogue is listed in Col. 4. It is less than for most of the stars. Only Cap, Gem and HR 5317 exceed this value. So rotational broadening cannot be affecting the CRI of the Ca II triplet lines in these stars. Since the latter star (i.e. Cap) in the pair HR 5317- Cap has higher , the relatively higher CRI of HR 5317 cannot be due to rotational broadening. In fact, the Ca II lines of HR 5317 as seen in Fig. 7 (click here)a seem a bit too weak for its high +ve [Fe/H]. We suspect that it is due to chromospheric emission filling in the star. Cet in Fig. 5 (click here)b has a rather low EQW perhaps again because chromospheric emission has partly filled in its absorption profile. As Cols. 8, 9 and 10 of Table 3 (click here) show, the EQWs of all the 3 lines are smaller in the star with the larger CRI than in the star with the smaller CRI. The amount by which they differ ranges from 10 to 30%. However, this may not be the best way of measuring the chromospheric emission that fills in the absorption. Following the scheme of Linsky et al. (1979), we compared the Ca II profiles of each pair (stars with similar spectral type, luminosity & metallicity) by dividing the spectrum of one with the larger CRI by that of the other with the smaller CRI point by point. Figure 9 (click here) shows some of these divided spectra in 8498, 8542 and 8662.

  
Figure 9: Line ratio spectra of , 8542, 8662 of a few sample pairs

 

Ratio spectra FWHM(Å) emission area of

peak(Å) emission

HR 2269 / 3 Cet 0.84 0.78 2.14 2.45 3.25 1.67 2.06 2.54 3.58

HR 2269 / 1.12 1.20 - 2.3 2.3 - 2.58 2.75 -

56 Peg / 1.50 1.15 1.23 2.36 3.25 3.05 3.54 3.74 3.74

56 Peg / 1.06 1.56 1.54 2.6 2.16 2.65 2.77 3.38 4.08

/ 1.17 0.9 1.43 1.84 1.48 1.75 2.15 1.33 2.50

/ 1.10 1 20 1.05 1.23 1.46 1.25 1.35 1.75 1.31

/ 0.70 0.69 0.63 1.52 2.10 1.98 1.06 1.45 1.25

/ 0.84 0.65 - 1.43 2.16 - 1.20 1.4 -

/ 0.83 0.71 0.81 1.64 2.47 2.17 1.36 1.75 1.76

/ 0.8 0.76 0.88 1.90 2.37 1.87 1.52 1.80 1.58

/ o UMa 0.89 0.83 0.6 1.37 1.65 3.05 1.21 1.37 1.83

/ 0.77 0.69 0.67 1.37 2.27 1.88 1.05 1.57 1.25

/ - 0.6 0.7 - 1.54 1.82 - 0.92 1.27

/ 0.86 0.62 0.77 1.18 1.4 1.4 1.01 0.87 1.08

/ 82 Eri 0.82 0.63 0.50 1.75 1.52 1.60 1.44 0.96 0.80

/ 0.92 0.64 1.05 1.27 1.71 1.4 1.17 1.09 1.47

/ 0.86 0.85 1.10 1.32 1.83 1.42 1.14 1.56 1.57

HR 5317 / 1.25 1.00 - 1.87 2.62 - 2.34 2.62 -

 

These ratio profiles actually show net emission in each of the 3 lines, the amount that fills in the absorption profile of the star having larger CRI, i.e., with larger chromospheric activity. They also show the spectral range over which the chromospheric filling in occurs in the active chromosphere stars. Columns 2, 3 and 4 of Table 4 (click here) list the FWHM of the ratio profiles of the 3 lines in the 13 star pairs given in Col. 1. It would be interesting to see how the FWHM relates to the luminosity of the stars in question. The ratio plots give FWHM for the 3 lines in supergiants in the range , in giants from 0.7 Å to 1.1 Å and in dwarfs & subgiants from 0.6 Å to 0.9 Å. The FWHM does increase with luminosity, quite in analogy with the behaviour of the Ca II H and K line emission widths to luminosity. The peak values of the emission profiles of all the 3 lines are displayed in Cols. 5, 6, and 7 of Table 4 (click here). The product of FWHM and the peak value of emission (which is actually a more direct measure of the amount of emission) was calculated for all the dwarf, giant and supergiant pairs and is shown in Cols. 8, 9 and 10. We find that this quantity is more sensitive to luminosity than FWHM alone. In supergiants, it lies in the range 2.1 - 4.0 for the 3 lines, in giants between 1.1 - 1.8 and in dwarfs & subgiants between 0.9 - 1.5. Most of the ratio profiles differ appreciably from unity only in the line core (within of the line core) implying substantial emission in the chromosphere but not in the photosphere. According to Linsky et al. (1979), a better way of detecting photospheric emission apart from the chromospheric emission, is to subtract the profiles of similar stars. Figure 10 (click here) shows the line difference plots for 8498, 8542 and 8662 for a few star pairs. The pairs 56 Peg- Cep, 56 Peg- Gem, HR 2269-3 Cet, HR 2269- Per are much more spread out (out to more than of the line core) than the other pairs, suggesting a strong evidence of emission beyond the line core, implying it originates in the photosphere.

  
Figure 10: Line difference spectra of , 8542, 8662 of the same pairs

Thus, if the CRI of the Ca II triplet lines is not affected by parameters like rotational broadening, macroturbulence, microturbulence etc., then the CRI of especially the 8542 Å and the 8662 Å lines could be viewed as a diagnostic and a measure of chromospheric activity in stars. The nearly one-to-one correlation between FWHM/area of the ratio profiles and the Ca II K line emission intensity index supports the above inference. Linsky et al. (1979) also strongly support the contention that a direct comparison of 8542 CRIs of stars with similar spectral type and luminosity class should at least be a qualitative discriminant of chromospheric activity. The direct spectroscopic diagnostic is the chromospheric radiative loss rate which is known to be significantly higher in active chromosphere stars than in quiet chromosphere stars. Linsky et al. (1979) have estimated the 8542 chromospheric radiative loss rate in stars they have observed and find that stars with larger 8542 chromospheric radiative loss rates are also those with larger Ca II K line intensities. Since the 8542 CRI's of the stars in each group of stars with similar spectral and luminosity type have one-to-one correlation with the K-line index, we contend that 8542 CRI is a good discriminant of chromospheric activity. Cayrel de Strobel (1992) has observed the Ca II 8498, 8542 lines in a small sample of F, G and K dwarfs and slightly evolved subgiants randomly selected from the ``Catalogue of Nearby Stars" of Gliese (1969). Observed profiles of these lines in 6 early K dwarfs having about the same effective temperature show that the shallower the lines are in the spectrum of a star, the more active the chromosphere of this star is.