Einstein observations of hot stars

67—73, 1983 Ado. Space Res. Britain. Vol.2, No.9, Printed in Great All pp. rights reserved.

0273—11771831090061—07$03.50/O Copyright

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EINSTEIN OBSERVATIONS OF HOT STARS Joseph P. Cassinelli Washburn Observatory, University of Wisconsin-Madison, 475, N. Charter Street, Madison, WI 53706, U.S.A.

ABSTRACT Surveys with instruments on the Einstein Observatory have shown that essentially all 0 and B main sequence stars are X—ray sources as are many, if not all, 03 supergiants and Wolf-Rayet stars. The X-ray luminosities are sufficient to explain broad lines from the superionization stages seen in the UV spectra of the stars. High energy resolution spectra from the Solid State Spectrometer are shown to place severe constraints on various models for the location of the X-ray sources in the outer atmospheres of the stars. Coronal and embedded shock models for the X—ray emission are discussed and each is found to have some problems in explaining the X-ray emission of OB stars. X-ray line emission of Si XIII and S XV in ~ On is discussed and Interpreted as arising from magnetically confined loops. INTRODUCTION X-rays emission from hot stars was one of the first discoveries of the Einstein Observatory (Harnden et al. [1],Seward et al. [2]). The detection of X—rays resolved the controversy discussed by Cassinelli, Castor and Lamers [3]concerning the ‘supenionization stages such as 0 VI A104O and N V Al250 that are seen as strong broad lines in the UV spectra of Of stars and OB supergiants. These lines are from ionization stages that are higher than those predicted from models of line driven winds assumed to be in radiative equilibrium. The observed X-ray flux from the stars is sufficient that the high ion stages can be produced by K shell ionization followed by the Auger effect. The net effect of the X-rays is the removal of two ~lectronsfrom some of the ions in the winds such that, for example, trace amounts of O’~are produced from o+3, which is the dominant star of ionization in the winds of st~ 5sof spectral cla~sBO.SIa and earlier. Similarly N ‘~ results from Auger ionization of N , and C+) from C+1 in later type supergiants. Most of the X-ray observations of hot stars were carried out with the Imaging Proportional Counter (IPC) on the Einstein Observatory. The results have shown that the 0.2 to 4 key 7 forluminosity 0 stars with scatter only abqut X-ray luminosity, Lx, is proportional to the total for aearly type of stars. Seward aandfactor of 2. [4]have For Wo1f-Rayet starsLx/L they 2x10 have found a wider range of Lx/L, from 6xl0~° to Chlebowski shown that 2x10 . The characteristic ~e~peratu~e 0for early type stars can be roughly estimated from the IPC spectra to be T l0°~~ to 10 ~° Figure 1, from a review of Cassinelli and MacGregor [5]shows the X—ray luminosities of various types of stars in the H-R diagram. The range of stars which obey the L~/L relation is indicated. The dashed lines show the regigns ~ th~.,H-R dia~amin which early type stars show the superionization stages o~0+ , N , C and S1 ~. These ions can be used to show the L~/Lratio decreases below 10 (to a value of —io°~~ at B9) for supergiants later than 82 (Odegard and Cassinefli) [6]. As a contrast note that for the late type main sequence stars the X—ray luminosity is correlated with rotation rates and presumably with magnetic field strengths (Pallavicini et al.) [7]. It will be interesting to see if the X—ray emission of hot stars is in any way affected by magnetic fields. This possibility will be addressed later. I now want to focus on the problem of current major interest. Where is the source of X-rays located in the outer atmospheres of the hot stars? Two possibilities have been considered so far. The X—rays could arise from a corona at the base of the wind or they could come from hot wisps embedded within the wind. A) Coronal Models A model for the X-ray emission that was studied prior to the launch of Einstein is the Slab Corona plus Cool Wind Model (Hearn [8],Cassinelli and Olson [9],Cassinelli [10]). In this picture, a thin corona is assumed to be present at the base of the wind with a temperature of several million degrees. The temperature is assumed to drop very sharply with 67

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Fig. 1 X-ray emission In the HR diagram. The solid dots indicate stars detected as X—ray sources by the Einstein satellite. For several stars the logarithm of the X—ray luminosity is shown in the ovals. Upper limits are also shown for a few supergiants and late-type giants. The extent of the presence of broad lines of the “superlonization” Ions in the UV spectra at 0 VI, N V. C IV, and Si IV are Indicated by the dashed lines for the early—type stars. For early—type stars L is proportional to the stellar luminosity, as indicated, while fo~ late-type stars the X—ray luminosity is proportional to rotation speed squared as might be expected from coronal emission that depends on a dynamo mechanism. This figure is from the review by Cassinelli and

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Einsteifl Observations of Hot Stars

69

radius at the top of the corona because radiative cooling is very effective at the large densities in the winds of OB stars. The temperature above the slab corona is <30,000 K as expected in winds in radiative equilibrium. The model predicts a strong attenuation of X-rays at E < 1 key by the cool wind material. Even with the relatively low resolution spectra of the IPC, Long and White [11]were able to show that there is a larger X-ray flux at E < 1 key than predicted by the slab coronal model. This is very evident in the spectra of OB supergiants observed with the Solid State Spectrometer (SSS) from Cassinelli and Swank [12],shown in Figure 2. The slab model predicts a sharp drop at 0.6 key because of the K shell absorption edge of Oxygen. The SSS spectrum clearly does not show that edge. I

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Energy (keV) Fig. 2 Shows the Solid State Spectrometer (SSS) spectrum of c On BOla. The heavy line shows the SSS spectrum predicted by a model in which the X-rays originate at the base of the wind In a thin slab corona. The absorption edge clearly seen at —0.6 key is caused by K shell ionization of oxygen in the cool wind. The lack of this absorption in the observed spectrum suggests that much of the soft X—ray emission comes from sources well above most of the wind opacity. Several modifications of the coronal model have recently been presented to account for this problem. Stewart and Fabian [13]suggest that the mass loss rates assumed for the stars is too large and that the winds are clumpy. The net effect of their modifications is that the optical depth of the cool wind to X—rays is reduced and it is possible to get good fits to the IPC spectra by adjusting the wind optical depth, T. Waldron [14]has calculated models for the recombination region of the wind and has found other reasons why the slab model may have overestimated the attenuation. The wind just above the corona can be heated somewhat by the absorption of coronal X-rays and the velocity at the base of the wind should be higher tha~assumed by Cassinelli and Olson [9]. By using the emission measure of the corona (N Vol) as an adjustable parameter Waidron [14]was able to get good fits to the IPC sp~ctraof hot stars. There are problems with these modified coronal models. It is possible to fit the IPC spectra, by a very fine adjustment of the optical depth of the wind [13]or the emission measure of the corona [14]. However, all the 0 stars observed have about the ~ame X-ray spectral distribution and certainly not all of the stars should have, r, and Ne’ Vol throttled so perfectly to account for this. More important, however, is that even the modified coronal models that fit the low resolution IPC spectra don’t predict the negligible 0.6 keV absorption edge that is observed in the higher resolution SSS spectra.

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B) Embedded Shock Models Nelson and Hearn [15],MacGregor, Hartmann and Raymond [16],Carlberg [17],and Kahn [18] have all found that line driven flows are unstable. Lucy and White [19]conjectured that the line driven Instabilities grow to produce a two phase wind medium consisting of blobs embedded in a diffuse gas. Bow shocks form at the interface between the radiatively driven blobs and the ambient gas. As these shocks are located somewhat above the base of the wind there should be less attenuation by the overlying wind of the X-rays produced in 6 the K) bow to shocks. a The temperature of theFurthermore shocks is somewhat (—lxlO provide goodcharacteristic fit to the observed spectra. the shockstooarelowstill so close to the base of the wind that there Is too much absorption at 0.6 keV to provide a good fit to the SSS spectra. Lucy [20) has proposed a second embedded shock model. in this model the velocity of the unstable flow is assumed to develop a saw-tooth pattern, illustrated in Figure 3, associated with the periodic passage of shocks. The flow is, in fact, accelerated by the radiatively driven shocks. A major difference from the blob model is that shocks are assumed to survive until shadowed by following shocks. This shadowing and decay of the shocks does not occur until well out Into the flow and hence very little attenuation of the X-rays is expected. ~—

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FIg. 3 The phenomenological description of the hypersonic motions adopted by Lucy (20]. The functions V(r) and U(r) denote the shock’s velocity in the star’s frame and theIr propagation speed, respectively, at radial distance r. The quantity v Is the basic parameter of the model and Is defined such that va Is the velocity interval within which matter In a shock’s wale 1.s irradiated by unattenuated photospheric continuum. (adopted from Lucy [20]) The basic parameter In the model is the quantity, v, which describes the Increment in velocity In the sawtooth pattern from one shock to the next, Av = va 5, In terms of the sound speed, a5, (20 km/sec) in the cool Intershock flow. Lucy [20]argues that v should have a value of —0.4. He finds, however, that this value of v leads to shocks that are far too weak to explain the observed X—ray spectra of hot stars. To overcome this problem, Lucy suggests that there are a few much stronger shocks which are characterized by a second much larger value of v. Lucy has kindly carried out calculations with his model to find the parameters required to provide a fit to our SSS spectrum of £ On. The results are shown in Figure 4. The figure shows the spectrum predicted from a model with v having the expected value of 0.4. The flux Is seen to be many orders of magnitude too low and the distribution is far too soft. However, a very good fit to the spectral distribution is pro1x =the(l-c vided by a model with v • 3.4. Combining results for the two values of v(v1 = 0.4; V2 = 3.4) as suggested In Lucy’s paper 2)11 + £212, we find £2 1/25 and a relative frequency of the strong shocks to be £2 “I = 1/200.. This means that one shock out of 200 must be a strong shock. (1-c,) v~ Thus, relatively infrequent strong shocks can account for the observed SSS spe~tru0of c On. Serious problems remain however. Firstly, the required strong shock value of v 3.4 is far beyond the maximum of v 1.0 predicted in the Lucy’s periodic shock model. So the production of strong shocks is not yet fully explained. Secondly, the relative frequency

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Energy (key) Fig. 4 Shows the SSS spectrum of £ On and several fits using Lucy’s[2O] periodic shock model. The shocks are parameterized by the quantity, v, described in Fig. 3. The preferred value of v Is 0.4, but as seen here this leads to an X-ray flux far below and far softer than the observed distribution. The spectrum for the case with v • 3.4 gives rise to a spectrum that is sufficiently hard and to more X-rays than are needed to fit the observation. If a fraction of only 1/25 of the total X—rays are assumed to come from these strong shocks a good fit is achieved. (reprinted from Cassinelli and Swank [12]) of the strong shocks is so small that the stars should perhaps be strongly variable in the X-ray region. Figure 5 shows results from the recent search for variability by Cassinelli, Hartmann, Sanders, Dupree and Myers [21]. The X-ray flux of c On was found not to vary over a span of —7 hours. Furthermore the X-ray flux level was the same as that of an observation one year earlier.

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TIME (iOns) Fig. 5 Total IPC Count rate for c On from a simultaneaous UV and X-ray observation made in 1980. The solid arrow indicates the mean count rate and the associated statistical uncertainty in the mean. The dashed arrow shows the count rate from c On from an observation one year earlier. On the basis of this date and the lack of any changes in the UV line profiles of supenionization stages it may be concluded that the X-ray emission is very steady in time. (Reprinted from Cassinelli et al. [21]) X-RAY LINES OBSERVED WITH SSS AND A POSSIBLE INDICATION OF MAGNETICALLY CONFINED HOT LOOPS We have seen that the X-ray spectra of early type stars indicates that at least the X—rays near 0.6 key appear to be produced in sources or ‘wisps’ embedded in the stellar wind. We might wonder if there is any indication for X-rays from a base coronal region as proposed in the original model. Figure 6 shows the SSS spectrum of ~ On 09.5Ia. Note the emission features at 1.8 and 2.2 key. • 0~ These are the locations of the strong X-ray lines ~ CORI of the helium-like ions Si XIII and S XV. The lines are very prominent in the SSS spectra of ~—. Supernova Remnants, [22],Capella [23],and in several other early type stars such as 7 Sco 09V .~ (Seward [24]). Cassinelli and Swank [12]have found that the fit to the SSS spectrum is ~‘ —2 significantly improved if a two source component r model is assumed. The results indicate that the region responsible for the Si XIII an9 Si XV lines 9 are from a hot source with l.5x10 cm-3. K and an 2 Vol T= = 2xl054 — (A) emission measure of Ne The temperature of 15 million degrees is I • particularly interesting because it is much 0.6 1.0 2.0 3.0 4.0 largeç than the maximum coronal temperature of —7x10~K derived from Parker Coronal theory. ENERGY (key) That is, a gas at 15 million cannot be confined by the gravitational field of the star. In Fig. 6 analogy with similar situations for hot gas in solar corona and in the outer atmospheres of other cool stars, we have postulated that the hot gas is magnetically confined in hot loops above the photosphere. Using the Rosner, Tucker and Valana ~ theory for loop structures we find an expression for the loop field strength. B 4 r I gauss where F is the fraction of the star covered by loop foot prints. Unfortunately F is so uncertain that this equation doesn’t provide much of a constraint on B.

Einstein Observations of Hot Stars

Guessing F 0.01 we get B = 40 gauss, the loop half lengths are 1/8 R~,and the total number of loops ~ ~25o. Thus the line emission could be explained with a large number of small magnetically confined regions. This type of model is of particular interest because it suggests that hot stars have a multiphase solar-like coronal structure that appears to be prevalent across much of the cooler regions of the HR diagram. In summary: the Auger ionization mechanism appears to be the explanation of the high velocity superionization lines seen in the UV spectra of hot stars. The locations of the sources of the X-rays is still uncertain. There may well be a need for sources of X-rays both at the base of the flow and in hot wisps from shocks in the wind. There is a possibility that solar-like magnetic structures exist in base coronae. 1. 2. 3. 4. 5. 6. 7. 8. 9. 10. 11. 12. 13. 14. 15. 16. 17. 18. 19. 20. 21. 22. 23. 24. 25.

RE FERENCES F.R. Harnden, Jr., G. Branduardi, M. Elvis, P. Gorenstein, J. Gnindlay, J.P. Pye, R. Rosner. K. Topka, and G.S. Vaiana, Ap.J.(Letters), 234, L5l (1979). F.D. Seward, W.R.Forman, R. Giacconi, R.B. Gniffiths, F.R. Harnden, Jr., C. Jones, and J.P. Pye, Ap.J. (Letters), 234, L55 (1979). J.P. Cassinelli, J.I. Castor, and H.J.G.L.M. Lamers, Publ.Astr.Soc.Pac., 90, 496 (1978). F.D. Seward and T. Chlebowski, Ap.J., 256, 530 (1982). J.P. Cassinelli and K.B. MacGregor in Physics of the Sun (ed. P. Sturrock) in press (1982). N. Odegard and J.P. Cassinelli, Ap.J., 256, 568 (1982). R. Pallavicini, L. Golub, R. Rosner, G.S. Vaiana, T. Ayres, and J.L. Linsky, Ap.J., 248, 279 (1981). AT~. Hearn, Astron.Astnophys., 40, 277 (1975). J.P. Cassinelli and G.L. Olson, Ap.J., 229, 304 (1979). J.P. Cassinelli in IAU Symposium No. 83, Mass Loss and Evolution of 0-Type Stars, ed. P.S. Conti and C. deLoore; D. Reidel Publ. Co., Dordrecht p. 201, 1979. K.S. Long and R.L. White, Ap.J.(Letters), 239, L65 (1980). J.P. Cassinelli and J. Swank submitted to Ap.J. G.C. Stewart and A.C. Fabian, M.N.R.A.S., 197, 713 (1981). W.L. Waidron submitted to Ap.J., Wisconsin Astrophysics No. 148 (1982). G. Nelson and A. G. Hearn, Astron.Astrophys., 65, 223, (1978). K.B. MacGregor, L. Hartmann, and J.C. Raymond, Ap.J., 231, 514, (1979). R.G. Carlberg, Ap.J., 241, 1131 (1980). F.D. Kahn, M.N.R.A.S.,T~6, 641 (1981). LB. Lucy and R.L. White, ~p.J., 241, 300 (1980). L.B. Lucy , ~, 255, 286 (l982T~ J.p. Cassinelli, 1. Hartmann, W.T. Sanders, A.K. Dupree and R.V. Myers, submitted to Ap.J. 1982. R.H. Becker, A.B. Szymkowiak, E.A. Boldt, S.S. Holt, P.J. Serlemitsos, Ap.J.(Lettens), 240, L33 (1980). ~ Molt, N.E. White, R.H. Becker, E.A. Boldt, R.F. Mushotzky, P.J. Serlemitsos, and B.W. Smith, Ap.J. (Letters), 234, L65 (1979). F.D. Seward private communication. R. Rosner, W.H. Tucker, 6.5. and Vaiana, Ap.J., 220, 643 (1978).