0.0 11.0

10.5 0

-0.2 10.0

9.5 o -0.4 o '"ci I '"0 I I 9.0

-0.6 8.5

o 0 8.0 o o -0.8 80 7.5 90 o

-1.0 7.0 0.0 0.1 0.2 0.3 0.4 0.5 2.0 2.5 3.0 3.5 4.0 6(B-H)-O.5 (B-H)-O.5 Figure 2: Displaeement in (B-H) and H due to the inelination effeets at Figure 3: Simulated (B-H)-0.5-H-O.5 relation (or a sampie of model log NDo = -0.5, and extrapolation of 8-0.5 from large aperture galaxies, when no eorreetions for inelinations are made. Straight line photometrie data by means of mean growth eurves. shows the C-M relation (rom whieh the data has been generated. ing a sampie of primarily face-on galax­ photometer with the "Quantacon" in the Thus there is hardly enough to comment ies for observations in Band H, at aper­ B band, and after a change of setup to on yet. Hopefully I shall be in a better tures near log NOo = - 0.5, would seem infrared in the middle of the observing position after my next run at La Silla. to be a good opportunity to study the period, with an infrared photometer us­ C-M relation in more detail. When at the ing an InSb detector for the H band References same time one chooses to use galaxies photometry. Unfortunately, very poor Aaronson, M., Bothun, G.O., Mould, J. R., with known redshifts, the full step of weather conditions were encountered Huchra, J., and Schommer, R.A., 1986. making a contribution to our further during the major part of the first observ­ Ap. J. 302, 536. understanding of "Iarge scale deviations ing period. Useful observations under Fixen, O.J., Cheng, E.S., and Wilkinson, from the Hubble flow" is within reach. photometrie conditions could be per­ 0.1.,1983, Phys. Rev. Lett. 50, 620. Thirteen nights in two runs (Sep­ formed in B for less than 1 night and in H Giraud, E., 1987. The Messenger No. 49, 20. Lynden-Bell, 0., Faber, S. M., Burstein, 0., tember/October 1987 and May/June for 1 night so far. Since each galaxy has Oavies, R. L., Oressler, A., Terlevich, R. T., 1988) at the ESO 1-m telescope have to be observed through a couple of and Wegner, G., 1988. Ap. J. 326, 19. until now been allocated to a project, apertures in each band, only observa­ Rubin, V. R., Thonnard, N., Ford Jr., W. C., where this is the main goal. The data tions of an uninterestingly small number and Roberts, M.S., 1976. Ast. J. 81, 719. were acquired using a single channel of galaxies have been finished so far. Wyse, R.F.G., 1982, M.N.R.A.S. 199, 1P.

Searching tor Double Degenerates A. BRAGAGLlA, L. GREGGIO, A. RENZINI, Dipart. di Astronomia, Universita di Bologna, Italy S. O'ODORlCO, ESO

that the majority of binary systems must has evolved independently of the other, The Elusive Progenitors of Type I at some stage become a double WO as if it were a single star. An example of Supernovae binary, and that a good fraction of the this kind may be provided by the Sirius The majority of stars are members of stars listed in our catalogues as WOs system: in less than one billion years binary systems; this is a frequently re­ may indeed be double. also Sirius A will be a WO, and Sirius ferred aphorism. The majority of stars Very few WOs, however, are known A+B will become a wide double degen­ also terminate their evolution as white for being double, and indeed Greenstein erate (00) system. But in close binaries dwarfs (WOs), as only few stars are born (1986) lists just six such objects. These the evolution of each component is se­ massive enough to experience a super­ are very wide pairs, visual double de­ verely affected by the presence of the nova explosion. We can easily conclude generates, in which each component companion: The system can experience

35 common envelope phases during which 600 r---.r---<::--,--_,_-_._---:::.,._-n-__,~~-_,_-_._-.,.__,_..___,-__,_-_n10 first the envelope of the primary, and subsequently that of the secondary, are "- expelled from the system thus leaving "- bare, WO cores. The energy required for "- "- "- such an ejection is supplied (via friction) "- "- "- "- "- "- by the orbital motion of the binary: The "- orbit then shrinks during the common 2 3 envelope stages, and at the formation of 2.7 6 \J the second WO the separation can be I'l much smaller than the original separa­ " 2.87 ;:0 0 tion, when both stars were on the main 0 sequence. This is just contrary to the evolution of wide binaries, where mass 42.- loss from the system leads to even larger final separations. 3.5 The relation between initial and final \ \ \ separation is theoretically very uncer­ \ 2 tain, an uncertainty wh ich follows 5.0 primarily from the hydrodynamical com­ 7.0 plexities of the common envelope stages wh ich prevent a secure estimate of the efficiency (a) with which the orbi­ 6 9 12 15 TIME (Gyr) tal energy is converted into energy of the ejected envelope (I ben & Tutukov Figure 1: The evolution of the period and orbital velocity ofa double degenerate system of two 0.7 M white dwarfs, as angular-momentum losses by gravitational wave radiation progres­ 1984). The smaller this efficiency, the Q sively shrink the orbit. Each curve is label/ed with the initial semi-major axis in RQ units. Note closer the two WOs as the second WO is how the merging time is strongly sensitive to the initial separation of the 00 system. Note also formed, i. e. the smaller the initial sep­ that objects with initial orbital velocity in excess of - 300 km s -1 merge in less than one Hubble aration of the 00 phase in the evolution time (assumed to be - 15 x 109 yr). of the binary system. The duration of this phase, unlike the previous nuclear phases of stellar evolu­ from their strong X-ray emission In this frame, we have started an ob­ tion, is controlled by a new clock: gravi­ (Fabbiano 1986, Canizares et al. 1987, servational programme aimed at reduc­ tational wave radiation. The fate of 00 Loewenstein and Mattews 1987). Clear­ ing the uncertainties still enshrouding systems is then to spiral in, until the less Iy, the understanding of important as­ both the final stages of binary evolution massive WO fills its Roche lobe and the pects of galactic evolution requires a and the nature of SN I progenitors. The two components merge. If the initial kowledge of the rate of type I SNe over questions we aimed to answer are basi­ semimajor axis of the 00 is smaller than cosmological times, which in turn de­ cally two: (1) How frequent are 00 sys­ a critical value, then merging will occur pends on the detailed distributions of tems among known white dwarfs? and within a time less than the Hubble time, binary parameters, like masses and (2) what is the fraction of such systems and merging itself will become an as­ separations. Still, the mentioned that will likely merge in less than one trophysically interesting event. Indeed, if theoretical uncertainty in the efficiency Hubble time? Figure 1 shows the com­ the 00 consists of two carbon and oxy­ parameter a prevents a deterministic puted evolution of the period and orbital gen WOs and its combined mass ex­ estimate of SN I rates as a function of velocity of 00 components, as gravita­ ceeds the Chandrasekhar limit population age, and at this point only tional wave radiation progressively (-1.4 MG) then either a thermonuclear dedicated observations of SN I candi­ shrinks the orbit. One sees that the initial explosion (carbon deflagration) destroys date progenitors could help passing orbital velocity of the 00 must exceed the star completely (Iben and Tutokov beyond this bottleneck. - 300 km S-1, if merging is to take place 1984), or the collapse of the merged object leads to the formation of a neut­ ron star (Nomoto and Iben 1985), de­ pending on the uncertain timescale of the actual merging process. 80th 4000 possibilities are astrophysycally very in­

~ teresting, as either a type I supernova ~ event or a millisecond pulsar would then '" ">. 3000 L.. respectively be produced (Iben and ~ Tutukov 1984, Saio and Nomoto 1985). ~ The former case is perhaps more excit­ .::::.. 2000 ing, due to the strong impact that type I x 0;::" SNe have on the evolution of galaxies. Type I SNe, in fact, are thought to gen­ .,' 1000 ,N erate most of iron-peak elements pro­ duced in galactic nucleosynthesis (e. g. y/ Greggio and Renzini 1983, Matteucci 0 and Greggio 1986) as weil as most of 3500 4000 4500 5000 5500 the heating required to maintain the in­ Wavelength terstellar medium of elliptical galaxies to Figure 2: Two spectra of WO-17 taken 25 hours apart. The flux scale is arbitrary. Note the shirt the observed high temperature inferred of the Balmer lines, which corresponds to a radial velocity change of -220 km S-I. 36 within one Hubble time. In this example, twice and possibly three times, always different times. To get such differences, this corresponds to an initial period of using EFOSC with the B 150 grism we use a programme whose original - 12 hours. Therefore DD systems in (spectral range 3600-5590 A) giving a version had been written by Pier (1983) wh ich we are interested should exhibit resolution of - 6.5 A. and which has been adapted to work on radial velocity variations of a few The reduction was carried out with the WDs, where generally lines are very hundred km s " that can easily be de­ IHAP system in Garching. The spectra broad so that the wings are very impor­ tected using modern instrumentation. were wavelength calibrated with a tant. The programme uses one spec­ Correspondingly, our observational helium lamp generally taken at the same trum as template; after normalizing, strategy has been to start a spectros­ telescope position of the target object. flattening with continuum subtraction, copic survey of catalogue WDs, with the This data base allows to test the limiting and rebinning, the programme com­ aim of finding high orbital velocity DD accuracy of EFOSC for radial velocity putes the cross-correlation of this tem­ pairs. We stress that theory cannot real­ determinations. By comparing the diffe­ plate with a second spectrum of the Iy predict what fraction of WDs are actu­ rent calibration exposures taken same star. The position of the cross­ ally DD systems: This is in fact the main through the night at different zenithai correlation peak gives immediately the motivation for undertaking such a distances (from 1 to 2 air masses) we velocity difference between the two search, that ultimately we would like to note that the lines fall in the same posi­ spectra. extend to -100 WDs, so as to ensure tion on the CCD to within - 0.15 pixels, some statistical significance of the which would correspond to a systematic Two Good Candidates and Two sampie. shift of - 35 km S-I at A = 5000 A. The WDs with Red Dwarf Companions accuracy of the dispersion curves fitting the calibration lines of an individual ex­ We have determined V variations for Hunting White Dwarf Binaries r with EFOSC posure is somewhat better, giving an the 20 programme stars for which we RMS of -0.3 A (or -20 km S-1 at A = have at least two spectra. Among these Our target objects were chosen from 5000 A). Our velocity determinations are WDs two turned out to have red dwarf the Catalogue of Spectroscopically based on several absorption lines, and (RD) companions, and these two inter­ Identified White Owarfs of McCook and hence we estimate that the error is actu­ esting objects are discussed below. Of Sion (1984). The selection criteria of our ally smaller than 20 km s-'. For the ob­ the 18 residual dwarfs 9 were observed choice have been very simple: We jects observed in the first run, we used a in the first run, and 9 in the second. All wanted the WDs to be relatively bright in single calibration exposure taken at the the first run objects show Vr values order to minimize the observing time beginning of the night. This leads to a between - 50 and -150 km S-I, just required to get a high S/N spectrum, larger uncertainty in the velocity deter­ consistent with all being the result of our and we further restricted to WDs of the minations for the first run objects, that poorer accuracy during that run. The DA variety in order to achieve the high­ we then estimate to be at least of the situation is far more exciting for the sec­ est possible accuracy in radial velocity order of - 60 km s-'. ond run objects: here seven dwarfs ex­ determinations, thanks to their promi­ Given the main goal of the project, we hibit 6. Vr consistent with 1a errors nent Balmer lines. Basically, our obser­ are primarily interested in the actual val­ -20 km s-', one WD has Vr = vational strategy has been to take sever­ ue of the vr differences between two or 70 km s-', and for the last dwarf (WD­ al spectra of the same WD, a few hours more spectra of the same star, taken at 17, in our list of targets) the cross corre- apart (in practice, intervals between 2 and 25 hours), and then look for any radial velocity change. This should be .5 most efficient in the case of DDs with components of very different luminosity (very different cooling time), while old DDs should rather have similar luminosi­ ty, and then exhibit variations in line , profile, rather than in line position. 2 , 1': , EFOSC at the 3.6-m telescope was cho­ 0 , l,~'"rl'\ .....,v- 1\1 N'} . ~, .\1 I I sen as this provides the optimal observ­ 4 I \ \ '~' ing configuration: thanks to its high effi­ ~ ~ , , ciency, high S/N multiple spectra of sev­ ,\~ I eral objects can indeed be obtained in J; 'i one night. Moreover, each spectrum is ~ obtained with a relatively short exposure ~ « - 15 min), wh ich then ensures ade­ :; - 5 quate time resolution also in the case of '- short-period objects. Our observational strategy is very different from others that have been unsuccessful in singling out WD0034-211 DD candidates that may be SN I pro­ genitors. In particular, we are not bound only to short-period systems as in the -1 JOOO 5000 6000 7000 case of the survey of Robinson and Wavelength Shafter (1987). Figure 3: Continuum subtracted and f1attened spectrum of WO 0034-211 .. this is a composition Till now we had two observing runs at of one blue and one red spectrum, taken respectively with the B 150 and 0 150 grisms. the ESO 3.6-m telescope, respectively Shortward of -5000 A the WO component dominates the light, while the red companion in September 1985 (2 nights) and in dominates at longer wavelengths. The faint emission core in the Balmer lines suggests the January 1988 (3 nights). A total of presence ofan accretion disko In the red side of the spectrum the most prominent features are 20 WDs have been observed at least the narrow Na absorption, several TiO bands, and the strong Ha emission. 37 5000 pated: First, we can check the mass­ luminosity relation for RDs, which theoretically is extremely uncertain (cf. 4000 f- rvJt1Jw"1~ - Liebert & Probst 1987, Renzini & Fusi

~ Pecci 1988) while its knowledge is fun­ c ;::. "'V"~ damental for the derivation of the low >- 3000 f- ~j - ~ 4 , end of the initial mass function, and then -e for assessing the baryonic contribution ~ to dark matter. Second, the determina­ x 2000 l- - ;:: tion of the masses and separation of 0;: :(f WD + RD pairs can help setting con­ straints on the dynamical parameter a, 1000 - \;0 0419-487 and then help a better understanding of binary evolution, with implications also

o'---'~~~---L~~~~--L~~~~-L~~~~---'-~~~~ for the DD systems that may lead to a / supernova explosion, and in which we 3500 4000 4500 5000 5500 6000 Wavelength are most interested. Figure 4: The blue spectrum of WO 0419-487; note that also in this case the fines of the red The Perspectives companion are prominent in the red side of the spectrum. The Balmer fines are much narrower than in the case of the two other WOs; this is indeed a low surface gravity WO with a mass of Our spectroscopic study of WDs has only -0.29 MQ. so far singled out four very interesting objects: The two WD + RD pairs and the two WDs with large radial velocity differences. This is quite a high propor­ lation between the three available pair of several TiO bands typical of M-type tion when we consider that just 20 ob­ spectra gives tJ. V = 220, 176, and stars, as weil as the prominent Ha emis­ r jects have been investigated so far. The 45 km S-1 respectively, weil above our sion. The two blue spectra that are avail­ natural developments of this project are estimated errors. Figure 2 shows two able do not indicate an interesting then both extensive as weil as intensive: spectra of WD-17, where the shift in line change in the radial velocity. However, On the one hand we need in fact to position is also apparent. We conclude the Balmer emission suggests that the enlarge the sampie of observed WDs, that two new DD systems may have RD component may be filling its Roche on the other hand, we would like to been discovered, one of which exhibits lobe, in wh ich case the system should obtain further data for those systems radial velocity changes wh ich are inter­ be very close. Orbital shrinking by both which are found to exhibit special estingly close to those expected for SN I a common envelope stage and subse­ characteristics. For example, we need progenitors. Glearly, further observa­ quent gravitational wave radiation might to determine spectroscopic periods, as tions are required to confirm these pre­ be blamed for the present configuration weil as whole radial velocity curves. In liminary results, and to determine the full of the system, wh ich appears to have particular, the system WD-17 deserves tJ. Vr amplitude and period of the two DD the characteristics expected for a cata­ further study. Its large radial velocity candidates. clismic binary caught in quiescence. variation makes it the best DD candidate We now turn to the dwarfs with com­ The other WD + RD pair (WD 0419­ for a SN I progenitor, comparable to posite spectra. The spectrum of two 487) is also an exciting object. Its spec­ WD0135-052 whose DD nature had programme WDs showed in fact clear trum is shown in Figure 4, where TiO been anticipated by Greenstein (1986) signs for the presence of a RD compan­ bands are also evident at longer and spectroscopically confirmed by ion. These two WDs were indeed al­ wavelengths where the RD light Saffer et al. (1987). In conclusion, mod­ ready listed by Probst and O'Gonnell dominates. We have then applied the ern instrumentation and techniques (1982) among the seven WD + RD pairs cross correlation technique separately allow a fresh approach to apparently that they have discovered photometri­ for the blue (3850-4650 A) and yellow unconspicuous objects such as white cally, thanks to their IR excess. The in­ parts of the spectrum (5000-5500 A), and red dwarfs. Although certainly all clusion of these two objects in our pro­ so as to obtain the tJ. Vr of the WD and of WDs had a bright past, most will just gramme was however unbiased, and the RD, respectively. This gives tJ. Vr = fade away as time goes by. Few among -124 km S-1, and + 232 km S-1, indicat­ then we can refer to a spectroscopic them, however, might have an explosive rediscovery of their binary nature. This ing that the mass ratio is close to 0.53. future. Time is ripe for a closer look on will be relevant for the statistical aspects Since the mass of the WD primary is these dwarfs. of the present research. In any event, known from surface gravity determina­ the spectroscopic study of the two pairs tions to be - 0.29 MG (Koester et al. Acknowledgements is revealing some very interesting prop­ 1979), we can then estimate the mass of erties of the systems. the RD secondary as - 0.15 MG' The We would like to thank Dr. Luciana In one of them (WD0034-211, ob­ system appears to have evolved Federici, Paola Focardi, and Bruno served in the first run) a faint emission through a Gase B Roche-Iobe overflow, Marano for a preliminary manipulation of core was evident in all the Balmer lines, when the primary was experiencing its the spectra obtained in 1985, and for possibly indicative of an accretion disk red giant phase of evolution, with a their introduction to the use of the originated by transfer of material from a 0.29 MG helium core. We can then cross-correlation programme. We nearby companion. In the second run guess the approximate dimensions of would also like to thank Dr. Jim Liebert we have then reobserved this object in the original orbit, and when knowing the for having made available to us the pre­ the red spectral range (5000-7000 A), present dimensions we could obtain a print of his study of WD 0135-052. using the 0150 grism and calibrating lower bound to the crucial efficiency with the He+Ar lamp. The composite parameter a. References spectrum (from 3600 to 7000 A) is Two very important applications of Canizares, C. R., Fabbiano, G., Trinchieri, G. shown in Figure 3: Note the presence of the WD + RD pairs can then be antici- 1987, Ap. J., 312, 503. 38 Fabbiano, G. 1986, PASP, 98, 525. Matteucci, F., Greggio, L. 1986, Astron. As­ Astron. Ap. in press; ESO Preprint No. 565. Greenstein, J. L. 1986, Ap. J 92, 867. trophys. 154,279. Robinson, E.L., Shafter, AW. 1987, Ap. J Greggio, L., Renzini, A 1983, Astron. Astro- McCook, G. P., Sion, E. M. 1984, A Catalogue 322,296. phys. 118, 217. of Spectroscopically Identified White Saffer, R.A., Liebert, J. W., Olszewski, E. Iben, I.Jr., Tutukov, AV. 1984, Ap. J Suppl. Dwarfs, Villanova University Press. 1987, Bull. A.A.S.19, 1041. 54,335. Nomoto, K., Iben, I.Jr. 1985, Ap. J, 297, 531. Saio, H., Nomoto, K. 1985, Astron. Astro­ Koester, 0., Schulz, H.,. Weidemann, V. Pier, J. R. 1983, Ap. J Suppl. 53, 791. phys. 150, L21. 1979, Astron. Astrophys. 76, 262. Probst, R. G., O'Connell, R. W. 1982, Ap. J Liebert, J., Probst, R. G. 1987, Ann. Rev. As­ Lett. 252, L69. tron. Ap. 25, 473. Renzini, A, Fusi Pecci, F. 1988, Ann. Rev.

A Search tor Flare Stars With the GPO Astrograph R. ANIOL, H. W. OUERBECK, W.' C. SEITTER, Astronomisches Institut der Universität Münster, F. R. Germany M. K. TSVETKOV, Department ofAstronomy, Academy ofSciences, Sofia, Bulgaria

1. Flare Stars in Young Stellar The elose eorrespondenee between A new ehapter for the GPO was Aggregates flare stars in aggregates and in the solar opened with inereasing interest in direet vieinity, as weil as some analogy be­ plates. Most of them are used for as­ The first systematic investigation of tween the flares in these stars and the trometrie purposes: for asteroid obser­ flare stars by the Mexican astronomer solar flares suggest theoretieal models vations, proper motion studies, and for Guillermo Haro marks the beginning of for the underlying physieal proeess(es). the input eatalogue of the HIPPARCOS long-term, multi-site studies of these However, a general pieture whieh in­ astrometrie satellite mission. Reeent objects in young stellar aggregates. His eludes all aspects of the problem of highlights of the southern sky, Halley's first series of observations with the eorrelating the properties of flare stars, Comet and the Supernova 1987 A, Schmidt telescope of the Tonantzintla UV Ceti stars and older main sequenee tempted more observers to "try" the Observatory already showed the elose stars, has not yet been developed. GPO. We will report on one of our two relation between stars in the T Tauri photometrie surveys, the seareh for flare stage and eool dwarf stars with en­ stars with the GPO. Figure 1 displays haneed UV aetivity (= flare stars), both 2. The GPO Astrograph the observational aetivity of the last two types of objeets being still in their early The GPO astrograph was one of the deeades. stages of evolution. first teleseopes on La Silla. At its previ­ After 30 years of subsequent investi­ ous loeation Zeekoegat in South Afriea gations, some of the eonclusions by 3. The Suitability ot the GPO tor a (1961 -1966) and during its first years Haro are substantiated on the basis of Flare Star Search on La Silla, it was mainly used to obtain extensive observational material, while objeetive prism speetra of stars in the Until now, patrol observations of flare the new results also gave a elearer pie­ direetion towards the Magellanie star fields were predominantly earried ture of the nature of flare stars. Some of Clouds. One of the rare speetra of Sk out with Sehmidt teleseopes in the the main aehievements and results to be - 69° 202, the star that later beeame northern hemisphere. The Sehmidt tele­ substantiated by further studies are: SN 1987 A, was taken with this instru­ seopes in the southern hemisphere - More than 1,000 flare stars in the ment. were to a large extent oeeupied by atlas aggregates of Orion, the Pleiades, NGC 7000 (North Ameriea nebula), Praesepe, Coma, and others were diseovered. - The physieal eorrespondents of flare stars in aggregates are the UV Ceti stars in the solar neighbourhood. - By observing multiple flares in given stars, the lower limit of the number of 1200 flare stars in various aggregates eould be estimated: Orion - 1,500; co'- Pleiades - 1,000, NGC 7000 - 400, Q) >- Praesepe - 300, Coma - 100 stars. Cf) 800 - It was found that, for a given aggre­ -....Q) co gate, the speetral type of the bright­ Q. est flare star and the maximum flare amplitude both eorrelate with the age of the aggregate. 400 - In stellar evolution, the T Tauri stage is followed by the flare star stage. - There are two major types of flares, fast and slow ones. It is eoneluded that all flares have the same physieal 1970 1975 1985 origin, sinee there are also rare flares Figure 1: Number of plates per year taken with the GPO telescope since the beginning of its of intermediate speed. operation on La Silla. 39