Observations alpha Tau (K5III) and beta Gru (M2II) made at high resolution are used to confirm line identifications of features blended at low resolution. The high resolution spectra allow selected pairs of lines to be used to find the electron density and the opacity. Methods for determining these factors and the usual emission measure are presented. The electron density and opacity can be used together with the emission measure to place constraints on the structure of the atmosphere. The line formation processes are briefly discussed. Photo-excitation by strong lines appears to be important in these late type atmospheres

Brown, A.

Ferraz, M.

Jordan, C.

English

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THE STRUCTURE OF CHROMOSPHERES AROUND LATE-TYPE
GIANTS AND SUPERGIANTS
A. Brown_ M. Ferraz and C. Jordan
Department of Theoretical Physics
Oxford University, I Keble Road
Oxford, OXI 3NP, England
ABSTRACT
Observations _ Tau (K5111) and _ Gru (M211) made at high resolution have been
used to confirm line identifications of features blended at low resolution.
The high resolution spectra allow selected pairs of lines to be used to find
the electron density, Ne, and the opacity, T. These can be used together
with the emission measure to place constraints on the structure of the atmo-
sphere. The llne formation processes are briefly discussed. Photo-exclta-
tion by strong lines appears to be important in these late-type atmospheres.
INTRODUCTION
Previous observations made with the IUE satellite have shown that giants and
supergiants later than about K2 do not have significant emission in lines
which in the sun would be formed in the chromosphere-corona transition region
(Refs I-5)% Instead, the observed spectra are dominated by lines formed at
Te <2 x l0T K. In particular strong Ol resonance line emission is observed,
which in the sun and _Boo has been shown (Ref. 6,7) to be produced via H Ly
excitation to a higher Ol level.
High-resolutlon spectra have now been obtained of _ Tau (K5111) and _ Gru
(M211) in order to study line fluxes and profiles and also to confirm identi-
fications of lines blended at low resolution. These spectra show that the
strong Ol lines themselves excite lines of SI (uv 9) (Ref. 8), and confirm the
importance of SI as a contributor to other features in late-type stars, as
noted previously for_Ori (M2 Ib) (Ref. 2). A line at 1641.2A is attributed
to Ol pumped by the Ol resonance lines.
In thls paper methods of determining the opacity, and electron density in ad-
dition to the usual emission measure are presented. The mechanisms by which
particular emission lines are excited are briefly discussed.
OBSERVATIONS
The dates, exposure times and resolution of spectra used in the analyses are
given in Table i. All were obtained using the large slot. The observations
were made at the ESA VlLSPA Satellite Tracking Station with the assistance of
the Resident Astronomers.
The low-resolution spectra and samples of the high-resolution spectra have
been previously published (Refs. I, 8) and are not illustrated here. The
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GPHOT image provided by VILSPA is the basis of the further data reduction in
which spectra can be summed and Overexposed pixels removed. An algorithm
provided by Gondhalekar was used to correct for the error in the original
intensity transfer function for the SWP camera. The flux calibration used
is that provided by Bohlin and Snijders (Ref. 9) and by Cassatella and
Selvelli (Ref. I0).
METHODS OF ANALYSIS
The q_antlty usually derived from absolute llne flux is the emission measure,
f NeZ dh. For lines of neutral and singly ionized atoms the contribution
function may be broad, _nd the region of formation not clearly defined.
Also, below Te _ 2 x I0 K, where hydrogen is not fully ionized, the relevant
quantity is f NeN H dh. Thus it is preferable to find f NeNHg(T ) dh, where
g(T) is the temperature dependent part of the contribution function, as a
function of temperature, and build up a self-conslstent model by using sev-
eral lines and independent determinations of Ne and r (_ f NH dh).
The electron density Ne may in late-type giants and superglants be found
from the re_atlve intensity of lines within the 2s22p 2 2p'2s2p3 4p multi-
plet of CII around 2325A. The transition probabilities have been _alculat-
ed (Ref. 11) but cross-sectlons are available only by extrapolating data
for OIV (Ref. 12) and NIII (Ref. 13). 8These approximate values lead tocritical densities in the range 4 x I0 cm-3 to 6 x 107 cm-3. With more
accurate cross-sections the ratio of the Cll multlplets at 1335_ and 23254
may be used to determine Te, providing colllsional excitation is the dominant
mechanism.
The opaclty may be determined from the relative intensities of lines from
a common upper level, where one of the lines is optlcally thin (Ref. 14).
Using a simple probability of escape approach the ratio Qf the fluxes F,
in two such lines is given by
FI/F2 ffi A2blql / Xlb2q2 (I)
where _ is the wavelength, b is the branching ratio, and q is the pro-
bability of escape, which is I for an optically thin line. Then since for
a Gausslan llne profile (Ref. 15),
- 1 - erf_in to) 1/21
q (2)
where 7 o is the opacity at llne centre, q and then T o can be found•
The total opacity 7 =2 7o is given by
7 - 1 2 x 10-14 A(A)f MI/2 N
• 12 f_HN1 _N_ion _.f ds (3)Nion 1/2
where f12 is the absorption oscillator strength, M is the atomic weight,
NE, NI, Nio n and NH are the number densities of the element, lower level,
ion (or atom) and hydrogen, respectively. Ti is the temperature corre-
298
sponding to the Doppler width of the line. Although some approximations
are necessary initially, the quantity fN H ds can be determined, providing
checks on models such as that by Kelch et al. (Ref. 16) for _ Tau. Table
2 gives some examples of pairs of lines which can be useful. The method
may also be applied to numerous lines of Fell above 2000_. The oscillator
strengths required for CI are given in Ref. 14; for Ol values by Garstang
(Ref. 17) have been used; for SI values are tabulated by Wiese et al. (Ref. 18).
It should be noted that the analysis of emission llne fluxes and profiles
in these late-type stars may be complicated by the long ionizatio_ and_
recombination times expected at Te < 2xlO_ K and densities of <i0 _ cm-_;
times of i03-I04s are estlmated for neutral and singly ionized species.
RESULTS
The present paper emphasizes methods since further work on the atomic models
and atomic data is required before quantitative analyses are possible. Some
preliminary results are given below.
In _Tau the Ol 1304/1641 ratio leads to fNH ds N 8 x 1021 cm-2, and the
CI 1657.4/1994 ratio to f NH ds < 6 x 1021 cm-2, assuming all the popul-
ation to be in the ground term. The Sl lines at 1296.17_ and 1302.87A have
approximately the same intensity indicating _ >llin these lines, and giving
NH ds >2 x 1019 cm-2. These values of 6-8xi02 are larger than expected
from the model of the chromosphere of _ Tau based on the Mgll and Call lines
(Ref. 16). Also the absolute flux of the CII 2326_ line leads to an emission
measure of N 1028 cm-5 at 104K, a_ain an order of magnitude larger than
predicted. H_wever, the CII 1335_ llne is observed to be weaker than expected
from CII 2325A and improved excitation cross-sectlons are required before
drawing definite conclusions regarding the temperature of formation. The
Sill lines at 1808_ and 1817_ indicate other problems with the atomic model or
cross-sections. They give an emission measure a further order of magnitude
greater than found from CII 2326_. It has been pointed out than in the sun
the 2D level can be excited predominantly from the metastable 4p level rather
than the 2p level (Ref. 19). Cross-section calculations by Roberts (Ref. 20)
supported this suggestion. However in _ Tau and _ Gru where Ne<109 cm-3,
the 4p level does not acquire sufficient population for this explanation of
the strength of 2p-2D to be appropriate. Either the cross-section for 2p-2D
is still underestimated or other processes such as recombination or ionization
are populating the metastable level. The absence of Silll at 1892_ suggests
however that the population of Silll is low.
Simple estimates show that the Ol emission in _ Tau is too strong to be
accounted for by colllsional excitation or recombination but that the H Ly
excitation process could provide sufflclent emission as in _ Boo (Ref. 7).
Similar results are de_ved for _ Gru. The opacity for Ol is even larger,
gi_ng f NH as _ 3x_ _ c_-z. The Sl lin@s provide limits of
I0_< f NH ds < 4x10 cm- . The CII 2325_/1335 ratio again suggests a low
299
Te or an overestimate of the 2325_ emission measure. The Sill 1808, 1817
lines require an emission measure an order of magnitude larger than that
for CI12325_. For _ Gru _he C_I 2pU4p line ratios can be used to determine
Ne, giving a value of ~I0 cm-_.
The relative intensities of the Fell lines above 2000_ allow some general
conclusions to be drawn. Although the Fel (mult. 44) lines pumped by the
Mgll emission (Refs. 21, 22) are observed in both_Tau and _ Gru these lines
are apparently not broad enough to pump Fell uv 32 or uv 373. Mult. uv 32
is present and its strength can be attributed to pumping by uv 1 which has all
opacity of _ 103. There is similar interlocking between the other multiplet_
observed. In _ Gru uv 399, which terminates on the upper level of uv 63 is
particularly strong, and the population of the upper and lower levels of uv
399 are comparable, indicating a selective population mechanism. It has beeu
pointed out previously (Ref. I) that H Ly _can excite levels in the 5p con-
figuration of Fell, the classifications of which have been extended by
Johansson (Ref. 23). These 5p levels can then decay to the 5s levels giving
rise to uv 399 and similar multlplets. (Penston has privately pointed out
the strength of uv 399 in RR Tel. and suggested recombination as a cause).
Johansson found that charge exchange with Nell is important for Fel-_Fell
in laboratory sources. Since the Fell multiplets are important not only for
stellar spectra but also for quasars and Seyfert galaxies it is of interest
to investigate these possibililes further.
Finally, the possible ro_e of H Ly _ in photo-excltlng (or effectively ion-
izing) high levels of Sl is pointed out. The relative strengths of the Sl
1425 and 1475 multlplets do not appear to be consistent with collislonal
excitation, using f-values and solar relative intensities as a guide to the
cross-sectlons. The strong lines of SI observed ie. mult. uv 3, uv 2, uv I_
and uv 9 originate from the lowest terms formed by recombination from the
three possible SII parents. The strength of the SI emission suggests a
recombination or cascade spectrum following photo-lonlzatlon (or excitation)
by radiation around H Ly _.
REFERENCES
I. Brown, A., Jordan, C., and Wilson, R. 1979 Proc. Symp. 'The First
Year of IUE' (Ed. Willis) 232.
2. Carpenter, K.G. and Wing, R.F. 1979 Bull. Am. Astr. Soc. i_i, 419.
3. Dupree, A.K., Black, J.H., Davis, R.J., Hartmann, L., Raymond, J.C.
1979 Proc. Symp. 'The First Year of IUE' (Ed. Willis) 217.
4. Dupree, A.K. 1979 'Highlights of Astronomy, 1979' In Press.
5. Linsky, J.L. and Haisch, B.M. 1979 Astrophys. J. Letts. 229, L27.
300
6. Shine, R.A., Lites, B.W., Chipman, E.G., Roussel-Dupre_ D.,
Brunet, E.C., Rottman, G.R., Orrall, F.Q., Athay, R.G.
and White, O.K. 1976 Bull. AAS, 8, 331.
7. Halsch, B.M., Linsky, J.L., Welnsteln, A. and Shine, R.A. 1977
Astrophys. J. 214, 785.
8. Brown, A. and Jordan, C. 1980 Mon. Not. R. astr. Soc. In Press.
9. Bohlln, R.C. and Snljders, M.A.H. 1978 SRC IUE Newsletter No. 2.
I
I0. Cassatella, A. and Selvelll, P.L. 1979 SRC IUE Newletter No. 4.
II. Dankwort, W. and Trefftz, E. 1978 Astron. and Astrophys. 65, 93.
12. Flower, D.R. and Nussbaumer, H. 1975 Astron. and Astrophys. 45, 145.
13. Nussbaumer, H. and Storey, P.J. 1979 Astron. and Astrophys. 71, L5.
14. Jordan, C. 1967 Sol. Phys. _, 441.
15. Zanstra, H. 1949 Bull. Astron. Inst. Neth. II, I.
16. Kelch, W.L., Linsky, J.L., Basrl, G.S., Chlu_ H-Y., Chang, S-H.,
Maran, S.P., Furenlld, I. 1978 Astrophys. J. 220, 962.
17. Garstang, R.H. 1961Proc. Camb. Phil. Soc. 57, 115.
18. Wiese, W.L., Smith, M.W. and Miles, B.M. 1969 'Atomic Transition
Probabilities', Vol. II, NSROS-NBS 22, (US. Govt.
Printing Office, Washington, D.C.).
19. Jordan, C. 1969 Astrophys. J. 156. 49.
20. Roberts, D.E. 1970 J. Phys. B. !, 676.
21. Van der Hucht, K.A., Stencel, R.E., Halsch, B.M. and Kondo, Y.
1979 Astron. and Astrophys. Suppl. 36, 377.
22. Gahm, G.F. 1974 Astron. and Astrophys. Supp118, 259.
23. Johansson, S. 1979 Physlca Scrlpta 18, 217.
301
Table 1
Spectra Obtained
Star Type Date Exposure Time (min) Resolution Camera
Tau KSIII 78/9/29 20 LO SWP
78/10/1 40 LO SWP
79/1/25 90 LO SWP
79/1/27 150 LO SWP
79/9/29 390 HI SWP
78/1/10 i0 LO LWR
79/9/29 I0 HI LWR
Gru M211 79/10/1 175 LO SWP
79/10/7 370 HI SWP
79/10/1 20 HI LWR
Table2
Opacity Sensitive Line Sets
Atom Transition AJ %(A)
CI ,2p2 3p_2p3s 3po 2-1 1658.1
l-1 1657.4
0-1 1657.9
2p2 ID_2p3s 3po 2-1 1993.6
2p2 3p_2p3s ipo 2-I 1614.5
I-I 1613.8
0-i 1613.4
2p2 1D_2p3s 1po 2-1 1930.9
Ol 2p4 3p-2p33s 38° 2-I 1302.17
I-I 1304.86
0-i 1306.03
2p4 1D-2p33s 3S° 2-1 1641.30
Sl 3p4 3p_3p34s,, 3po 2-1 1296.17
i-i 1302.87
3p4 3p-3p34s 3S° 2-1 1807.34
I-I 1820.36
0-i 1826.26
302


The structure of chromospheres around late-type giants and supergiants

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