N95- 27078

13

DISCUSSION ON SELECTED SYMBIOTIC

R. Viotti and M. Hack

I. INTRODUCTION made of its variations" (Mayall, 1969).

This pessimistic remark should be consid- Because of its large variety of aspects, the ered as a note of caution for those involved in symbiotic phenomenon is not very suitable for the interpretation of the observations. In the a statistical treatment. It is also not clear following, we shall discuss a number of indi- whether symbiotic stars really represent a vidual symbiotic stars for which the amount of homogeneous group of astrophysical objects or observational data is large enough to draw a a collection of objects of different natures but rather complete picture of their general behav- showing similar phenomena. However, as al- ior and to make consistent models. We shall ready discussed in the introduction to the sym- especially illustrate the necessary steps toward biotic stars, in this monograph we are espe- an empirical model and take the discussion of cially interested in the symbiotic phenomenon, the individual objects as a useful occasion to i.e., in those physical processes occurring in describe different techniques of diagnosis. the atmosphere of each individual object and in their time dependence. Such a research can be I1. Z ANDROMEDAE AND THE DIAGNOS- performed through the detailed analysis of TICS OF THE SYMBIOTIC STARS individual objects. This study should be done for a time long enough to cover all the different II.A. INTRODUCTION phases of their activity, in all the spectral ranges. Since the typical time scale of the symbiotic phenomena is up to several and Z And has been considered as the prototype decades, this represents a problem since, for of the symbiotic stars, from its light history and instance, making astronomy outside the visual the spectral variation during outburst. For a region is a quite new field of research. It was a long time, its has been the basis for fortunate case that a few symbiotic stars (Z theoretical studies of the symbiotic stars. The And, AG Dra, CH Cyg, AX Per, and PU Vul) optical spectrum of the has been studied had undergone remarkable light variations (or extensively since the pioneer work of Plaskett "outbursts") in recent years, which could have (1928). Of particular interest is the work of been followed in the space ultraviolet with Swings and Struve (1941; see also Swings, IUE, and simultaneously in the optical and IR 1970) describing the spectral evolution during with ground-based telescopes. But, in general, and after 1939 outburst, and of Boyarchuck the time coverage of most of the symbiotic (1968), who discussed the 1960 outburst and objects is too short to have a complete picture the following decline phase. Ultraviolet obser- of their behavior. In this regard, one should vations started after the launch of the Interna- recall Mayall's remark about the light curve of tional Ultraviolet Explorer (Altamore et al., Z And: "Z Andromedae is another variable that 1981) and continued to the time of writing the shows it will requires several hundred years of present review. In particular, this allowed the observations before a good analysis can be study in the UV of the two minor outbursts that

663 occurredin 1984and 1985.Theseobserva- emissionlines of several different atomic spe- tions,togetherwith theextensivemonitoring cies and with a large range of ionization en- of theIR andof theemissionlineprofiles,still ergy. A compendium of the ions whose transi- makeZ Andoneof thebest-studiedsymbiotic tions have been identified in the optical spec- stars,andanidealtargetfor investigationthe trum of Z And was given in Table 11-2. Both symbioticphenomenon. permitted and forbidden transitions are pres- ent, as well as lines typical of stellar chromo- Thevisualluminosityof Z Andis variable spheres, Be, Of, and WR stars, planetary nebu- betweenV = 8to 11ontimescalesfromafew lae, coronal regions, etc. These features are daysto severalmonths,withoutclearevidence also seen in the ultraviolet. Altamore et al. of anyperiodicity.Thefull lightcurveisrepro- (1981) identified low-excitation lines of OI, ducedin Figure11-2in Chapter11.Figure13- MglI, and Fell, and high-ionization lines of 1 is a condensedplot of IO0-daymeansof Hell, OV, and MgV]. But the main result is that photographicand visual observationsfrom the intensity of the emission lines is largely 1887to 1969(Mayall,1969). variable on long time scale. Variation of H(x was reported by Altamore et al. (1979), while Altamore et al. (1982) found that in November 1982, the UV continuum and emission lines Thelightcurveis characterizedby (1)four mainphasesof higherluminositystartingin were about 40 percent stronger than in August 1895,1914,1939,and1959,andby(2)periods 1980. They also noted that the IR spectrum did of lowluminositylastingroughlyonedecadein not show a significant change since 1981. From between.In thefollowing,weshalldiscussthe a more thorough analysis of the UV observa- tions collected since 1978 with the IUE satel- behaviorofZ Andduringquiescenceandactiv- ity. lite, Fernandez-Castro et al. (1988) found that the UV continuum and the emission line fluxes vary quasi-periodically on a time scale of about 760 days. The amplitude is larger for the ll.B. THEBEHAVIOROFZ ANDDURING Balmer near-UV continuum and for most of the QUIESCENCE emission lines, and lower for the far-UV con- tinuum and the CIII] line. The time variation of the SilII]/CIII] ratio suggests a variable mean Themostrecentquiescentphase(1972to electron density of the emitting region from 0.6 about1984)of Z Andwasalsothelongest,and to 2.2 x 10 "_ cm -3 in phase with the UV light thisfact hasalloweda detailedstudyof the curve. At maximum, the density and the emis- behaviorof a symbioticstarduringminimum sion measure are larger, while the effective light. The optical-ultravioletspectrumof Z emitting volume is smaller. Thus at maximum, Andat minimumisrich in strongandnarrow the emission mostly comes from a compact

1_ 1_ 20000 2_ 3_ 3_ 40000 61 I I I I I I - 6

10- -10 8- _ -8 12 - -12 I I I I I I I I I 1890 1900 1910 1920 1930 1940 1950 1960 1970

232848 Z ANDROMEDAE, 1887-1969

Figure 13-1. Plot of the lO0-day average magnitudes of Z And during 1887 to 1969 (from Mayall, 1969).

664 region,andatminimum,theemittingvolumeis lengths, but this needs to be confirmed by fu- morediluted.Theseobservationsareconsis- ture observations. tentwith a modelof highly ionizeddiffuse regionthat is periodicallyocculted,leaving visibleonlytheouterlow-densityregions. II.C. THE ACTIVE PHASES OF Z AND

Muchaboutthenatureof Z And canbe learnedfromthesimultaneousstudyin differ- The active phase is characterized mostly by entwavelengthregions.Figure13-2illustrates an increase of one to two magnitudes of the thevariationsof Z AndfromUVto IRduring visual after a long-lasting quiescent thelastquiescentphase. phase. The brightening is rather slow, the rise time being about 100 days/mag, even com- ThelargeUVvariabilityisclearlypresentin pared with the slow novae. This "outburst" is theU-bandandin theHcf-variationandcanbe generally followed by a sequence of minima tracedbackto beforethefirstultravioletobser- and maxima resembling a damped oscillator. vation.Thevariationis lessevidentin thevis- The time interval between two successive ual.Duringminimum,V variedbetween10.2 maxima (or minima) is not constant, but varies and11.0in anapparentlyirregularmode.This from 31(1 to 79(1 days (Mallei 19781 in an ir- variationduringminimumhasbeennotedby regular way and is not in phase with the UV- severalauthors,(see Kenyon,1986),and variability during quiescence discussed above. KenyonandWebbink(19841,fromtheanalysis of all theminimaduringquiescence,foundthat During the rise to outburst and the subse- theyareclearlyperiodicwithameanperiodof quent oscillations, the optical spectrum under- 756.85daysin agreementwiththeUVandHa goes large changes, which have been exten- results.The radialvelocityof the M-giant sively described in a number of papers. As the componentwas measuredby Garcia and stellar increases, the high-ioniza- Kenyon(19881overabout6years.Theyfound tion lines fade, and at maximum light, the spec- a smoothedsinusoidalcurvewith a periodof trum displays a strong blue continuum with 750+ 8 days, and a semiamplitude of K = 8.1 prominent emission. These lines + 0.5 km s t. This fact provides further evidence have absorption cores that dominate the emis- of binarity, but the and photom- sion at the higher members of the Balmer se- etric curves do not have the relative phasing (a ries. In some cases, profiles are present quarter of phase) expected from eclipse or re- in the H and Hei lines. At maximum, the ab- flection effects. This might suggest a more sorption bands of the M spectrum are hardly complex geometry of the system, as also sug- visible. However, the weakening of the cool gested by the ultraviolet observations as dis- spectrum is only apparent since, as it has been cussed below. Finally, concerning the , found, for instance Boyarchuck (1968), the Taranova and Yudin (1981) and others reported TiO bands are only veiled by the enhanced blue small fluctuations, clearly not in phase with the continuum, while the luminosity of the cool UV variations. These can be attributed to the component has not changed within the errors. irregular behavior of the late-type component The high-ionization lines and the TiO absorp- that is a common feature of the normal (normal tion bands strengthen again as the blue contin- = not a symbiotic or peculiar system) late-type uum decreases during the light fading. giant and supergiant stars. The little variability of the cool component was also noted by Alta- Recently, two minor outbursts of Z And more et al. (1979, 1981) who, from their analy- were reported. After nearly 12 years of quies- sis of a collection of blue-infrared objective cence, the star brightened to V = 9.6 in March prism plates taken during October 1977 to June 1984 and again to V = 10.1 in September 1985 1979, found that the near-lR continuum re- (Mattei, 1984, 1985). The 1984 outburst dis- mained constant within +0.1 mag. Larger vari- played a general rise of the line intensity. The ations seem to be present at longer wave- largest increase was measured for the high-

665 ionizationlinesof Hell, CIV,andOiii, while 19790 1981 0 19830 NV remainedunchanged.Of particularinterest I is thebehaviorof thecontinuum:longwards NV 1400A theflux appearedlargerthanprevi-

11 25 ously,whileit wasweakerin thefar-UV,indi- I075 I catinga generaldecreaseof thecontinuum 1 temperature.This result is not unexpected, since,asdiscussedabove,generallyatoutburst the excitationof the spectrumdecreases. C IV

However,thisbehaviorwasnotfollowedbythe 1050 I OI emissiontripletat 1302-06A, whichmark- - 1100 I edlyfadedduringoutburst.Thisresultcould

O Ill) • l I_ T l _ r - 11 25

1

Ic_ Hu Am(_ - 1175 i I \ 0

11 25 11

C I_l] •

12 11 75

11

12 S_ Ill] • -1150

- 1200

I I 4OOO 50OO 1 23_ JULIAN DAY

165u 5 H 19790 1981.0 1983.0 '_----_____.__..._e_- ----_ I l l

6 115 23p

2900 A - 120

4 L 36. __ - 12.5

< (5 4 9_ ., o,

115 1336 A 3 N

4 - 120 */ /

I {I I l I I I I 40O0 5000 JULIAN DAY_ JULIAN DAY

Figure 13-2. Multifrequeney monitoring (_'Z And during quiescence. (a) Near-UV to IR and Hot intensity (Altamore et al., 1979). Taranova and Yudin (1981). Lunel (see Altamore et al., 1981). Ultraviolet line and continuum fluxes during 1978-1982 (Fernandez-Castro et al., 1988).

666 probablybe relatedto thefact thatthe 1984 smallerthantheopticallythinvalueof two,and eventwasnot similarto the mainoutbursts in onecase,evensmallerthanone(February observedin thepast.In themonthsfollowing 1986).Thisresultwill bediscussedbelowin theoutbursts,aftertheluminositydecline,OI SectionII.D.Strongbroadwingsareevidentin andNV strengthenedin theUV (Cassatellaet theJune1985spectrum,while in theother al.,1984).At highresolution,thebluewingsof phases,theyaremuchweaker. CIV andNV displayedan absorptionlineat -120kin s_,whichsuggeststhepresenceof a II.D. DIAGNOSTICS low-velocity,warmwindlikethatobservedin AG Dra(Viottiet al., 1983).Thehigh-resolu- tionUVmonitoringofZ Andaftertheoutburst Oncea consistentamountof homogeneous revealedlargevariationof theemissionline datais availablefor a giventarget,including profileswhichhavebeennotedsinceApril alsoitstimebehavior(andtakingintoaccount 1984.Figure13-3showstheevolutionof the all thepossibletimescales),it isthenpossible CIV doubletsincethe 1984outburst(Cas- tomakethenextstep,thatis,totry to buildup satellaetal., 1988).Evidentin thefigureisthe an empiricalmodel.In general,theIUE ar- differentshapeof thelinesatdifferentepochs. chivesprovide|or mostof thesymbioticob- PCygniabsorptionsareclearlypresentin three jects(andfor manyothercategoriesof astro- spectra.Theintensityratio of 1548/1551is physicalobjectsaswell)thebesthomogeneous

3. O0 JUNE 1985 jUNE 19e4 1.80 I ml03l .103

.801 I. 50 ,J IO0,S o.oo _'_,'_ e,,_'.,a..a,.,._,_ 1540.0 IS$O.O tseo.o 1540.0 JSSO.O ISSO.O

t. 20 8.001 FEB. 1988 OEC. 1988

mlO 3 I103 I

• 8O ,.oo| x f "-i I _, ,

k 0.00 0.00 I ..... l_J 1540.0 1"1'10.0 1_180.0 ISdO.O |qSO.O 1510.0 LAHBOA L _80A

Figure 13-3. Evolution of the ClV doublet in Z And afier the March 1984 outburst (Cassatella et al., 1988). Note the P Cygni absorptions which are clearly present in three spectra, and the strong emission wings in June 1985. In February 1986 the 1548/1551 line flux ratio is smaller than the optically thick limiting value of one.

667 'iO I I I set of calibrated, good quality data. In the a 30 N lit] 4po multiplel s _b_ ._ meantime, several theoretical computations of T = 810 _ °K the atomic parameters of important UV transi- 06 j_/S tions have been recently performed, so that the 20 • c line intensities can be used to derive the physi- 04 sl 4 cal parameters of the emitting region. In the 10 following, we shall discuss the case of Z And 02 also as an illustration of the impact of the new UV observations on our knowledge of the na- o I [ N l ,0 10 _ 10" 10 _° 10 _ 10 _: ture of peculiar objects such as the symbiotic stars. Figure 13-4. The NIII multiplet ratio in Z And and RR Tel (AItamore et al., 1981). The different curves give as As discussed above, the optical and UV a function of the electron densiO, for Te=8xlO 4 K the theoretical line ratios." (a) 1754.0/17522; (b) spectrum of Z And includes prominent emis- 1748. 7/1749. 7 (dashed cur_,es, right-hand scale)," (c) sion lines of both permitted and intercombina- 1748.7/1749.7," (d) 1746.8/1748.7 (solid curves, tion or forbidden transition of different ions. In left-hand scale). The observed ratios for Z And (dots) principle, this should be used to determine the and RR Tel (triangles) are indicated. electron density of the emitting region(s). In addition, the presence of different ionization (Nussbaumer and Schild, 1979). This is again stages of the same element and of ions of differ- in agreement with the above results, and might ent elements can be used as diagnostics of the also suggest that all the intercombination lines temperature and chemical composition. Table are formed in the same region. This may not be 13-1 summarizes the main line ratios, that can true for all the symbiotic stars. In fact, Nuss- be used to derive the physical parameters of the baumer and Schild (1981), in their analysis of symbiotic system. the UV spectrum of the symbiotic V1016 Cyg, found from the CIII] and NIVI double Altamore et al. (1981) analyzed the NIIII ratios an electron density of 2 x 106 and 4 x 106 multiplet near 1750 A and determined an elec- cm 3 respectively. But the intensity of the NIII] tron density (for the N++ region) of (1.7 _+0.9) multiplet would require a much higher N of x 10_° cm -3 (Figure 13-4). The line ratio of this about 10 +9 cm -3, which could be related either to multiplet is sensitive to changes in the electron a large density gradient in the emitting region density in the range from 10 +9 tO 10 nl cm -_, but or to a particular excitation mechanism. is rather independent of the electron tempera- ture (Nussbaumer and Storey, 1979). For in- Altamore et a1.(1981) also considered the stance, the NIII] line intensities in the symbi- CIII]/NIIII intensity ratio, and found an elec- otic novae RR Tel and V1016 Cyg suggest a tron density of 1.9 x 10 _c_and 1.2 x 10÷9cm -a for smaller density of 1.5 x 10÷9 (Altamore et al., Z And and RR Tel, respectively, consistent 1981) and of about 10+9 cm 3 (Nussbaumer and with the results obtained from the NIII] multi- Schild, 1981), for RR Tel and VI016 Cyg, plet. More recently, Fernandez-Castro et al. respectively. As shown in Table 13.1, the elec- (1988) have used the SilII]/CIII] flux ratio in Z tron density can also be determined from the And to give an estimate of the electron density relative intensity of the OIV] and CII] multi- during different phases of its UV variability. plets. From the observed OIV] 1404.81/ They found N to be variable in the range from 1401.16 ratio of 0.38, Altamore et al. derived 0.56 to 2.2 x 10 _°cm _, again in agreement with log N e = 10.1, in good agreement with the NIII] Altamore et al. results. However, when the flux estimate. The CII] multiplet is too weak in Z ratios of the lines of different ions are used, the And for density diagnostics. Only the high- derived density estimates largely depend on the density component of the CIII] doublet is vis- adopted electron temperature, as well as on the ible in the UV spectrum of Z And, which pro- ionization equilibrium of each ion and on the vides a lower limit to Ne of about 10÷_ cm 3 C/N abundance ratio. Singly ionized

668 TABLE 13-1. ULTRAVIOLET INDICATORS OF THE PHYSICAL PARAMETERS IN SYMBIOTIC STARS.

ions lines parameter range ref.

CII] 2325.4/2328.1 N 107 - 10 _ ( ] ) .... 2325.4/2326.9 ...... 2324.7/2326.9 ......

NIII] 1754.0/1752.0 N 10"- 10 '_ (2) .... 1748.7/1752.0 ...... 1748.7/1749.0 ...... 1748.7/1746.8 l0 s- 10 j_'

OIV] 1407.4/1401.2 N 10L105 ; 10_-10 l' (3) 1401.2/1404.8 10L105 ; 10_-10 "_5

CIII] 1908.7/1906.7 N 10 _- l0 b (4)

NIV] 1486.4/1483.3 N c 104 - 1() '_1 (5)

SillI] 1892.0/I 882.7 N 3.10 -' - 3. l0 s (6)

[NelV] 1601/2423 N c 104 - l0 s (7) ...... T (for NI0 _)

IFeVII] 2015/3759 N 107- 10"' (7) ...... T (for N<107 or >10 "_)

NIIII/CIII] 1749.7/1908.7 N 10 _ - 10 II (I)

SiIIII/CIIII 1892.0/1908.7 N 10'_= 10 II (1)

CII/CIlI] 1336/1909 T (9) (10)

CIll]/C1V T 104- 3x104 (9) (10) 1718/1240

NIV]/NV T (9) (10)

Oml/IOIIll 1661/6300 T (5)

Silll/Silll] 1299/1892 T 1.2 - 6.104 K (6) (for N = 10"-I0 z_>)

HeII 1640.4/Fc(I 336) T* (1 I)

Notes to the table. (I) Stencel et al. (1981). (2) Altamore et al. (1981). (3) Nussbaumer and Storey (1982). (4) Nussbaumer and S('hild (1979). (5) Nussbaumer and Schild (1981). (6) Nussbaumer (1986). (7) Nussbaumer (1982). (8) Nussbaumer and Stern'el (1987). (9) Stikland et al. (1981). (10) Nussbaumer and Storey (19841. (11) Fernandez-Castro et al. (1988).

669 and have ionization energies of 24 and velocity difference between the high-ioniza- 30 eV respectively, which are different enough tion resonance lines and the intercombination to give significant differences in the ionization lines. From a study of several symbiotics fractions C++/C and N++/N. This also depends Friedjung et al. (1983) found a systematic red- on whether the ionization is by electron colli- shift of the former lines with respect to the sions or by photoionization. Altamore et al. latter, of the order of +10 to +20 kln s -_. They (1981) have computed the CIII]/NIII] ratio interpreted this as the result of radiative trans- assuming ionization by electron impact like fer in a warm, low-velocity expanding me- that in the solar transition region, and for a dium, which should be identified with either cosmic C/N abundance ratio. This last assump- the cool star wind, or with its base. This model tion can be justified from the fact that the study is also supported by the larger width of the of several symbiotics--including Z And--by higher ionization lines observed in some sym- Boyarchuk (1970) in the optical, and by Nuss- biotics, as discussed in Chapter 11 Sections baumer et al. (1988) in the ultraviolet generally IV.D. and VIII.D. Friedjung et al. found that in suggest a close to cosmic abundance for the Z And, the radial velocity difference is proba- emitting regions in symbiotics. Altamore et al. bly variable but always positive well beyond also found that the CIII]/NIII] ratio could be the measurement errors. Fernandez-Castro et largely affected by radiation field (e.g., the al. (1988) confirmed this result and found a diluted hot stellar radiation). The flux ratio range of variability between +15 to +30 km s-_, observed in Z And would imply an upper limit but without a clear correlation with the UV line of 1.5 x 10 _6erg cm 3 Hz _ for this radiation in and continuum flux variations. the N++ emitting region, or a dilution factor smaller than 7.5 x 10-_. To decide about possible models, better would be to have a direct estimate of the elec- In a planetary -like model of symbi- tron temperature. In most cases, this is diffi- otic stars, such as the one proposed by Nuss- cult, since the temperature indicators generally baumer and Schild (1981) for V1016 Cyg, the are weak lines, and their flux ratio may also electron temperature should be close to 104 K. depend on other parameters. Stikland et al. Photoionization models for Z And were dis- (1981), Nussbaumer (1982), and Nussbaumer cussed by Fernandez-Castro et al. (1988) and and Storey (1984) discussed, among others, Nussbaumer and Vogel (1988). Alternatively, methods of electron temperature determination the high-ionization lines could be produced in from emission line ratios. Some useful T e indi- a solar-type transition region with much larger cators are listed in Table 13.1. Fernandez-Castro et T e, as for instance suggested by Altamore et al. al. (1988) derived the electron temperature for Z (1981) for Z And. This latter hypothesis is And from the NIV ! 7 ! 8/NV 1240, CII 1336/CIII] based on several arguments: the small width of 1909, and CIII 1176/(7III] 1909 flux ratios. They the emission lines implies formation in a low- found T -- 15,000 K, lower than the value as- velocity dense region (which excludes line sumed by Altamore et al., (1981) which seems formation in a high-velocity hot star wind, and to favor a photoionization model. However, the in a rotating disk). Comparison of the derived above flux ratios involved emission lines that CIII] emitting volume of I x 10÷_vcm _ with the are rather faint in the UV spectrum of Z And. SilV maximum line thickness indicates that the Thus, the T e estimates of Fernandez-Castro et emission should come from a thin shell, rather al. are rather uncertain, and could well be lower than from an extended sphere. In addition, as limits. discussed above, the stellar radiation in the

NIII] emitting region should be very weak, Once the electron density and temperature hence the region far from the hot star. are known, the line fluxes can be used to derive the total amount of the emitting ions, and the Another observable which could put con- size of the emitting regions. Altamore et al. straints on the possible models is the radial (1981), assuming the C++ and N++ regions

670 homogeneousand transparent,obtainedan and continua provides an estimate of emittingvolumeof 1.0x 10 36 and 8.2 x 1035 cm 3 the emission measure Ne2V. Assuming T c = for C++ an N++ respectively. These values 15,000 K, Fernandez-Castro et al. (1988) found obviously depend on their assumption of a Ne2V equal to about 2-6 x 1()+_ cm _, a value solar-type emitting region with T-- 80,0(10 K. much larger than the above derived value of 4

It is easy to see that such a warm region cannot X 10 +57 cm _ for C++, N++ region, assuming be responsible for the ultraviolet continuum cosmic abundance (Altamore et al., 1981). observed in Z And. This implies the presence of Perhaps the intercombination lines and the an optically thick region (or disk) or a hot star Balmer continuum are formed in different re- with a radius of about 2 x 10 m cm. The tempera- gions. ture of the hot source can be obtained using the Zanstra method. Assuming that the emitting Variability may give important information, region is optically thick to the hot source con- especially about the spatial structure of the tinuum Fc shortward of 912 _ (ionization from system, if the variations are associated with an the ground level of H) and of 228 A (He+ ioni- orbital motion of a binary system. In fact, ob- zation), and that the hot component of Z And servations at different epochs allow one to radiates as a blackbody, we have for a case B observe the system with different lines of sight. He++ recombination (Fernandez-Castro et al., The best period is when the symbiotic system is 1988): in quiescence, so that the periodic phenomena are not masked by the symbiotic activity, (13.1 whose time scale is generally comparable with l(Hell 1640)=3.303x10-2s T _ x the . Fernandez-Castro et al. Fc(2') (1988) investigated the periodic UV variability s(x,, during quiescence and found that the UV con- tinuum and emission line fluxes are variable on

where it is assumed that the Hell 1640 A line is a time scale of about 760 days, with maxima optically thin, and f(_.",T) is the relative num- and minima in phase with the UBV and H¢_ variability (see Figure 13-2). From the time ber of photons provided by a blackbody at a variability of the SilII]/CIIII flux ratio, Fernan- temperature T between _, = 0 to _," = 228 A. dez-Castro et al. suggested that the mean elec- Using the continuum flux at _,' = 1336 A and the tron density in the emitting region varied from Hell 1640 A intensity, dereddened for E(B-V) = 0.35, Fernandez-Castro et al. (1988) derived 0.6 to 2.2 x 10 TM cm _, being higher at maxi- mum. This can be explained with a model of a Hell temperature of about 105 K. This tem- line formation in an asymmetric nebula near perature remained nearly constant in spite of the cool giant, which is photoionized by the UV the large UV variability. In fact, the NV/CIV flux ratio, which can be considered as an ioni- photons of the hot star. This nebula is occulted at minimum, and emission is seen only from the zation temperature indicator, did not signifi- outer, less dense parts. This fact and the small cantly change during the period studied by Fer- nandez-Castro et al. size and high density of the emitting region suggest that it can be identified with the inner parts of the cool giant wind ionized by the hot The hot continuum is responsible for the star radiation. Therefore, the UV emission short wavelength UV continuum. But there is lines are mostly formed in the cool giant wind, an excess of the continuum flux in the range from 1500 A to the visual that can be attributed as also is suggested by the systematic radial to bound-free and free-free hydrogen emission, velocity difference between the high ionization resonance lines and the intercombination lines as indicated by the marked Balmer discontinu- discussed above. ity at 3650 A (Altamore et al., 1981; Blair et al., 1983). Continuum recombination from He++ may also contribute to the UV. Fitting of the observed Balmer continuum with hydrogen

671 II.E.POSSIBLEMODELSFORZ a rejuvenated is favored. In any ANDROMEDAE case, its structure could have been largely af- fected by the mass-transfer processes during the earlier stages of evolution of the binary Let ussummarizeour present knowledge system. In particular, we do not know if it pres- about the Z And system. The UV to IR energy ently has the same mass and radius of a white distribution is characterized by two maxima, dwarf, and this makes our modelling rather one in the near-IR, and another in the unseen uncertain. In Z And, the dwarf component EUV (Fernandez-Castro et al., 1988). During seems to accrete mass from the cool star wind, the active phases, the second maximum is not through overflow. The accreted shifted to longer wavelengths in the UV or in matter could form a disk around the dwarf star. the optical. In the latter case, the apparent But so far, there is no evidence of an "outburst" in the visual is larger. According to disk in Z And, whose presence could, for in- Fernandez-Castro et al., the integrated fluxes stance, be indicated by broad high-ionization of the two spectral "bumps" are comparable, around 450 and 880 L for the hot and cool emission lines. This problem will probably be O solved with future high-quality observations of components, respectively. The "nebular" f-f and b-f continuum contributes another 44-141 the emission line profiles in the visual and UV. Mass accretion M onto the degenerate star L ° to the total power from Z And. Therefore, results in an increased luminosity L c of the the radiative power emitted in the UV is not a star. The actual values of M and L critically small amount of the visual-lR power. As dis- depend on many poorly known parameters, the cussed in Chapter 12, this point rules out the model of active cool star, in which the emission mass-loss rate and wind velocity of the cool giant, the orbital parameters, the mass and line spectrum and blue continuum are the result radius of the dwarf star. Thus, our picture of Z of a large surface activity of the cool giant. And (and of all the other symbiotic stars, as Thus, the double-bump energy distribution is well) is necessarily limited because of the a strongly suggestive of binarity. This hypothe- priori assumptions we have to make. The mat- sis is also supported by the optical and UV ter flowing from the is ionized by the variability during quiescence, and by the long- hot star radiation, and emits in the radio. The term radial velocity variations of the cool com- radio flux from Z And observed by Seaquist et ponent, even if the phasing is not the expected al. (1984) has a spectral index ot = +0.62, close one. An important point for the following dis- to the theoretical value for a photoionized wind cussion is whether the cool component is filling (cf. Seaquist and Gregory, 1973). Using the its Roche lobe. Taking into account the pro- formulations of Wright and Barlow (1975) and posed orbital period of 750 d, the luminosity taking for the wind velocity the value of 40 km class of the cool component, which implies that s', and the observed radio flux at 4.885 Ghz, its mass should be of a few solar masses Fernandez-Castro et al. (1988) derived for the (Querci, 1986), and assuming for the hot com- cool giant a mass-loss rate of about 2xlO 7 M o ponent the mass of a white dwarf, it can be yr', in good agreement with the typical values easily found that the cool giant is well inside its critical Roche surface. A different conclusion of M giants (e.g., Goldberg, 1986). Using their parameters for the Z And system, Fernandez- will be drawn in the case that the Z And system Castro et al. also derived an accretion rate of contains a cool bright giant, for instance, as suggested by the IR spectroscopic observations 4.5 x 10 '_ M ° yr', and an accretion luminosity of Kenyon and Gallagher (1983). This again of 1.2 L ° . This luminosity is about two orders stresses the importance of future detailed stud- of magnitude lower than the recombination ies of the cool spectral component, in order to continuum, and, therefore, it cannot be ac- derive its . counted for by accretion processes. Con- versely, accretion seems not to play an impor- The nature of the hot component is so far tant role in the energy budget of Z And. It can unveiled, but a model of a hot subdwarf, or of anyhow be responsible for some of the symbi-

672 oticphenomena,and,in particularfor there- sidered that the particle density in the wind current"outbursts".In fact,theabovederived rapidly falls down outwards. Since the contin- accretionrateis wellbelowtheminimumac- uum and line emission in the HII region is

cretionratefor steadyburningdiscussedin proportional to N 2, the regions near the ioniza- Section12.IV.A,butimpliesoccurrenceof re- tion front are the main contributors to the nebu- currentthermonuclearevents• lar spectrum. Therefore, in deriving the physi- cal parameters of the emitting region, it is A schematicmodelfor Z Andduringquies- crucial to take into account the geometry, cenceisshownin Figure13-5.TheUVphotons which could be far from homogeneity and from from thehotcomponentionizethecool-star spherical symmetry. In addition, because of the winduntila limitingsurfacewhichis separat- different dependence on the electron density, ingtheHI andHII regions•Accordingto Tay- the mean regions of formation of different lor andSeaquist(1984),theshapeof thissur- emission lines and continua could be signifi- faceisdeterminedbyaparameterX definedas: cantly different. The corresponding electron densities and emission measures, therefore, can be largely different, even if the lines are (13.2) formed in the same medium. This can explain the different values found in Z And, as dis- X = (4n;aL h/or) • _2m21n.(V/l_)2 cussed above, and in many other symbiotic stars as well. Formation of emission lines in a where a is the binary separation, Lph, the num- low-velocity medium, such as the red giant's ber of hydrogen ionizing photons per second wind, is also in agreement with the line narrow- from the hot component, _, the recombination ness. According to Fernandez-Castro et al., the coefficient to all but the ground state of hydro- observed time variability of the line fluxes gen, _, the mean molecular weight, m., the should be the result of the orbital motion: the hydrogen mass, and V and 1_, the red giant's effective emitting volume is different, if the wind velocity and mass-loss rate, respectively. line of sight is different. As discussed by Nuss- For Z And, Fernandez-Castro et al. (1988) ob- baumer et al. (1986) for V1329 Cyg, one should tained X = 14, implying that the ionization thus expect a "modulation" of the line profile front is close to the red giant surface with a with the orbital phase. Observations have not shape as shown in the figure. It should be con- yet shown this effect, at least during quies-

Figure 13-5. Schematic binary model for Z And, according to Fernandez-Castro et al. (1988)•

673 cence,probablybecauseof theirfairly poor lines of NV, and SV in the UV, and of [FeV], quality.[Profilechangesasthosereportedby [FeVII, and the 6830 A feature in the visible. Cassatellaet al.(1988)shouldbeattributedto HelI is quite strong in emission in both the the stellar"activity",ratherthanto orbital spectral domains. Also, the far-UV continuum motion.]Partialoccultationfromtheredgiant is very hot and intense. Moving to even shorter maypartlyaccountfor theobservedfluxvari- wavelengths, we find an intense X-ray tlux, the ability.Theeffectobviouslydependson the most intense among symbiotic stars (taking (unknown)inclination.Alternatively,theorbit into account the X-ray spectrum and the inter- couldbe eccentric: in this case, one would stellar absorption). This is also in contrast with expect a larger nebular emission and a higher the fact that X-rays were mostly detected in D- electron density at the passage near the perias- type symbiotics, while AG Dra is S-type. Per- tron, as observed. However, it is hard to fit all haps, the most interesting aspect of this object the observational data with this fairly simpli- concerns the major outburst (followed by three fied model. Nussbaumer and Vogel (1988) other ones of lesser strength), which occurred recently showed that the coexistence of two in recent times, allowing a detailed study of a winds can substantially modify the ionization symbiotic star during the whole duration of an structure. The presence of two winds will also active phase. modify eclipse effects in different ionization stages. III.B. THE LIGHT HISTORY

In conclusion, the binary model may ac- The light history of AG Dra is quite similar count for many of the properties of Z And. to that of Z And, in spite of the several differ- Being a detached system, the (s) ences between the two objects, as discussed play an important role, both during "quies- below. Robinson (1969), from a detailed study cence" and during the "outbursts". Many of the of its light curve, identified 11 outbursts be- system parameters-inclination, stellar masses, tween 1890 and 1966. Since then, AG Dra etc.--are uncertain and prevent a more thor- remained at minimum until the end of 1980, ough study of the Z And complex. Yet the when it brightened again from V = 9.8 to 8.5 in object appears a very promising target for the a few days. During the following fading phase, study of many important aspects of the symbi- the star had a new minor maximum in 1982, otic phenomenon. But high quality observa- then it reached the minimum luminosity in tions are essential for a real progress on the matter. mid-1983. More recently, two new low-ampli- tude "outbursts" were observed in March 1985 and January 1986 (Mattei, 1987). The sche- III. THE HIGH-VELOCITY SYMBIOTIC matic light curve of the star is show in Figure STAR AG DRACONIS 13-6.

III.A. INTRODUCTION This long-term behavior is close to that of Z And, i.e., recurrent phases of activity followed by long periods of quiescence. As in Z And, The symbiotic star AG Dra has many inter- small amplitude variations were found in AG esting peculiarities with respect to what is Dra during quiescence, with an amplitude in- generally considered as a "classical" symbi- creasing towards shorter wavelengths otic. First, it is a high-velocity (Roman, 1955), (Belyakina, 1969; Meinunger, 1979), and high-galactic latitude object (b= +41", Table which are possibly periodic with a period of 11-12); i.e., it belongs to Pop 11. Its cool spec- about 554 days (Meinunger, 1979). This behav- tral component appears less cool (K-type) than ior is also present in the space-UV, where large in most other symbiotic stars, whereas the amplitude UV variations were observed before excitation of the emission line spectrum is and after the 1980-82 outburst (Viotti et al., qvite high, as indicated by the high-ionization 1984a; Viotti, 1988a). Recently, Kaler (1987)

674 studiedthe variation of AG Dra in II interme- anomaly as discussed by Lutz et al. (1987). In diate and narrow photometric bands, between this regard, Iijima et al. (1987) found that the 3473 A and 8200 A. The observations were absorption spectrum can be classified as G7V, made from March 1977 to October 1980, dur- in agreement with earlier classifications (e.g., ing the quiescent phase preceeding the recent Wilson, 1943), but not in agreement with the active phase, and covered about 2.4 cycles. The red-IR colors. They attributed this discrepancy periodic variations are present in several bands to a metal deficiency of the cool star, as also and are by far the largest in the near-UV with an suggested by its Pop II nature, which may af- amplitude of about I mag. The amplitude de- fect the usual spectral classification criteria. creases towards the red, but the variations are The optical-IR energy distribution at minimum also present to some extent in the narrow bands is, in fact, consistent with a K-giant spectrum. centered on strong emission lines. Kaler also Viotti et al. (1983a) derived a K3-511I spectral noted indications of a "secondary eclipse" of type from the broad-band photometry. During the K giant at some wavelengths. outburst, the K-type spectrum is veiled (Huang, 1982) by the intense blue continuum. Of particular interest is the recent activity of the star with one major outburst in 1980 and The optical spectrum of AG Dra (Figure 13- three minor maxima in 1982, 1985, and 1986 7) shows prominent emissions of H, Hel, Hell, (see Figure 13-6). These events occurred at the OlIl, and of , from FelI up to [FeV] and right time to pertorm multifrequency observa- [FeVI], and the unidentified high-excitation tions of a symbiotic star during activity from features at 6830 and 7088 A (Boyarchuk 1966; both ground-based and space observatories. Bopp and Smith, 1981: Huang, 1982: Blair et This active phase, in fact, occurred when the al., 1983; lijima el al., 1987. I-tot is very strong IUE satellite was fully operational, which gave in emission. Smith and Bopp (1981) aml Oliv- the opportunity to collect a complete set of ul- ersen and Anderson (1982) found large profile traviolet spectra throughout the whole light variations of the line which they attributed to curve (Viotti et al., 1984a; Lutz et al., 1987). In activity on the K-star surface. In the UV, the addition, two X-ray satellites, HEAO-2 and star displays weak intercombination lines of EXOSAT, were operating during this period, CllI, OIIl, OIV, and SilII, and strong permitted and, as discussed in Section I I.IX, X-rays from lines of NV, C1V, and especially Hell 1640 A AG Dra were positively detected on tbur differ- (Figure l l-31a). Other identified species in- ent occasions (Anderson et al., 1981; Cas- clude the low-ionization OI and Mgll lines, and satella et al., 1987). the high-ionization OV and SV lines (Viotti et al., 19831. The Hell 1640 A line presents ex- tended wings, while the NV doublet displays a III.C. THE OPTICAL AND ULTRAVIOLET P Cygni profile with absorption components SPECTRUM extending to -170 km s _ (Viotti et al., 1984a). Two separate components can be identified in the UV continuum: in the near-UV, a flat con- Near minimum, the yellow-red spectrum of AG Dra presents the absorption line spectrum tinuum probably of nebular origin, and short- typical of a luminous K-type star, with several wards of 1600 A, a steep and strong continuum, narrow absorptions of neutral and ionized close to the Rayleigh-Jeans tail of the energy metals (e.g., Huang, 1982: Lutz et al., 1987). spectrum of a hot star (Figure 11-30). However, The relative strength of these absorptions sug- the high-resolution observations do not reveal gests an early-K spectral type and a luminosity any absorption of possible photospheric origin. class III. However, some anomalies are pres- Actually, at such an we ent, such as the strength of the Bali and Srll should expect to observe strong photospheric lines, which might imply a higher luminosity absorptions of Hell 1640 A and of the high- ionization resonance lines of NV and CIV. class, for instance, as suggested by Huang (1982), or, more probably, a composition These features should be completely hidden by

675 1920' 1930 1940 1950 1960 1970 1980 i i i I I i I I I t I I m I pg

o° °:

t • 10 oo

o" , • . • "" ,"'.... • "" =s " . . .'.,; .- .. "['; • ...... ;': • . • .. : ": ._.'. o.,. . 11 • " . . , . o • , • ... ,,, o • -.:,...... -....

-, , i I , , , J l i , , , I , , _ , I i I , ' I " 25000 30000 35000 40000 45000 JD 240000 +

AG DRACONIS

f3

15

17

b. 19

21

23

23 25

4- 27 25

o

_z ,.'5, 29

a 31

_ 31 ....._--....., ,,...... _/."_.:''_,. _ '33

'%'"'° I °'° 33 , 35

37 L _J-.- 39

39 ...... 0.5 mag. [ I HorizJGl . CO. ,,0J,,

Figl_r¢ ]3-5. ((z) The ligh! history ofAG Dr_z from 1890 to 1956 (Robinson, 1966).

676 I I I I MAX MIN

9

10

1 I I I 1979 1980 1981 1982 1983

1985 1986.0 I I

10 m

Figure 13-6. (b) The recent light curve of AG Dra from 1979 to 1986 (Mattei. 1987). Three outbursts have been recorded in November 1980, February 1985, and January 1986, The vertical arrows show the epochs of ultraviolet ( I UE ) and X-ray (HEAO-2, EXOSA T) observations.

677 the prominent nebular emissions. But, accord- tude of K = 5.3 + 0.3 and with about the same ing to the IUE observations of hot subdwarfs of period as that of the photometric one (see Fig- Rossi el al. (1984), we should also expect to ure 11-14). In this case the phase shift between observe excited lines of highly ionized C, N, the radial velocity and photometric curves is and O, and especially FelV and FeV near 1400 the expected one and clearly confirms the bi- A, which have not been seen in the UV spectra narity of the object. of AG Dra. The only absorptions are the inter- stellar ones (see Figure 11-29c), which are rather strong indeed if compared with the low III.D. THE RECENT OUTBURSTS interstellar extinction of E(B-V) = 0.06 of the star (Viotti et al., 19831. This fact, however, is The main outburst that occurred in Novem- not unusual among halo stars. ber 1980 represented an excellent occasion to The radial velocity of the absorption lines is study in many spectral regions the behavior of high, about -14(I + -146 km s t (Roman, 1953; a symbiotic star during an active phase. The Garcia and Kenyon, 1988). This point and the optical photometric variations during outburst high galactic latitude clearly indicate that AG and those of the UV spectrum are described Dra is a halo object. Huang (1982) observed a among others by Kaler et al. (19871 and Viotti variability of the absorption and emission line et al. (1984a), respectively. The outburst was radial velocity. More recently, Garcia (19861 most energetic in the ultraviolet. The ampli- and Garcia and Kenyon (1988) found that the tude of the first rise was of two magnitudes in K-star lines vary periodically, with an ampli- the u band (near 3500 A), and only 0.5 A at 8200 A (Kaler et al., 19871. The IUE observa-

o .... _.... _.... ,.... I .... I .... it.... l .... i tions (Viotti et al., 1984a) show that the UV

o AG DRA 5 AUG 80 " / I continuum underwent a large increase between October and November 1980, at the time of the a" o ol optical outburst, with a subsequent further increase until January 1981. The overall rise 2 ° was of about a factor 10, much larger than in the visual (Figure 13-8a). In the following months, go the trend in the UV nearly followed the visual, with a minimum by mid-1981, and a second maximum in December. Then, the continuum /'( _f)l'l ,Z_.DII flux gradually faded to minimum. A similar WAV[I F N(JTH IAj trend was displayed by the emission lines. The

_' 1 .... ] .... I .... _ .... I .... l"''l'_"l .... l _" main difference was the absence of the secon- AG DHA ', F _ [] 81 dary minimum in the high-ionization NV dou-

x ,s_ blet, and the constancy of the ionization before and after the outburst, as indicated by the NV/ CIV line ratio (Figure 13-8b). This behavior of AG Dra was confirmed by the optical observa- ,,,o

cu tions showing the persistence of the high-ioni- I x D I zation emission lines (e.g., the Hell 4686 A

© ,l*_tnl .... lJ,_,l,_l .... lia,,I .... IJ,,,h line) also after the outburst (see Figure 13-7). ;_lll)(l 4 (,H'JC 4t_[X) r_(X_"J Ill}(#) fJ()_)() fl!)(J(J IfX)() l!lf)()

WAV[ I F [q(JlH iA) An intense X-ray flux from AG Dra was first detected with HEAO-2 in April 1980 when the Figure 13-Z The optical spectrum of AG Dra at two star was at minimum (Anderson et al. 19811. different activity phases. August 6, 1980 and Februat 3, 6, 1981. Note the strength of the Hell 4686 The star was pointed to again with HEAO-2 A line, and of the Bahner discontinuity in both spectra alter November 1980, but a technical problem (Blair et aL, 1983). prevented the observations (Seward, 19851.

678 More recently, AG Dra was observed with 19800 1_10 1_20 "_30 EXOSAT during the 1985 and 1986 outbursts - 10 , I I 1 , 11 i I and during the minimum phase in between. These observations, already described in Sec- tion I I.IX.C, indicate a modulation of the X-

ray flux by the stellar activity, while there is no -11 indication of a dependence on the 554-day

period. J u.. O o It is interesting to note that IR observations -12 collected before and after the outburst only showed small variation, indicating that the K star remained substantially stable during this period (Viotti et al., 1983b, Piro et al., 1985).

-13 %,, i I I I l Radio emission at 6 cm was first detected 10 U-PHASE_ 2 0 3 0 from AG Dra in June 1986 by Torbett and 19800 1981 0 19820 1983 0

Campbell (1987) who found a flux larger than -9 I I I I I I 1 I I 0.5 mJy. Previously, Seaquist el al. (1984) reported only an upper limit of 0.41 mJy, for a He II 1640 6-cm observation made in February 1982.

Torbelt and Campbell also resolved AG Dra -10 into two close components separated by about 1.3". This increased radio activity, and the presence of a structure could be related to the recent activity of AG Dra. It would be interest- -11 II,,,,,,,_ I C IV 1550 -- ing to follow the further development of this f star in the radio. ,Y_

III.E. INTERPRETATION -1

jP __

In contrast with the majority of the other symbiotic object, our observational data on AG N V/C IV _' Dra are fairly complete to make a clear picture ". of the system. First, AG Dra is binary. The cool component is a K giant, which, taking into account the orbital parameters, is not filling its Roche lobe. Assuming that it has the same ] i I I 10 20 30 absolute luminosity of single red giants, for U-PHASE o E(B-V) = 0.06, a distance of 730 pc is derived, Figure 13-8. The ultraviolet variations of AG Dra 500 pc above the galactic plane (Friedjung, during 1979 to 1983 (Viotti et al., 1984). (a) The UV 1988). It should be considered that, if one as- continuum near 1340 A (curve B) and 2860 A (curve sumes that the cool component fills its Roche C) compared with the visual light curve (curve A) derived from the IUE FES count rates. (b) The lobe, it should be a bright giant with a distance variation of the high ionization UV emission lines and of about 3 kpc. Such a large distance and the of the NV/CIV line flta ratio. corresponding large height on the galactic plane are, however, in contrast with the weak- system, Kenyon and Webbink (1984) inter- ness of the interstellar CIV lines in the UV. preted the UV energy distribution of AG Dra as Concerning the other stellar component of the due to a hot star without an . In-

679 deed, the fact that the red giant does not fill its together with the geometrical dilution and the Roche lobe implies that the accretion rate is too increased far-UV opacity, would contribute to low to produce a disk. According to Viotti et al. recombine nitrogen ions to NV and would then (1983), the temperature of the hot component is produce the observed P Cygni profile. Still about 100,000 K. The corresponding effective further out, in almost stationary regions, the radius in mid-1981 was 0.02 R o , typical of a CIV ions are produced that would not show a P white dwarf. A higher temperature of about Cygni profile. In this case, the absence of CIV 160,000 K was also determined by Kenyon and could also be a geometrical effect, for the more Webbink (1984), and Iijima et al. (1987). extended CIV region is not homogeneous and However, the hot component of AG Dra cannot would not hide the stellar disk. It should be be considered a normal hot star, since, as dis- considered that, according to this model, we cussed above, no "photospheric" absorption should expect the OVI lines in the far-UV to lines have been so far identified. have a P Cygni profile broader than that of NV. This is an interesting study for future astro- Let us turn our attention to the "nebular" nomical satellites. component. Emission lines and the strong Balmer continuum can be attributed to a nebula A more plausible model for the low-velocity excited by the hot star radiation (Boyarchuk, wind observed in NV is ejection from the cool 1966b). Such a simple nebular model, how- giant surface. The observed wind velocity of ever, fails to explain the large ionization range about 170 km s-_ is close to the stellar escape observed in AG Dra, with neutral to several velocity, but quite large with respect to the times ionized species. Other "ingredients" wind velocities generally observed in normal need to be added to explain the many peculiari- cool giants. The high ionization should be the ties. The HelI 1640 A line has broad wings result of the presence of an extended solar-type which, according to Viotti et al. (1983), seem transition region or, more probably, of ioniza- not to be produced by Thomson scattering. An- tion from the hot star radiation. Like the model other possibility is that the wings are formed in proposed for Z And, the near-UV continuum an accretion disk, or in a high velocity (103 km and the emission lines are probably mostly s-t) warm wind. To verify these possibilities, emitted from an extended region of the cool high S/N observations of the emission line star, and their emission is modulated by the profiles are required, which are not available orbital motion of the system as a result of vari- with IUE. The P Cygni profile of NV, which is ation of the visibility of the region (e.g., Viotti present during all the (orbital and activity) et al., 1984a). The variability during quies- phases of AG Dra, and the simultaneous ab- cence can also be explained if the is ellip- sence of such a profile in CIV, suggests the tical. In this regard Iijima (1987) suggested that existence of a low-velocity (-170 km s _) warm the photometric variations are due to variation wind in the system (Viotti et al., 1984a). This of the mass-transfer rate during the orbital velocity is too low to be associated with a stel- cycle. However, the results of Garcia (1986) lar wind from the hot star, which should have a and Garcia and Kenyon (1988) are in better much larger velocity, unless the structure of its agreement with a low eccentricity of the sys- atmosphere is very peculiar because of the tem. accretion processes. As discussed by Viotti et al., a dense, "torrid" wind (with a temperature The light history of AG Dra was character- higher than 105 K), for instance, could be pro- ized by recurrent phases of activity and long duced from the polar regions, or near hot spots. periods of quiescence. By combining the pho- Since all nitrogen is in the form of at least NIV, tographic magnitude estimates of AG Dra since no NV lines are formed there. At a certain dis- 1920, lijima et al. (1987) suggested a recur- tance from the star's surface, the wind slows rence period of the outbursts of roughly 15 down. This fact causes a decrease of the density years, corresponding to about ten 554 d cycles. with the radial distance r slower than r2; that, Iijima et al. suggest that the outbursts are pro-

680 ducedby mildhydrogenflasheson amassive Unlike the other S-type symbiotics, AG Dra (-1.2M° ) whitedwarfundergoinglargemass is a strong X-ray source. The X-ray spectrum is accretion(-10 7M: yr'). However,suchan very soft, with an integrated luminosity of 2.1 accretionrateis notconceivablefor AG Dra, x 10 '2 erg cm 2 s _ in the 0.2-1.0 keV range since the cool star seems not to fill its Roche (Anderson et al., 1981). If the X-rays are pro- lobe. IUE observations have shown that during duced by the cool giant, the ratio of the X-ray the 1980 outburst, the far-UV continuum in- flux to the bolometric flux (in the usual units creased significantly. Viotti et al. (1984a) for the HEAO-2 observations) is equal to 2.1 x found that the variation occurred at nearly 104 which is three orders of magnitude larger constant temperature, implying that the effec- than that observed in the Hyades K giants tive stellar radius should have increased by a (Stem et al., 1981a). AG Dra should have an ex- factor of two to three during outburst. Kenyon ceptionally enhanced chromospheric activity and Webbink (1984) considered that the 1980 giving origin to an extended corona, for in- outburst was thermonuclear, but substantially stance, as observed in some dwarf stars (Stem less developed than the large-scale events ob- et al. 1981b.). From the analysis of the HEAO- served in other symbiotic stars, such as the 2 data, Anderson et al. (1982) obtained a symbiotic novae. The white dwarf would not plasma temperature of 1.1 x 106 K, and an have developed a very extended envelope, so emission measure of 2.6 x 1055 cm _, much that its effective temperature would have re- lower than that derived from the Hell 1640 A mained high. It should be important to find line, but close to the values for the intercombi- possible probes of the hot star structure close to nation lines (e.g., Viotti et al., 1983). Alterna- the time of the outburst. Figure 13-9 shows the tively, X-rays are lormed in a hot "area" emit- variation of the ratio of the flux of the Hell ting as a with a temperature of 1.5 1640 A emission line and of the far-UV contin- x 105 K, and a radius of 1.4 x 103 km (Anderson uum at 1340 A before and after the 1980 out- et al., 1982). This area can be identified with an burst (Viotti et al. 1984a). This ratio, as dis- active region on the K-giant surface. Actually, cussed in the previous sections, is a measure of Oliversen and Anderson (1982b) proposed for the far-UV temperature of the star, if one as- AG Dra a model of a (single) star with active sumes that the Hell line is radiatively excited regions of enhanced surface brightness to ex- and that the 1340 A continuum belongs to the plain the modulation of the U-light curve. Al- hot star. The figure shows that the ratio was the though we cannot exclude that the cool compo- same just before and after the outburst (the nent in symbiotic system has some kind of observations were made in October 23 and enhanced activity (and this point needs to be November 15, respectively, when, according to further investigated), it is difficult to accept Viotti et al., the visual magnitude changed by that this activity regards a large fraction of the -0.7 mag, and the UV fluxes by -1.6 mag). The stellar energy output. HelI/Fc(1340A) ratio largely decreased in the period following the first light rise, and Garcia (1986) considered the X-ray lumi- reached again the preoutburst value in 1984. nosity as converted gravitational energy due to The 1982 minimum can be explained by a capture of matter from the cool star's wind. But lower color temperature of the hot star (about the required mass-loss rate from the K giant 80,000 K), after the outburst, but it remains turned out to be too high for a normal giant difficult to explain the long delay of the (10 7 M ° yr'). The X-rays are most easily ex- change. It should be considered that before the plained as the high-energy tail of the hot com- outburst, there was a slight increase of the Hell/ ponent spectrum. Slovak et al. (1987) found Fc(1340A) ratio, which could be an indication that HEAO-2 and far-UV IUE data (during the of a heating of the stellar surface before the quiescent phase of mid-April 1980) can be fit- event, which can be associated with the onset ted by a black-body with a temperature of of the thermonuclear outburst. 191,000 K and a luminosity of 174 L o. AG Dra

681 U-PHASE 1.0 2.0 3.0 ' I I I ' I

He II 1640. (,_) Fc (1340) 3OO OUTBURST

100 .._._ j f j/

3O

I , I I I , I 1980.0 1981.0 1982.0 1983.0 1984.0

YEAR

Figure 13-9. The Hell 1640A/Fc (1340A) flux ratio in AG Dra before and after the 1980 outburst. was observed again with EXOSAT in June of 1985 and 1986, without a similar decrease 1985 by Piro et al. (1985), who found that the of the ionization of the emission line spectrum observations can be fitted with a (see Figure 11-34). The cause of this rather Bremsstrahlung model with log N(H) = 20.2, unexpected behavior is not clear, but seems to kT = 24 keV and a flux of 3.4 x 10 _3 erg cm2 suggest that X-rays do not represent the tail of s _ in the 0.2-1.0 keV range. (It should be re- the hot star spectrum. They are probably called that all these flux estimates strongly produced in nearby region. If the region is depend on the adopted interstellar extinction. heated by the stellar radiation, its angular For instance, Anderson et al. (1982) assumed extension should be small, so as to capture only log N(H) = 20.5, while Slovak et al. (1987) a small fraction of the stellar EUV photons, in have made no reddening corrections. Piro et al. agreement with the very large Hell/X-ray emis- used the N(H) value derived by Viotti et al. sion measure ratio. (1983) from the interstellar Lyot). If the X-rays are produced near the hot component, and if the orbit of the AG Dra system is seen nearly edge- To summarize, the presently available ob- on, we should expect an eclipse at phase 0.5 of servational material on AG Dra is best inter- the Meinunger's light curve. In fact, AG Dra preted in the framework of a detached binary was observed again with EXOSAT in Novem- model, formed by a rather normal K-type ber 1985, at the time of the expected eclipse, giant, and a degenerate star that is "heated" by but the X-ray flux was the same (Cassatella et slow mass accretion from the giant's wind. The al., 1987). Cassatella et al. also found a large accretion is probably responsible for the recur- decrease of the flux during the small outbursts rent outburst. The K-giant could be peculiar for

682 having an anomalous chemical abundance, in At present, AG Peg displays a typical symbi- agreement with its Pop lI nature, and a rather otic spectrum with a cool (red) continuum and high-velocity wind. But it is not clear whether strong TiO absorption bands. The emission these facts are associated with the symbiotic lines of low and high ionization are very promi- phenomenon, it is also not yet clarified if there nent in the visible and UV spectrum, with a hot is any enhanced surface activity of the giant. continuum extending to the far-UV (e.g.+ More detailed models of the AG Dra system Boyarchuck, 1967; Hutchings et al., 1975). AG require a careful analysis of the available and Peg was the first symbiotic star observed in the future data, especially the time variability in ultraviolet with OAO-2 (Gallagher et al., different frequencies. Of particular importance 1979). The star was extensively studied with would be the systematic study of the cool spec- the IUE satellite (e.g., Keyes and Plavec, 1980: trum at high resolution, to derive the basic data Penston and Allen, 1985). Some high-resolu- (abundance, surface gravity, turbulence, rota- tion profiles of UV features are shown in Figure tion, etc.), but also to improve the orbital para- l l.31c. In many respects, the emission line meters of the system, and high-quality emis- spectrum is similar to that of a WN6 star, and sion line profiles collected during the orbital AG Peg is often classified as M+WR, although cycles and at different activity phases. the luminosity of the hot spectral component actually is lower than that of normal WR stars. IV. THE SYMBIOTIC NOVAE The WR features are more probably associated with interactive phenomena in a binary system, 1V.A. INTRODUCTION as discussed later. It should be noted that with- out the knowledge of the light curve of AG Peg so many years ago, its nova-like nature would In describing the light curves of symbiotic not have been recognized on the basis of its stars in Section 11.III, we have shown that there exists a small group of objects character- 6 -- V AG Peg ized by the fact that they have undergone one _ l "_" 8 ! _ single outburst in their known light history. # I This group (or subgroup) was first identified by I .... I I I I I Allen (1980b), who called them symbiotic no- 1850 75 1900 25 50 6 _ v ,_ RR Tel - vae. Attention was directed onto these stars after the recent outburst of a few northern ob- 10 jects, namely VI016 Cyg, V1329 Cyg, and HM Sge, whose behavior was found to be similar to 14 those of slow novae. These objects were exten- I I 1 l I I sively studied in all the wavelength ranges for 40 1945 50 55 60 65 70 lO - V V1016 Cyg several years after their outburst. Thus our present knowledge of their behavior is rather 12 complete. Other objects have displayed, in the # past, a similar behavior, the most remarkable 14 ones are RR Tel, which will be discussed in the I /1 I 1 I I I next section, and AG Peg. AG Peg is the oldest 60 1965 70 75 80 85 90 12 known . In the middle of the B _. V 1329 Cy_l past century, the star underwent a major nova- f 14 like outburst, with a very slow increase of the _t// visual luminosity from the 9th to the 6th mag- 16 I I I I 1 I 1 nitude in one to two decades, and a still longer 60 1965 70 15 140 85 90 decline to the present magnitude, which is Figure 13-10. "/+he com]_tll'tltil't' [iL_]l[ cl1/'x'£,_ Qf S(_tltl' close to the reported preoutburst luminosity symbiotic novae: AG Peg. RR Tel. VlOlO Cyg, and (Figure 13-10). V1329 Cyg ( = HBI 475).

683 presentbehavioralone.Thestaris, in fact, (1988b, 1989). quite different from the other "classical" sym- biotic novae, such as RR Tel, Vi016 Cyg, and HM Sge, all of which are D-type without the prominent M-spectrum of AG Peg. It is quite IV.B. RR TELESCOPII conceivable that several other symbiotic ob- jects actually belong to the category of symbi- otic novae, because they underwent in the past RR Tel was discovered as variable by Mrs. a nova-like outburst. But, owing to the long Fleming (1908) many decades before its main time scale involved, and the low frequency of outburst. The light curve, based on 600 Har- the phenomenon, their main outburst has not vard observations from 1889 to 1947 was de- been recorded. scribed by Mayall (1949), and is schematically shown in Figure 13-10. According to Mrs. Allen (1980b) listed seven stars having the Mayall, RR Tel showed little evidence of peri- character of very slow novae: AG Peg, RT Ser, odic variations from 1889 to 1930, the ob- RR Tel, V1016 Cyg (MHet328-116), VI329 served range being about i.5 mag with maxima Cyg, (HBV 475), HM Sge and V2110 Oph (AS ranging from 12.5 to 14 mag. After 1930, the 239). More recently, a new event--the outburst periodicity of the variation became clearer, and of PU Vul--was recorded, and the star added to a mean period of about 387 days could be de- this small class of objects. Figure 13-10 shows rived with an amplitude of about 3 magnitudes. the schematic light curves of some symbiotic This behavior is typical of a long-period vari- novae. The basic parameters are summarized in able, and is presently barely visible at optical Table 13-2 (from Viotti, 1988b). In the follow- wavelengths with a period of about 374 d (Heck ing, we shall discuss in detail the case of RR and Manfroid, 1982; Kenyon and Bateson, Tel, which, for its luminosity and spectral 1984). The period, however, seems to be vari- evolution, can be considered as the best repre- able between 350 and 410 d (Heck and sentative of the category of symbiotic novae. Manfroid, 1985). As already discussed in Sec- The main observational properties of most of tion I I.F, the Mira-type pulsation is evident at these object were already presented in the dif- IR frequencies (Feast et al., 1983a). In late ferent sections of Chapter I1. The general 1944, the periodicity stopped and the star rap- properties of symbiotic novae, for instance, idly brightened from m = 14 to 10 in a few Pg were discussed by Kenyon (1986a), and Viotti days, then rose to 7 mag by mid-1945. In the

TABLE 13-2. THE SYMBIOTIC NOVAE

Star To _ Tmax 2 Magnitude Spectrum Type Pre3 max 4 post s cool max IR

AG Peg 1855 1871 9 6 8.3 M3 S RT Ser 1909: 1923: >16 9.5 13 M5.5 A8 S V2110 Oph ...... 1940: 11: 22 >M3 D RR Tel 1944 1948 14v 6 !1 M5 F5 D V1016 Cyg 1964 1967 14 11 11 >M4 neb D Vi329 Cyg 1966 1967 14v 11.5 13-14 >M4 neb S HM Sge 1975 1975 >17 11 11 >M4 neb D PU Vul 1978 1982-83 15v 8.8 8.8 M4 A7 S

Notes to the table. (1) of beginning of the outburst. (2) Year of maximum luminosity. (3) Preoutburst mag- nitude. (4) Maximum luminosity. (5) Present (1986-88) postoutburst magnitude.

684 followingyears,RRTelremainedatmaximum peculiar variable Eta Car (Thackeray, 1953). luminosityuntil 1949,reachingthesixthmag- According to the Bosque Alegre spectra shown nitudeduring1948.Thenthe stargradually in Figure 13-11, the spectral variations should fadedto thepresentV = 10in about14years have taken place in less than one week, and (seeKenyon,1986).Althoughthestarwas maybe in a few days. quitebrightat maximum,it wasdiscovered only in 1949,three years after the outburst. The spectral evolution of RR Tel in the fol- Therefore no spectroscopic information is lowing years is best described by A. D. Thack- available on the early development of RR Tel. eray's review (Thackeray, 1977). Since 1949, Nevertheless, since late 1949 the star under- the star has shown a gradual increase of the went a major spectral evolution which was mean ionization of the emission line spectrum. followed at the Bosque Alegre (Argentina) and Hell and NIII appeared around August 1950 Radcliffe (South Africa) observatories. Figure (see also Pottasch and Varsavsky, 1960). These 13-11 illustrates the spectral evolution of RR authors also identified Hel absorption with a Tel between April 1949 to July 1971 based on velocity of -685 km s _ (in 1951) and -865 km a collection of spectra obtained at Bosque s_ (in 1952). Then, between 1951 and 1952, Alegre. [OIll] and [Nelll] flared while the permitted Fell lines faded (Thackeray, 1953). The in- crease of the level of ionization continued through 1953 and 1954 and is best illustrated by The first spectrum taken at Bosque Alegre in the sequential appearance of higher and higher April 1949 showed a strong continuum with ionization stages of iron, from FeIIl to FeVII. many absorption lines of singly ionized metals. This sequence is shown in Figure 13-12. First, Some absorptions are flanked at longer wave- the spectrum only showed the low-ionization lengths by a weak emission component, indi- lines of permitted Fell. Then, the forbidden cating a marginal P Cygni profile. This emis- [FeII] lines appeared and gradually strength- sion became more prominent in July 1949. By ened with respect to FeIl, indicating a decrease mid-September, the continuum appeared weaker and the emission lines dominated the of the density of the emitting medium (cf. Viotti 1976). spectrum of RR Tel, although the mean line excitation was still low. The first spectra taken at the Radcliffe Observatory of South Africa The sequence continued with [Felll], whose were discussed by Thackeray (1950), who lines brightened in 1952; [FelV] (1952-53); found strong absorption of Call and hydrogen [FeV] (1954-56); [FeVI] (1956-59); and finally in June-August 1949. TilI was present in ab- [FeVI1] (1959-64). This behavior is similar to sorption, and HI3 was absent, probably filled in that found in novae after the optical maximum, by emission. Mayall (1949) reports on low- but with the important differences that in RR quality, low-dispersion spectra taken when RR Tel (and in other symbiotic novae), the trend Tel was at maximum brightness. All these ear- was much slower. In any case, RR Tel was the lier ob,;ervations agree in giving an F-super- first nova-like object in which this phenome- giant spectra[ type (cF5, according to Thack- non was studied in such detail. It should also be eray, 1950). In a later paper, Thackeray (1977) noted that at the time of Thackeray's observa- reports that the relative shift of the absorption tions, the spectrum of three to five times ion- lines was - 100 km s _. The remarkable spectral ized metals was very poorly known. His careful change, which occurred between August and study of the spectral evolution of RR Tel, and September 1949, was first noted by Thackeray the theoretical computations of B. Edlen and (1950), where reported that in his spectra all R.H. Garstang led to the first identification of the absorption lines disappeared and a rich many forbidden lines of FelV to FeVI, and of emission - line spectrum appeared with promi- other highly ionized metals. From this point of nent hydrogen, CalI and especially Fell emis- view, RR Tel has also been a good laboratory sion, close resembling the spectrum of the for atomic spectroscopy.

685 I 19/.9- iV- 6 Vii - 9 I0 15 16 17 Viii- 16

26

IX- 5 6 II X - 2.0 Xl - 11 195o- 111- 9 V -'/ I0 Vi - 6 VII- 5 VIII - /* Xl - 2o 1951- VI -Z6 Z9 IX - 2 1953- VII - 16 I95/*- Vl - Z_

VIII - 12 1955- IX - tZ 1957 ° Vl - 6 1956- IX - 11 1960- X - 13 196_- X - 16 1955- VII- 20 1971 -Vll - 3

19"/I- Vll- 5

Figure 13-11. (Plate) the blue spectrum of the symbiotic nova RR Tel between April 1949and July 1971. This is a reproduction of spectrograms taken at the 150-cm telescope of Bosque Alegre Observatory (Argentina). The original reciprocal dispersion is 41 A mm i. Note the dramatic spectral change occurred in September 1949. Courtesy of Professor Jorge Sahade.

686 I ! I I I I I I I I I

I

20 I Fe III t t t t I I I I I t ! 40 M

20 _T Fe IV i

120 I

I

60 I

,' I I I I I I I I I I

240 I

I Fe VI -

120 m

2OO I

100

I 1 I I 1 I 1 I 1 1952 1960 1968

Figure 13-12. The evolution ¢_the intensity qf tire lira's of the different hmizatiml stages of iron in RR Tel during 1952 to 1964 as described by Thackerary (1977). Note the gradual increase of tire ionization qf tire emitting envelope as indicated by the appearance of the evolution of the successive ionization stages of iron (Viotti. 1987).

In the ultraviolet, the star presents a very high-temperature features at 683(I and 7088 A). rich with a wide range of Penston et al. (1983) found that the width of the ionization, from neutral species to five times emission lines varies between different atomic ionized calcium (Penston et al., 1983). More species and increases from 40 to about 80km recently, Raassen (1985) suggested the identi- s _ from low- to high-ionization lines. Similar fication of a line at 2648.9 A with the 3P 2 - 1D 2 correlation between line width and ionization transition of [FeXI], which could thus be the energy was previously reported by Friedjung highest ionization stage so far observed in RR (1966) and Thackeray (1977) for the optical Tel (possibly excluding the yet unidentified lines. This behavior was also found in other

687 symbiotic novae, such as VI016 Cyg and HM deep 1980 minimum of PU Vul, and in AG Peg Sge. and RT Ser after decline from maximum. In V1016 Cyg and HM Sge, which have not (yet) declined, the M spectrum in the visible is masked by the strong continuum and line emis- sion from the circumstellar regions. In these stars, the presence of a late-type component is IV.C. GENERAL PROPERTIES OF THE supported by the Mira-type variations in the SYMBIOTIC NOVAE near-IR. As in RR Tel, the Mira star could be hidden by a dense circumstellar dust shell Like RR Tel, the other symbiotic novae also (Kenyon et ai., 1986). Indeed, in the D-type have the common property of having under- symbiotic novae, the Mira component should gone a single major outburst and of showing a be subject to large mass overflow, followed by symbiotic spectrum. Since there exists in the formation of dense gas and dust clouds. Table few objects classified as symbiotic novae a 13-3 (from Viotti, 1988b) summarizes the typi- large variety of behavior, it is important to cal time scales of the cool components of investigate their general properties and symbiotic novae. The spectral types found whether they represent an extreme case of from the literature are given in Table 13-2. The symbiotic stars or have to be associated with luminosity class is II1 for AG Peg and PU Vul, the category of novae. In the following, we and for the Mira components of the D-type summarize the different aspects of the phe- objects as well. nomenon.

a. The preoutburst phase and the red compo- b. The outburst. nent. As seen in Figure 13-10 and in Table 13.3, The preoutburst phase is known in some de- the rise to maximum was fast in V1329 Cyg, tail only for V1329 Cyg and RR Tel. Both HM Sge, and RR Tel, and very slow in V1016 appeared largely variable, on a long time scale. Cyg and especially in AG Peg and RT Ser. It is In VI329 Cyg, the variability is interpreted as noticeable that the rise time is apparently not due to eclipses of a binary system, and this is related to the other features (e.g., the IR-type) also supported by recent radial velocity meas- of the symbiotic novae. The amplitude of the urements (e.g., Grygar et al., 1979; Nuss- outburst, ranges from 3 mag (AG Peg) to more baumer et al., 1986). On the contrary, in the than 6 mag (HM Sge, RT Ser). It is also impor- cases of RR Tel, the long-term variability is tant to consider that this difference is not due attributed to a Mira-type pulsation, as also (or not only due) to the actual amplitude of the confirmed by the recent optical and IR moni- outburst, but rather to the relative brightness of toring (Heck and Manfroid, 1982; Feast et al., the late-type component (with respect to the 1983a). The orbital period of RR Tel, as in luminosity of the symbiotic nova at maxi- other D-type symbiotics, is believed to be mum), which is high in AG Peg and V1329 much longer. Cyg, and very low in RT Ser, V2110 Oph, and In these two objects and in V1016 Cyg, the HM Sge. The actual visual luminosity increase preoutburst spectrum was M-type. It is quite of the red-giant companion is unknown. possible that the luminosity (and spectrum) of the preoutburst M star was the same as that of The spectrum at maximum is another in- the present red component of the symbiotic triguing problem. As in classical novae, some systems. This obviously implies that the out- objects (RT Ser, RR Tel, and PU Vul) dis- burst was not (at least directly) caused by the played an intermediate (A-F) equivalent spec- red star, but rather by its companion. An M- tral type, possibly of supergiant class, but with- giant spectrum was also observed during the out the highly violet-displaced absorption lines

688 TABLE13-3.CHARACTERISTICSTIMESCALESOFSYMBIOTICNOVAE.

Star Rise Decay Mira Orbit K (a) (b) (c) (d) (e)

AG Peg 16 40 == 816.5 5.1 RTSer 14 7: ...... RRTel <0.3 9 374.2 .... VI016Cyg 2-3 >125 472 .... V1329Cyg 0.3 12-20 == 950 62 HMSge <0.4 >65 500-600 .... PUVul 1 >38 ......

Notes to the table. (a) Rise time of the visual luminosity in years. (b) The e-folding decay time _" tile visual luminosity in years. (e) Period (in days) of the Mira pulsation. (d) Orbital period (in days). (e) Semiamplitude of the radial velocity curve (in km sJ). that are seen in novae. No similar absorption expansion velocity of the main emitting region. spectra have been observed in the other symbi- There are, however, a number of interesting otic novae. The simplest explanation is that the exceptions. Crampton et al. (1970) found in absorption-spectrum phase occurred during a V1329 Cyg the [OlIl] and [NeIll] lines having period not covered by the observations, and we a multiple structure, with emission peaks rang- missed it. In fact, VI016 Cyg was first ob- ing from -240 to +250 km s t. This line multi- served spectroscopically near the end of its plicity was also present in later spectra of the long-lasting rise to maximum (indicated by an star (Grygar et al., 1979; Tamura, 1988) (Fig- arrow in Figure 13-10). The first spectra of ure 13-13; see also Figure 11-13). V1329 Cyg were taken at the end of 1969, one year after its light maximum, and showed a rich As discovered by Crampton et al. (1970) and emission line spectrum, which could be com- confirmed by Baratta et al. (1974), the low- pared with the post-maximum spectrum of resolution spectra of the star show several classical novae. In this regard, RR Tel repre- broad and shallow emission features, which sents a rather fortunate case, since its A-type were identified with WN5-type lines having an spectrum suddenly disappeared in late 1949, expansion velocity of about 2300 km s _ (figure just a few months after the discovery that the 11-10). As discussed above, broad WR struc- star had "exploded". Suppose that the discov- tures have also been seen in AG Peg and RR ery would have been made 4-5 months later, in Tel, while high-velocity P Cygni profiles were late September or in October 1949, instead of observed in RR Tel during decline. Therefore, in April. We would have missed the absorp- in some symbiotic novae, at least, there is evi- tion-spectrum phase, and perhaps have placed dence for the presence of high temperature, RR Tel in a different subcategory of symbiotic high-expansion velocity regions, but this novae (see, e.g., Kenyon and Truran, 1983). should only represent a small portion of their Thus it is important to take into account these emitting envelope. This result is probably re- selection effects in any modeling of the phe- lated to the velocity gradient found in some nomenon. objects from the analysis of the line width of the narrow emissions, with the higher energy Like the classical novae, the symbiotic ions (and the corresponding higher tempera- novae after the outburst display a rich emission ture-emitting regions) having larger expansion line spectrum with wide ionization energy velocity, as suggested by the observed correla- range. But unlike novae, the emission line tion between emission line width and ioniza- profiles are narrow, in general, indicating a low tion energy (see Chapter 11 Section VIII.D.) It

689 i ! I ! I ! i i ! by the radiation of a hot source (a hot star or the INT. hottest parts of an accretion disk) in the system. Collisional ionization is another possible mechanism to explain the very high ionization 15 features and the X-rays. It might occur in shocks lbrmed by the interaction of the winds of the two stellar components, tor instance, as suggested by Willson et al. (1984), or by colli- sion of the stellar wind(s) with the circumstel- lar matter. In general, the analysis of the high- ionization emission line fluxes and of the far- 5 m UV continuum leads to a model of a hot central source with temperatures (around 105K) and radii (0.1 R ° or less) typical of the nuclei of

L Planetary Nebulae (e.g. Nussbaumer and 0 KM/SEC. +5OO -5_) Schild, 1981; Tamura 1981; Mueller and Nuss- baumer, 1985; Hayes and Nussbaumer, 1986).

c. The decline phase.

JULY 1970 The behavior of symbiotic novae after the outburst is very different from case to case. Four objects showed a gradual fading of the visual luminosity that took several years to decades (Figure 13-10). The e-folding decline time varied from 7-9 years for RT Ser and RR Tel to 12-20 years for V1329 Cyg and 40 years "_lfa V L _ 01 I - 400 km.s for AG Peg (Table 13-3). In the case of VI329 400 Cyg, the decline time was derived from a fit of the UV emission line flux variation, taking into Figure 13-13. Emission line profiles of the forbidden account the 950 d periodicity (Nussbaumer et lines in the spectrunz of the symbiotic nova W329 al., 1986). We recall that Allen (1981) and Cyg. Top." Mean profile of [O 111]and [Ne 111]lines in October 1969 (Crampton et aL, 1970). Bottom." Willson et al. (1984), from the analysis of the profile of the/0111/4363 line in July 1979 (Gr),gar X-ray flux in three symbiotic novae, VI016 etal., 1979). Cyg, HM Sge, and RR Tel, suggested a very slow decrease of the X-ray flux after the out- is also conceivable that the WR features ob- burst, with an e-folding decay time of 5 to 50 served, especially in VI329 Cyg, originate in a years. But this result needs to be confirmed. dense expanding envelope or wind of the hot star, which behaves like some WR-type nuclei Three recent symbiotic novae, V1016 Cyg, of Planetary Nebulae. In view of possible HM Sge, and PU Vul, have not significantly models, it would be interesting to know if these faded since their outburst, although a small WR features are present far away in time from visual luminosity decrease has been recently the outburst, or in the spectroscopic records of noted for PU Vul (Gershberg and Shakhovskoj, the preoutburst phase. 1988). Their decay time is probably similar to that of AG Peg, or even much larger. Once The emission line spectrum should be again we recollect that in symbiotic stars, formed in an extended circumstellar or circum- emission lines largely contribute to the broad- system nebula, which is supposed to be ionized band photometry, so that the observed light

690 curve of an object does not necessarily describe decline discussed above. During this phase, the its global time behavior, but should also reflect visual and ultraviolet spectrum remained local fluctuations of the physical structure of dominated by the hot component, but with a the emitting envelope. Therefore, it is possible gradual evolution from F5 to A2 during 1983- that some precious information remains 1986 (Gershberg and Shakhovskoj, 1988). In masked in the broadband light history. late 1987, Maitzen et al. (1987) noticed an increase of the emission line strength. The As in RR Tel and in classical novae, the 1980 minimum was also followed by Friedjung spectral evolution of VI016 Cyg and HM Sge el al. (1984) beyond the visual range. As shown was characterized by the gradual increase of in Figure 13-14, the amplitude of the minimum the ionization level of the emission line spec- was much larger at shorter wavelengths, and trum. WR features were first detected in HM just detectable in the near infrared. Also the Sge two ),ears after the outburst (Ciatti et al., duration of the eclipse was longer in the UV. 1978). As these faded, HeII and [FeVII] Friedjung et al. interpreted the deep minimum emerged with very intense lines (Blair et al., as a result of temporary obscuration of the 1981). We recall that, unlike in the novae, the "exploded" star by dust condensated from the spectral evolution of Vl016 Cyg and HM Sge ejected shell. Dust might also have been pro- occurred at nearly constant visual luminosity. duced by the red giant. Alternatively, the hotter star has been eclipsed by the M giant (Kenyon, Symbiotic novae were also followed at radio 1986b). From the duration of the minimum, wavelengths, sometimes over a period of sev- Kenyon derived an orbital period of about 700 eral years (see Section I I.Vll). Radio observa- years. In both hypotheses, the M spectrum tions of HM Sge started in 1977, two years after observed at minimum should be that of the coo[ the outburst, and disclosed a gradual evolution, giant component of the system, which at maxi- with a steady increase of the radio flux at 15 mum is completely masked by the radiation of GHz from 40 mJy in 1977 to 150 mJy in 1985. the early type component. The radio spectrum remained optically thick all the time, indicating that, unlike classical IV.D. POSSIBLE MODELS FOR SYMBI- OTIC NOVAE novae, the expanding HII region was still dense several years alter the outburst (Kwok et al.,

1981, Kwok, 1988). An optically thick radio Let us now examine the above "main proper- spectrum is also displayed by VI016 Cyg. The ties" of symbiotic novae in the light of possible star was first observed in 1973, nine years after models. Other aspects of the problem are dis- the outburst, but no significant flux change has cussed in Viotti (1989). Although the preout- been detected since then. Probably it reached a burst phase is very poorly known, it is clear that stationary stage before 1973. We finally recall the large increase of the visual luminosity, the large and probably irregular radio variabil- mostly due to the appearance and strengthening ity of the other symbiotic nova V1329 Cyg. of emission lines, was associated with a sudden increase of the flux of far-UV photons from the One particular case is PU Vul. This star, M-giant's companion. The hot source should after the main brightening phase that took have largely increased its brightness tempera- about 1 year, remained at maximum for another ture and bolometric luminosity with respect to year, showing an F-supergiant spectrum (e.g., the previous unknown stage. As discussed in Nakagiri and Yamashita, 1982; Kolotilov, Section 12.iV.A, such an event can be ex- 1983). Then, during the first half of 1980, the plained as a result of sudden thermonuclear visual luminosity gradually dropped from V = burning of the hydrogen-rich matter accreted 8.8 to 13.5 (Figure 13-14), and the M-type by a degenerate star from the red giant wind. spectrum appeared. PU Vul remained at mini- According to the models developed, among mum for about 200 days, then gradually flared others, by Paczynski and Rudak (1980), Fujim- up again to V = 8.5, followed by the very slow oto (1982a,b), Kenyon and Truran (1983),

691 ' I ' ' I i

n PU Vul

6-- K L l

8 -

8 --

10

12

14 "T-

0

i a

+ 0.8 B-V _ _ _ __'- i

0 i

m

+0.8 m

-13 m

-14 m

-15 i

! t I i i I i , I = 1977 1980 1983 1986

Figure 13-14. The infrared, optical, and ultraviolet light curve of PU Vul during the recent outburst (Friedjung et al., 1984; and Kenyon, 1986, adapted).

Kenyon (1988), and Livio et ai. (1989), the first case is supposed to occur in at least some accretion rate should be smaller than that re- Z And-type symbiotic stars, symbiotic novae quired to have a stable burning of the matter as should represent the dividing line between Z it is accreted, and larger than that which char- And-type variables and classical novae. Due to acterizes the classical nova outburst. Since the the larger separation of the components in

692 symbiotic novae, the accretion from the red- finally consider that during the high-tempera- giant wind (the systems are clearly detached) ture phase, the exploded star should have a should be smaller than in other symbiotic sys- dense hot wind, which might have produced the tems, even if we expect denser winds in those WR features observed in several cases. This symbiotic novae containing a Mira. point, however, has not yet been investigated in detail, especially in order to find possible dif- According to Kenyon (1988), thermonu- ferences with the hot components of symbiotic clear flash models imply that, it" the accretion systems whose high surface temperature is not rate is fairly low, the white dwarf envelope is the result of accretion processes. In particular, completely degenerate. Under these condi- the chemical composition of the wind should tions, the luminosity of the star first increases reflect the recent violent history of the star, and at nearly constant radius, then a slow expansion this problem should require more studies. at constant bolometric luminosity follows until an A-F supergiant configuration is reached (see The outburst of symbiotic novae might be Figure 12-4). Later, after a rather long time, the explained by instabilities of an accretion disk star evolves again to high effective tempera- (of., Duschl, 1986b), but this model should tures. This de+enerate .[lash model possibly require high accretion rates, which do not ap- applies to AG Peg, RT Ser, RR Tel, and espe- pear to be realistic for detached systems. Alter- cially PU Vul. However, in the case of RR Tel, nately+ the outburst can be the result of a sud- it is difficult to explain the rapid evolution of den onset of a strong stellar wind from the cool its spectrum from F-supergiant to emission giant (Nussbaumer and Vogel, 1988). The wind line. will produce an extended envelope surround-

If the accretion rate onto the white dwarf is ing the system. The luminosity increase is the result of the ionization of the envelope by the larger, the accretion results in a non-degener- UV radiation of the hot stellar component. This ate envelope, and produces a relatively weak model has to be worked out in more detail, with shell flash. These weak non-degenerate flashes special attention to the time scales involved in (Kenyon, 1988) do not evolve into the A-F the processes. Periodic enhancements of the supergiant stage discussed above, but the star accretion rate could occur if the orbit is highly remains hot throughout the eruption. This eccentric, and the red giant is going to fill its could be the case of V1016 Cyg, V1329 Cyg, Roche lobe at periastron. Such a model was and HM Sge, which have probably not devel- proposed by Kafatos and Michalitsianos (1982) oped the intermediate-type spectrum in the earlier stages of their outburst. In any event, to explain the outbursts of R Aqr (See the next section V.), but obviously it does not apply to such a phase, if it occurred during an early those symbiotic novae, AG Peg and VI329 unobserved phase, should have lasted quite a short time, one year or less, which is difficult to Cyg+ whose period appears too short for the time scale involved in the symbiotic nova explain in the light of the proposed models. The phenomenon. Other observational and theo- amplitude of the outburst in the visual should retical aspects of the phenomenon that need to be larger in the former case of a degenerate be further investigated are the identification of outburst, as a consequence of the small bolom- other symbiotic novae whose main outburst has etric correction at maximum. Although the not been observed, the recurrence of the phe- amplitude of some well-documentated objects nomenon, and the structure of the circumstellar (RR Tel and PU Vul, on one side, VI016 Cyg nebula. But we especially need the basic para- and VI329 Cyg, on the other) apparently seems meters of the binary systems. to support this model, in the reality, the preout- burst visual magnitude in all these objects is V. R AQUARII: A SYMBIOTIC MIRA that of the red giant, since the preoutburst spec- WITH JET trum is M. The actual amplitude of the white dwarf outburst should be larger, and probably much larger, than that given in Table 13-2. We R Aqr is one of the most peculiar astrophysi-

693 cal objects,sincethecharacteristicsof many cases the variability nearly disappeared. For differentastrophysicalcategoriesarepresent instance, this has happened in the years 1905- in thesameobject(Michalitsianos,1984).R 10, 1928-30, and 1974-78. Thus, the light curve Aqr is symbiotic for its composite spectrum presents a kind of a long time scale "modula- characterized in the visual by many emission tion" of the amplitude of oscillation. Willson et lines and a late-type Mill component. R Aqr is al. (1981) suggested that these irregularities also a Mira-type variable with a period of 387 should be caused by eclipses of a close binary days, but the light curve presents important system orbiting in a highly eccentric orbit with irregularities. A SiO maser emission was also a period of 44 years. detected. The star is interesting for being at the center of a with a mid-ioniza- R Aqr was monitored in the infrared (JHKL) tion nebular spectrum. The central part of the at SAAO during 1975 to 1981 (Catchpole et al., nebula, studied at radio wavelengths, is highly 1979; Whitelock et al. 1983b). These observa- variable with jet-like features. Finally, as dis- tions confirmed the visual periodicity of 387 cussed in Chapter II Section IX.E, R Aqr was days. The light curve in the L-band (about 3.6 recently detected as X-ray source with EX- lam) is slightly different from the visual curve, OSAT (Viotti et al., 1987). Thus, it is difficult with a steeper decline after maximum, and a to put R Aqr in one specific category. In addi- slower rise to maximum, which is reached tion, the star is rather different from the "clas- slightly later than in the visual. This is fairly sical" concept of symbiotic stars. However, we normal for a . Whitelock et al. may consider that the symbiotic phenomenon (1983b) noted that the infrared fluxes appeared is particularly evident in this object and that its depressed during 1975-78. They attributed this study could give an important contribution to to an obscuration by an opaque dust cloud as the problems that we are discussing in this suggested by Willson et al. (1981). i_nonograph. This is the reason for having de- voted a full section to this interesting object. The Mira character of R Aqr is also indi- Many aspects of R Aqr have also been dis- cated by the positive detection of SiO maser cussed by Querci (1986) in the previous vol- emission (Lepine et al., 1978), which is nor- ume on M-stars of this monograph series. mally associated with LPV's. So far, R Aqr is the only symbiotic star showing detectable maser emission (Lepine et al. 1978, Cohen and V.A. THE MIRA VARIABLE Ghigo, 1980). The negative detection of OH (Wilson and Barrett, 1972) and H20 lines (Dickinson, 1976) is probably related to the R Aqr, as indicated by the letter "R", was the inhibition by the hot close companion of the first variable discovered in the Aquarius con- Mira. More recently, Hollis et al. (1986) re- stellation. It was found as variable by Harding ported interferometer SiO observations indi- in early 1800, and since then it has been studied cating that the maser emission occurs in the by several astronomers. Thus, its light history nebulosity about one arcsec away from the has been fairly well known for almost two optical position of the Mira (see Figure 13-17). centuries. The light curve from 1887 to 1980 is This result is clearly in disagreement with a model reproduced in Figure 13-15. of collisionally pumped SiO emission (e.g., Elitzur, 1981 ). R Aqr is a red giant that shows large and quasi-regular light variations rather typical of a V.B. THE NEBULA Mira variable. The mean period is 387 days. The mean light curve generally presents a broad minimum lasting 6-7 months, followed The planetary nebula around R Aqr is essen- by a rapid rise to maximum. There are ample tially composed of two distinct structures: the variations from cycle to cycle, in both the outer nebula with an oval shape with is formed shape and amplitude of the light curve. In some by two arcs symmetrically extending to the

694 13000 14000

I I I | I I I I .,_

7_ ! I I I i i 1 i /_,A | 4 _. 9

11 ° °_ 1,5000 , , .... , o, ,1,_ , , , , .' ' , ' '_ I

9

11

17000 , ,

g

11 J 19000 ,

9

11

21000 ^ , , , , , , , , 22000 , , , , ,

' ..... , ......

; ,,,, ,_,,,_,,,,,,,, '1

Figure 13-15. The light curve of R Aqr from 1895 to 1941 Mattei, 1979). The visual magnitude is periodically variable between V = 6 and V = 11. with large variations from cycle to cw'le. Note. in particular, the am)malies during 1905-10, 1928-30, and 1974-78, which could be associated with enhanced activity of the hot component. The mean Mira period is 387 days.

695 East and West from the central star giving to 0.5" (Hollis et al, 1985; 1986). the nebula the aspect of a double lens (Figure 13-16). The radio jet is cospatial with the optical jet,

The R Aqr nebula has been recently studied and, because of the higher spatial resolution, by Solf and Ulrich (1983) who lbund that the can be studied with much more accuracy. Radio observations suggest an ordered geome- nebula is composed of two separate shells, try of ejecta: the distance of each knot, C2, A which are expanding at velocity of 30-50 km s-_. These shells should have been ejected from and B, from the central source CI is linearly dependent on position angle (Hollis el al, the central object 185 and 640 years ago. The 1986), and this should be associated with the spectrum of the nebula is typical of a (low mode of expulsion of the jets. According to excitation) planetary nebula. The problem is to Kafatos et al. (1986), components B, A, and C2 find the central ionizing source, which could be were formed during successive outbursts of the identified with the unobserved hot companion system, C2 being the most recent ejection, of the red giant. probably related to the mid-1970s event dis- A few years ago Wallerstein and Greenstein cussed above, while the two further ones (1980) first reported the detection in a 1977 should have been ejected long ago, during plate of R Aqr of a "spike" of emission nebulos- previous active phases of the object. At any ity that appeared as an elongation of the stellar rate, it should be considered that no expansion image towards North-East, never reported pre- of the radio knots has been so far detected (e.g., viously. Using Lick plates, Herbig (1980) and Hollis et al. 1985). Sopka et al. (1982) confirmed the jet-like fea- ture that was, however, not present in a 1970 High-resolution optical imagery of the R plate of R Aqr. Therefore, the jet should have Aqr complex should provide precious comple- appeared between 1970 and 1977. Sopka et al. mentary information on the nebula. Michal- also found the presence of an elongation in the itsianos et al. (1988b) have recently studied the radio map at 6 cm at the position of the optical large-scale structure of the nebula using a CCD jet. Later, higher spatial resolution radio obser- camera and narrow-band interference filters. vations obtained with the NRAO VLA of So- Paresce et al. (1988) used a coronograph in corro led to the identification of five separate conjunction with narrow band filters to imag- radio sources (Figure 13-17), the "jet" (source ine the immediate surroundings of R Aqr (I to B), a second jet closer to R Aqr, (A), a about 15 arcsec) at subarcsec spatial resolu- "counter-jet" (A'), while the central source C tion. These observations have put in evidence was resolved in two components separated by an S-shaped bipolar shape which comprises the

Figure 13-16. (Plate) The planetary nebula around the symbiotic-Mira R Aqr (Kafatos and Michalitsianos. 1984).

696 radiojet featuresdescribedabove.Theoptical radio jets. Again, this result has to be further imageis extendedin bothdirectionsandat investigated and compared with observations muchlargerdistancesthanobservableatradio with similar accuracy of other symbiotic and wavelengths.Observationshavealsorevealed Mira variables. thepresenceof severalknots,includingonenot seenat radiowavelengths.Thisbipolarsym- V.C. THE SYMBIOTIC SPECTRUM metryof theR Aqr innernebulasuggestsa The optical spectrum of R Aqr is rich in symmetriccollimatedflowfromR Aqr,associ- emission lines which are difficult to observe atedwitharotationorprecessionof thecentral when the star is near maximum light. The object.Foranyconsiderationof thiskind,the hydrogen lines and the nebular [OIIII and knowledgeof theprecisepositionof thestar- INelIII are strong in emission. As in other like counterpartis very important.Michal- symbiotics, the energy range is wide, as indi- itsianoset al. (1988b)derivedtheastrometric cated by the presence of low- (Fell, [EelI], etc.) positionof the Mira variablewithin about and high-ionization lines (Hell, NIlI, Clll). +0.05". The star position is about 0.15" SW of During the optical outbursts, the latter ones the central radio source C1 (Figure 13-17) and become stronger and broader. Zirin (1976) provides clues to the origin and ionization reported the identification of the coronal structure of the HII region surrounding the R IFeXIlll line at 10747 A in spectra made in Aqr system. As discussed above, the SiO maser 197(}-71. But there are no other observations of source is not coincident with the astrometric this line. position of the Mira variable, as one would have been expected, but it is placed 1" SE from The UV spectrum of R Aqr has been investi- C I and LPV, in the opposite direction of the gated since 1979 and has revealed the presence

iT I I I i

36"

38", o

' i II 40" _--

__ A I %, ,i / r 42"

435"

t

44" _--

A'

46" I 440" J

I I I 1 [ , "I '" 1 J

14s6 14s5 14s4 14_3 14'2 14 _ 1 14 s 0 14 s 32 14 s 30 14s.28 14s.26 14 _ 24 14 s 22

a(19500) = 23 _' 41 _

Figure 13-1 7. The high-resolution radio map of the central region of R Aqr, Left is the 6 cm map showing the

'jet-like" features A and B, and. marginally, the counteljet A. Right the central source is resolved into two

components_C l and C2. The astrometric position of the Mira variahle ( LPV) and o['the SiO maser emission is also indicated (fimn Michalitsianos el al., 1988).

697 of moderate-excitationemissionlines with doubletintensityratio I(1548)/1(1550)was prominentCIII] andCIV,andweakerO1,CII, foundequalto about0.5,i.e.,muchlowerthan SilV,OIV], andNIII] emissions.Theoverall the optically thick limit of unity. This far-UVspectrumis remarkablysimilarto that anomalousCIV doublet ratio intensity, effect of Mira itself asdescribedby Reimersand has been observed in other symbiotic stars at Cassatella(1985).Thehigh-ionizationlinesof least during some phases of their activity, such NVandHelIareweaklypresent.Kafatoset al. as in the case of CH Cyg (Marsi and Selvelli, (1986)foundthatthelineintensitiesfor the 1987; see Table 11-8). Michalitsianos et al. centralHII regionareratherstable,in spiteof (1988a) found that in RX Pup, the ratio I(1548)/ thelargeMiravariationsin thevisual.Onthe I(1550) is variable in time, and that it is in- contrary,the UV emissionlinesarelargely versely correlated with the CIV line intensity, variablein thejet A andB features.Thehigh- as well as with the visual luminosity. In Z And, ionizationlinesof NVandHellweregreatlyin- Cassatella et al. (1988a) noted that the CIV tensifiedin thejet in 1982andbecameeven doublet ratio was much smaller than one in strongerthanin thespectrumof thecentral February 1986, i.e., during the active phase source.Kafatoset al. (1986)notedthatthis started in September 1985, while it was slightly increaseof theionizationcouldbe relatedto larger than one during minimum. An anoma- thefirstdetectionof X-raysfromR Aqr(Viotti lous intensity for the NV resonance doublet et al., 1987).During1982-1986theemission was observed in the 1979 spectrum of AG Peg lineintensitiesvariedin a quasi-periodicway, (Figure 11-31c). The "anomalous" doublet ra- withminimain 1983and1985,andmaximain tio intensity cannot be explained by simple early1984andpossiblyin late1986(Kafatoset considerations on the line opacity. In some al, 1986,andunpublishedresults).Thisone- cases, it could be the result of intense high- and-halfyearmodulationis largerthanthe temperature interstellar lines, since in this Mirapulsationperiod,butcouldberelatedto case, the stronger emission component of the it. In fact,if onetakesintoaccounttherelative multiplet should also be the more depressed motionof thebinarysystemfollowingthere- one by the interstellar line. Actually, we have centcloseapproach,theincreasingdistance already noted in Section I I.VIII.B that, in betweenthetwostarsshouldcausea delayof some cases, the anomalous intensity ratios of thetimeof arrivalof thematterfromtheMira the OI resonance multiplet observed in the UV wind. spectrum of some symbiotic stars has to be at- tributed to the interstellar line absorption. In R Emissionlineprofilesobservedathighreso- Aqr, the interstellar lines are weak. We cannot lutioncantell usaboutthedynamicalstructure exclude that they could be partly responsible of thesystem.Forthisreasonandto haveas for the CIV structure in the core and in the much informationas possibleon R Aqr, nebula, but this possibility is excluded for the Michalitsianoset al. (1988a)recentlyat- largely anomalous CIV doublet ratio in the temptedto obtainhigh-resolutionultraviolet core. As Michalitsianos et al. (1988a) dis- imagesof RAqranditsNEjet.Becauseof the cussed, to explain the observed profile com- faintnessof the sources,theseobservations plex, radiative transfer effects should be con- requiredabouthalfa dayof exposure,butthe sidered, which require a complete analysis resultswerequiteinstructive.Michalitsianos under multiscattering conditions. According to etal.foundthattheCIVdoubletin thenebula spherically symmetric wind models for hot appearedbroad,possiblydouble,with a stars computed by Olson (1982) for resonance FWHMof about250kmsL Thedoubletinten- doublets whose separation is comparable to, or sityratioI(1548/I(1550)wascloseto theopti- smaller than, the wind velocity, the source callythinvalueof 2.InthecentralRAqrcore, function of the longer wavelength doublet theCIV doubletpresentedsomemulticompo- component depends on non local values of the nentstructurewith 2 or 3 sharpcomponents shorter wavelength source function. Radiation separatedby about40 kms-_.In thecore,the scattered in our line of sight by the blue compo-

698 nent can be scattered again by the red line, V.D. POSSIBLE MODELS FOR R AQR enhancing emission in the red wing of the 1550 A line.

The large amount of available data on R Aqr in The presence of a high-velocity wind can be all the spectral range should in principle aid in put in evidence by overlapping the two doublet building detailed models of the system. However, components as shown in Figure 13-18. In the fig- it is difficult to find models which are capable of ure, the shaded area represents those points of the describing in a consistent way the whole observa- line profile, in velocity space, where the mono- tional information. In addition, some funda- chromatic flux of the 1548 A line is smaller than mental parameters such as the distance of R that of the 1550 A line. Taking into account the Aqr and the interstellar extinction are still wavelength shift between the components, the uncertain. In the framework of binary models, shaded area corresponds to a velocity range from the Mira should have an unseen companion about -500 to -700 km s L This should be the producing the high-energy photons that ionize range of the P Cygni absorption oflhe 1550 A line the compact central HII region, and the nebula in order to reduce the emission of the 1548 A line. (Michalitsianos, 1984). The hot companion is Similar wind velocities can be derived for the probably hidden by a disk or by opaque matter other symbiotic stars showing this anomaly and in the orbital plane which is seen nearly edge- represent an indirect evidence for the presence of on (Figure 13-19), and/or by circumstellar high-velocity winds in symbiotic systems. dust. Its nature is still uncertain: present infor-

R AQUARII 6o7° I I - 77 I j, i I HIII REGION t

50 __ -27kms -1 (V R = -23kms -1) I 4.0

3.0 _ /,1550.8_,

Z ILl I-.- Z 2.0 UJ >

< ,,.J UJ rr 0.0

-1.0 - 57

- 2.0 -- -10

l I I ] I 1 I - 200 - 150 - 100 - 50 0 + 50 + 100 + 150 + 200

VELOCITY (km s-_)

Figure 13-18. Comparison of the profile of the two components of the CIV doublet in the spectrum of the central core of R Aqr. The 1548 A line appears much Jainter than the 1550 A line. Probably as the result of absorption by a high-vetoclty (500-700 kTn s -_) wind. The shaded area indicates the region in the velocity space over which 1(1548)/1(1550), 1 (Michalitsianos et al., 1988a).

699 mation is not sufficient to decide whether the radiation from the central source is largely high-temperature source is a hot, possibly reju- absorbed by the circumstellar matter. There- venated dwarf or the inner boundaries of an ac- fore, in order to explain the highly ionized cretion disk. In any case, the disk would be con- nebula, and the X-rays emerging from it, one siderably extended in the outer regions, where has to suppose that the ionizing photons are it should be much cooler and probably cause mostly emitted perpendicularly to the line of the temporary obscuration of the Mira dis- sight, from regions that are less occulted (Vi- cussed by Whitelock et al. (1983 a and b). The otti et al., 1987). Indeed, Kafatos et al. (1986)

R AQR

RADIO SOURCE

DISK MIRA

/ ,It°

HI /

He III H II

Figure 13-19. A model fi>r R Aqr. The neutral wind from the Mira giant is ionized by the UV radiation from the hot subdwatf and /or the accretion disk, and originates the intense central radio source. The hot source is obscured in the direction of the line of sight (perpendicular to the figure) by the accretion disk or by matter in the equatorial plane. Intense ionizing radiation is emitted from the poles inside a cone, and hits the circumstellar cloud

(the 'jets'), producing the high-ionization features (NV, Hell) and X-rays. Alternatively, the jets could he heated by shocks produced by the interaction of the hot source wind with the circumstellar environment.

700 explainedthepeculiarradiomorphologyof R proposed a high eccentricity of the orbit, in Aqr as a consequence of photoionization of order to allow large mass accretion through ejected material lying inside an ionizing radia- Roche lobe overflow at the periastron passage. tion cone, whose axis is perpendicular to the A geometrically accretion disk formed during orbital plane, and having an opening angle of this phase would produce at its inner boundary about 150 °. The hot source might also generate high-temperature photons, and possibly peri- the high-velocity wind revealed by the anoma- odic ejection of matter. The intensified radia- lous CIV resonance doublet discussed above. tion field would also cause ejecta from previ- ous outburst to brighten (Kafatos et al., 1986). The high-temperature emission from the This might explain the sudden (but apparent) nebula can be alternatively interpreted as emis- appearance of the A and B jets around 1970, sion from a hot plasma, which can be heated while according to Kafatos et al., the feature C2 material previously ejected from the central (Figure 13-17) would represent the most recent hot star. But, taking the electron density on the ejection. The disk may also be formed by the jet of 4 x 104 cm -3 derived by Kafatos et al. capture of the Mira wind, which is enhanced (1986) and an electron temperature of 3 x 105 during some periods by the close passage of the K, the cooling time should be around 5 x 106 s, binary components in a moderately eccentric much shorter than the supposed time elapsed (e < 0.5) orbit (Kafatos et al., 1986). Again, since the ejection. Such a warm matter, there- such a model better explains the episodic out- fore, would cool in a rather short time. A pos- bursts and ejections observed in R Aqr. To give sible heating mechanism could be the interac- a better insight into this problem, accurate tion of the ejected material with the circumstel- measurements of the Mira radial velocity over lar environment, producing shock waves. Vi- a few decades are needed. otti et al. (1988) found that, in this case, the V.E. PLANS FOR FUTURE OBSERVATIONS emission measure of the Hell emitting region, assumed to have a temperature of 3 x 10_ K, would be close to that derived from the X-ray R Aqr, for its relatively close distance (180- flux assumed to be optically thin thermal emis- 300 pc). and the many peculiarities represents an sion at the same temperature. The low-electron ideal target for future observations involving temperature derived by Kafatos et al. (1986) space experiments and high-technology ground using the CII] and CIII] line ratios would then telescopes. In particular, of special importance be referred to cooler, not shocked parts of the will be the high-resolution imagery at different nebula. A shockwave heating (or photoioniza- wavelengths. tion by a power law continuum) is also sug- gested by the [NII]/Hot ratio of 1.2 to 1.8 found Figure 3-20 shows a set of CCD images of R by Paresce et al. (1988) in the nebula. But so Aqr obtained using the Science far, it is not possible to decide which is the Institute coronograph, which occults the bright dominant heating mechanism of the jets. central star, thus, allowing a detailed study of the nebula very close to R Aqr with a subarcsec spa- The orbital elements of the R Aqr system are tial resolution. Figure 13-21 shows the derived unknown. The period might be of several dec- contour maps in the light of Hot and INIII 6584 ades, as suggested by the modulation of the emission lines, where many emission knots are light curve (Willson et al., 1981). Such a period easily detected. Ultraviolet and visual images and seems to be supported by radial velocity meas- polarimetry with subarcsec resolution will be urements (Andarao et al., 1985; Wallerstein, possible with the Faint Object Camera of the 1986). The separation of the stellar compo- , and should provide nents should be large to account for a large information about the location and physical struc- mass transfer from the Mira to the dwarf star, ture of the high-temperature regions, including and to feed the hot source. To overcome this the X-ray source, and about the presence and na- problem, Kafatos and Michalitsianos (1982) ture of circumstellar dust. Concerning this prob-

701 Figure 13-20. CCD images of R Aqr in a broadband R ftlter (left), and in narrowband fllters centered on Ha (center) and [N II] 6584 (righO. North is up and East to the left. Top row: original images, bottom row." final images in which the emission line contribution is subtracted from the R filter image, and the glow from the central star is subtractedJkom the narrowband ftlters (from Paresce et aL, 1988).

lem, we expect very interesting results from the tribution of the cool matter and dust. In this new high-quality infrared arrays. Recent IR im- regard, R Aqr probably represents a unique agery of R Aqr at 3.45 !am led to the discovery target to study the nature and structure of of an extended spherically symmetric halo that circumstellar dust, and the interaction of the extends to about 15 arcsec from the central star stellar radiation and wind with the circumstel- (Schwarz et al., 1987.) Higher resolution IR lar environment in an evolved object. imagery is needed to determine the spatial dis-

702 R Aqr Ha R Aqr [N II] k ! ' ' "---_-".! .... I .... I .... 1

_-_ , _. O

o _r'_ ,, ,',' .... _!_

LO 0 . ,,i; ',

z

(j)o ÷ O z z co 0 c_ .o z LU CO o (.-) I n," _ _ . o_ • , , _,_ ,< o .< 0 I I

if') T- I I

o I I ; I .... I .... 1 .... I.., , 1 -10 _ 0 5 10

ARCSECONDS W ARCSECONDS W

Figure 13-21. Contour maps of R Aqr in the light of H_t (le[i) and[Nil] 6584 (righO as derived from the images in Fig 3-20 (Paresce et al., 1988).

VI. CH CYGNI: ANOTHER SYMBIOTIC emission lines were present in his spectra. VARIABLE WITH A JET Smak (1964) observed CH Cyg using narrow filter photometry during a period of 82 days. VI.A. INTRODUCTION The visual magnitude varies l¥om 7.06 to 6.64, and the color indices indicate that the star was bluer when fainter. This behavior is common to CH Cyg is classified as an Mb star in the HD all M-type variables and could be ascribed to catalogue. It has long been known for its the TiO absorption bands, that affect the mag- semiregular light variability with a period rang- nitudes B and V but not U, and are stronger at ing from 97 to 101 days (Wilson, 1942; Gapos- minimum. The has chkin, 1952; Payne-Gaposchkin, 1954) and was been observed to vary between 7.9 and 9.1. classified M6 III. Beside this short-time variabil- ity, a long cycle of 4700 days or 12.8 years was All the existing observations ofCH Cyg before found, not very different from that recently de- 1963 indicate a normal M6 llI semiregular vari- rived by using all the radial velocity measure- able. In September 1963, Deutsch (1964)ob- ments tound in the literature by Yamashita and served that CH Cyg "showed a composite spec- Maehara (1979) of 5750 + 250 d and by Hack et trum. A hot, blue continuum was superposed over al. (1986) of 5,000 + 450 days. According to C. the late-type spectrum, together with emission Payne-Gaposchkin (1954), the median maxi- lines of H (strong and wide), He I (weak and mum photographic magnitude was 7.97 and the wide), IFe II] (strong and narrow) and Ca II (also median minimum, 8.44. Joy (1942) measured strong and narrow). The spectrum of the recurrent the radial velocity at different epochs and nova T Cr B closely resembled this in June 1945, found an almost constant velocity (from -51 to a few months before the outburst of 1946. No -59 km/sec) for the M6 absorption lines. No doubt, there is a nova-like in CH Cyg

703 system,too.However,a spectrogramof March detectable; the star reached V = 6.4 at the begin- 1961showednotraceofthehotspectrum,which ning of the outburst, and after three years at al- hasfadedappreciablysinceitsdiscoveryin Sep- most constant magnitude, reached V = 5.6 at the tember1963."Onehigh-resolutionspectrum end of 1981 and remained close to this value un- takenattheHauteProvenceObservatoryin Au- til mid-1984. A sudden luminosity drop of about gust1965(FaraggianaandHack,1969)showsno one magnitude occurred between July and August evidenceof symbioticcharacteristics.Theonly 1984 and was accompanied by strong spectral peculiaritywasanemissionatHalphaandHbeta variation (for instance, weakening of the blue andpossiblyatHgamma. continuum, see Section VI.D.I)

A secondsymbioticepisodestartedin June The variation of the B-V and U-B colors from 1967(Deutsch,1967)andwasover by late au- 1967 to 1985 are described by Hopp and Witzig- tumn 1970. A third episode started in 1977 and mann; by Panov et al. and by Mikolajewski and is almost over at the time of writing this report Tomov. (July 1988). Luud et al. (1978) detected no trace of H alpha emission in spectra taken on In the period of quiescence 1970-1976, B-V is May 1976, but in May 1977, H alpha and H beta generally bluer at minimum as it was observed by showed a double emission peak with V/R>I. In Smak. In general, U-B is in phase with B-V with August 1977, the presence of a blue continuum stronger fluctuations. Both indices become and of several emission lines was evident smaller (star bluer) at the beginning of the third (Fehrenbach, 1977; Morris, 1977). outburst. During the outburst, B-V fluctuated between 0.4 and 0.6 and increased from 0.49 to Spectra taken in July 1986 show the pres- 0.82 from Aug. 10, 1984, to September 30, ence of numerous strong emissions, although 1984. U-B varies from 0.6 in mid-1976 to about the hot continuum has practically disappeared -0.4 during the outburst, with oscillations be- between November 1984 and January 1985. tween -0.3 and -0.7. After July 1984, U-B The emissions are fainter in 1987, but they are gradually increased. On November 24, 1984, it still easily detectable. In 1988, only Hot and HI3 was equal to -0.29, and in May 1985, to about are strong in emission. The strongest [Fe II] 0.0. (see Figure 13-22, light curve and color emissions are very faint, while the Fe I1 permit- variation). ted emission lines have almost completely dis- appeared. In addition to these long-period variations, which are typical of semiregular late-type vari- This last outburst episode has been observed ables, short-time scale (minutes), small ampli- with the IUE satellite since April 1978 (see Sec- tude variability has been observed during the tion VI.D.3 on the UV spectrum). periods of activity and will be described in the fol- lowing sections. VI.B. THE LIGHT CURVE OF CH CYGNI VI.C. THE 1967 OUTBURST

The light curve ofCH Cyg from 1899 to 1975 has been described by Gusev (1976) and is illus- The outburst of 1967-70 was followed both trated in Figure 13-22. Until 1960, only small photometrically and spectroscopically by several amplitude variations are present. The following observers. period is described by Hopp and Witzigmann (1981), by Duschi (1983), Panov et al. (1985), A detailed description of the spectral vari- and Mikolajewski and Tomov (1986). After the ations from July 1967 to December 1970, cover- 1963 outburst, the oscillations became more evi- ing the whole duration of the second observed dent and regular with a mean period of 700-800 outburst, has been given by Faraggiana and Hack days during 1967 to 1977. During the third out- (1971). Photometric observations during the burst, the semiregular variations were no more same period have been made by several authors,

704 a)

I I I I 1900 1905 1910 1915 1920 1925 \ I

I I I I I

1930 1935 1940 1945 1950

J _t J •

-__ . . , , / \_.."_,/ /

I I I I I i i 1955 | I 97o 1975 1963 1967

1968 1971 1973 1976 t979 1982 b/ MAY 23 FEB. 19 NOV. 13 AUG 9 MAY 6 JAN 31 m v I 1 I I I 1 r ,

6.0 .°

.ll'.

7.0 • I ,_ .% j. .'. ,.. "I , I _ . _" " "_4 ". ;'"1. ), i.'. t ._ .,:," i _..: "t,. : -'_._""_ ., 8.0 '" t • "s --.. •

= I l_ I l I 0000 1000 2000 300O l 4000I 5000

JD 244... 1977 AUG. 27

c) 198411985 -1 -1 _MA'R. MAY JU'IL._IIISEPT. 'NOV. JAN! MAR. 'MAY

E 0 L ° U-B "¢£-_"- o L o--__ . ,A" 0 / 0 ...... _41. _ _ _ _ / _ 1 I- B-V _o--O'---oO.._ x _-A/ 1

6

7J- '=- -"¢t)--5.-4%__

/ I I I I I He'd- I _ I I 5800 6000 6200 6400 6600 JD 244...

Figure 13-22 a) The visual light curve of CH Cyg from 1900 to 1975. The arrows indicate the of starting oJ the two outbursts of 1963 and 1967 b) Visual light curve from 1968 to 1981. The arrow indicates the starting oJ the outburst of 1977. c) Visual light curve and B-V, U-B variation in 1984-1985. d) light curve and U-B variation from 1967 to 1982 (from Gusev, 1976; Duschl, 1983; Mikolajewski and Tomov, 1986; Mikolajewski, 1985; Panov et al., 1985; Mikolajewski and Wikierski, 1986, Hopp and Witzigmann, 1981).

705 (k) ---

1970 1975 1980

. | ,l ! I j I J | I 1; | • • U-B

0 o ._.._ CH Cyg o.f_._. -0.2 :(':• * @ :,,". :.;-.,L • + 0.6 • -._ .,- ""'_ t.;._" "_,. ,r " .: 0.2

pe e j - B-V • ,L; ..r,,,_,_.t_,_ '"v 1.0 5.6 • ...... ¢ • _,._ ._, _¥_, -,,_ :;_L m" .-- "" _q-'l_'_ ^ t,.,,'.._, ",- -- • _" ,rJ. :!,. 1.8 V

7.2

,._, "I _. :• ,_ "l t..: _,. i "t "_ej._"" "_: "1""J 8.0

tm * t_ I I f. 4 I J ! 2440000 ( / 2442000 2444000 tp.D.l The rapid variations (flickering) were first noted 1970; it appeared in June 1967 (Deutsch, 1967) by Cester (1968) and by Wallerstein (1968). The and reached a maximum in August 1968. The amplitude was of the order of 0.1 mag on a few color temperature of the blue continuum was minutes time scale. Cester (1969) found that the. about 10,000 K. No measurable Balmer disconti- variations were strongly correlated in different nuity was observable. colors, and the color indices were quite different from those typical of an M6 III star: B-V oscil- b) Emission lines of He I, Fe II, [FeIl] and [Sill lates between +1.3 and +1.0, and U-B, between are present. The Balmer lines and the H and K -0.05 and -0.6. As a rule, the amplitude of the lines of Ca II present a P Cygni profile. In July and flickering is larger at shorter wavelengths. For August, H and K presented two sharp absorption instance, Shao and Liller (1971) reported vari- cores at radial velocity of about -75 and - 160 km/ ations of 32% at 3,200 A and of only 2% at s, while the radial velocity of the photospheric 7,000 A. Figure 13.23 shows the flickering ob- lines (TiO bands and nonresonance lines of neu- served in U by Luud et al. (1970) on November tral metallic atoms) ranges between -50 and -60. 5, 1968. Walker et al. (1969) made photoelec- in July 1968, the nebular line 5007 [Ollll ap- tric spectral scans from 3300 to 5000 A during peared; 4959 and 4363 [OIII] were not visible. August 1967. Figure 13.23 shows the excess The emissions reached maximum intensity in continuum radiation relative to the standard August 1968 and again in August 1969. In 1965 M6 lllab spectrum of 45 Ari for several scans and 1966 and in September-December 1970, the made during the night of August 3, 1967. Vari- spectrum was a normal M6 III type, except that ations on time scale of a few minutes are evi- H alpha and H beta presented emission compo- dent. nents.

c) Low-excitation absorption lines of metallic The main characteristics of the spectrum and ions appeared in the ultraviolet continuum in July its variations during the second outburst can be 1968. In May 1970, several lines of neutral ele- summarized as follows (Faraggiana and Hack, ments appeared. At this epoch, the ultraviolet 1971): lines of low excitation ions and the strong reso- nance lines give radial velocities more nega- a) The M6 spectrum is veiled by a continuum tive than the other lines having the same low that partially fills the absorption lines and in- level by about 20 km/s and are not filled in by creases in intensity toward the violet. This contin- the blue continuum. This may suggest that they uum was absent in August 1965 and in September were formed in the blue continuum like a kind

706 a) Ore3 _ &m

Ore2 _

Oral _

OmO -

_Oral -

_ Ore2 I 17h20 17h30 17h40 17h50 18bOO 18'_10 18h20 UT b) Au ' ' ' 1 ' ' ' ' ' ' i , , , I '' ' ' ' ' ' ' ' ' ' I' ' .....

CH Cyg: 1984 July 31

4/

0.00 ! | I ! v ! " gv '

+0.10 --

+ 0.20 --

, i J I I I , , n J J J J J n J , _ I I I J I n n I n * I J J * I 20h 21 h 22r' UT Figure 13-23. a) Observations (,J U fli_kering m_ No vember 5, 1968 (J)om Luud et al., 1970), b) and on .hdy 31, 1984 (f_om Pam)v et al., 1985.

of shell absorption. Significant spectral vari- sity, while the line intensity was variable. ations have been observed over a few days. For instance, the Balmer lines were in emission on VI.D. THE OUTBURST STARTED IN 1977 May 13. 1970, and in absorption on May 16, 1970. The last outburst began in 1977 and was fol- lowed by several observers in a wide spectral The ratio [Felll/Fe I1 is constant during the range from UV and X-rays to optical, infrared, whole outburst, indicating no change in den- and radio.

707 c) 3400 36 38 4000 42 44 46 48 5000 1 I 1 1 I I 1 1 I

-- 3 AUG. 1967 SCAN 5 7h 38 m U.T.

-- • 0000 O0 • 00 00 - • • O O • • • • O _0 J-e-#- - _o_

SCAN 6 8h 02 m

D m

B m

0000000000000

m ......

SCAN 7 8h 16m 3 000• 000 2 "- 00000 • 1 000 0 OOe O°°. *° o -1

SCAN 8 8h 23 m 4 O o

3 m 00000000 • 2 • 000 •

1 • 00

0 •ee° qDa__o ....

-1 I I I I I I I 1 I 34O0 36 38 4000 42 44 46 48 5000 _.

Figure 13-23. c) Excess continuum radiation in the spectrum ofCH Cyg relative to the standard M6111ab spectrum of 45 Ari,for twelve seans taken on the night of A ugust 3, 1967. Ordinate is influx units with an arbitrary scale factor applied (from Walker et al., 1969).

708 c) 3400 36 38 4000 42 44 46 48 5000 1 I I I I I I 1 I

6 3 AUG. 1967 SCAN 1 5 - 7 h 3 m U.T.

4 -

3 - gO OoO OO • OoOO 2 -

1 - • Oldn, • oOo • • 0

SCAN 2 7 h 11 m

4 R

3 m • • oOo0 O0 0 2 O o 1 m O ° • O • • • 0 '9

-1 SCAN 3 7h 24 m

3 O o 2 oooO0 0 OO 1 oO0 OO0 O o • 0 "OoO eem e o -1 SCAN 4 7h 30 m

4 D

•O 3 m • • 2 O0 O0 • OO OO0 1 • O O • _ • oo_._, d,O • • 0 --'ml",-O _ _6"-- I L t I I L 1 1 I 3400 36 38 4000 42 44 46 48 5000 ),

Figure 13-23. c; cmztinued.

709 t..t.I o

I-.-- ms

o m_" I i .,:£

° T J w_ oQee • o No 11 el ee• _i k: ...-, • _ o • •• • D• •e •loeeQIoe e• • _ • •• ! • •• •t

,e

11 I J 1 1 I I 2440000 2000 4000 60O0 JD Figure 13-24. Ultraviolet light curve of CH Cyg showing the two minima in 1969 and 1985. (e) U magnitude, + V magnitude measured from the ground (x) V magnitude measured with the Faint Error Sensor, (FES) aboard IUE, (o) m (3000 A) derived from the IUE fluxes: m(3000) = -2.5 log F(3000) -22; (A) m(1440) = -2.5 log F(1440) -29,. The constant 22 has been arbitrarily chosen.

The International Ultraviolet Explorer (IUE) which are typically observed in late-type giants has given us the opportunity of follow the out- and supergiants, these measurements suggest a burst in the ultraviolet since April 1978. period of 13 to 15.5 years (Yamashita and Maehara, 1979; Hack et al., 1986). Our last VI.D.I. PHOTOMETRIC OBSERVATIONS observations, added to all those existing since 1961, suggest P = 14.4 years or 5250 days, not As we have observed in Section VI.B, the very far from the value suggested above. (see semiregular light variations with a period of 700- Figure 13-27, radial velocity curve). High- 800 days disappeared during the third outburst speed photometry confirmed the presence of when V gradually rose to 6.4 at the beginning of flickering during the present phase. Slovak and the outburst, then rose to 5.6 and remained African• (1978) observed an amplitude of at this value until July 1984, when its bright- about 0.10 mag in the ultraviolet and violet ness dropped by about one magnitude light (u and v filters) and of only 0.03 mag in y (Mikolajewski and Tom•v, 1986). One year filter. Actually, two features characterize the later, V was equal to 7.8. light curve: rapid flickering on time scale of 5 min and amplitude 0.02-0.04 mag and slow, The U magnitude and the ultraviolet flux large amplitude flares (0.10 mag) lasting 15-20 measured with IUE displayed a different be- minutes. Panov et al. (1985) give a summary of havior (Mikolajewski et al., 1987; Mikolaj- the flickering amplitudes observed during this ewski et al., 1988), i.e., a minimum lasting outburst (Table 13-4). The data obtained by about 150 days (May - Oct. 1985). A broad Cester (1969) during the outburst of 1967 are minimum, especially in the U band (Cester, given for comparison. 1972; Luud et al., 1977), was observed in 1969. The possibility has been suggested that these After the drop in brightness of July 1984, a two minima--separated by about 5700 days-- drastic change in the spectrum was observed at were two consecutive eclipses of the hot com- the end of 1984, with the almost complete disap- panion of a binary system (Mikolajewski et al., pearance of the hot continuum (see next section), 1987) (Figure 13-24). while a large rise of the radio flux was observed between April 1984 and May 1985 (Taylor and Although the evidence of orbital motion Seaquist, 1985; Taylor et al., 1986). given by radial velocities measured from 1942 to 1986 is weak because of the large scatter due A fresh outbreak of activity (the end of the to the irregular radial velocity variations, eclipse?) is indicated by the photometric observa-

710 TABLE13-4.FLICKERINGAMPLITUDESOFCHCYGNIFROM1977TO 1984(FROM PANOVETAL., 1985)ANDIN 1968(CESTER1969).

Observers _u BB BV

SlovakandAfricano(1978) 0.10 0.03

Ichimuraetal.(1979) 0.18 0.13 0.10 Luudetal.(1982) 0.20

Spiesman(1984) 0.23 0.21 0.15

ReshetnikovandKhudyakova(1984) 0.20 0.20 0.20

Panovetal.(July31,1984) 0.6

Panovetal.(Nov.24,1984) 0.24 0.17 0.12

Cester,1969(July13,1968UT23h15) 0.24 0.24

Cester,1969(Aug.20,196821h45) 0.24 0.12

Cester,1969(Aug.24,196822h28) 0.44 0.24

Cester,1969(Aug.25,196823h05) 0.24 0.12

Cester,1969(Aug.25,196824h00) 0.36 0.24 tionsofPanovetal.(1985):OnAugust17,1985, 1977to November1984is characterizedby a U=7.38,U-B=0.97,i.e.,thecolornormallyob- continuum,thatis strongerat shorterwave- servedforCHCygin quiescence(Smak,1964), lengthsbutisalsopresentin theredpartof the aswellasforseverallate-typesemiregularvari- spectrum,andpartiallyor almostcompletely ables.OnDecember9, 1985,U = 7.31,U-B= veilstheM6absorptionfeatures.Thestrengthof -0.021.Asweshallalsoseeinthefollowing,the thiscontinuumis evidentwhencomparingthe spectralchangesconfirmaslightstrenghtening regionsoftheTiObandheadsorthatofthebroad of thebluecontinuum,whichwasveryweakin absorptionof CaI at4227A observedduring January1985. theoutburstwiththoseobservedin 1970,after theendofthesecondoutburst,orin 1988,when VI.D.2.SPECTROSCOPICOBSERVATIONS theoutburstwasalmostover(Figure13-25).

High-resolutionspectroscopicobservationsof Thepreludetotheendoftheoutburstwasob- CH Cygniduringtheperiod1977-1986were servedfromDecember1984to January1985, madebyseveralgroups(Hacketal.,1982;1986; whenthevisualspectrumbecameagainthattypi- 1988;MikolajewskiandBiernikowicz,1986; calforanM6III star,afterthedisappearanceof Mikolajewskiet al., 1987;TomovandLuud thebluecontinuum.Theemissionlines,how- 1984;Wallerstein,1981; 1983;Wallersteinetal., ever,werestill presentandprominent;more- 1986;YooandYamashita,1984).A generalde- over,besidetheemissionlinesof HI, FeII, [Fe scriptionofthevisualspectrumanditsvariations Ill, 10 I1,HeI, ISII],emissionsof [OIii1 and wasgivenbyHacketal.(1986,1988).Thespec- [NeIII] appearedin November1984andrap- trum,fromthebeginningoftheoutburst,in May idly increasedin intensity.Thecomplexvari-

711 MAY 5, 1981 Hd 4227 Ca -0.40

-0.80

- 1.20

o _ 1.60

0 "J -2.00

- 1.20

- 1.60

-2.00

-2.40

- 2.80 - JULY 21, 1988

-3.20 , t , 1 , [ , I , I , I , I 38O0 3900 4000 4100 4200 4300 4400 4500 ANGSTROM

Hy MAY 5, 1981 H/3 5007 O III -0.40 I

-0.80 nO

- 1.20

- 1.60

- 2.00

- 1.20

- 1.60

-2.00

- 2.40

JULY 21, 1988 - 2.80

-3.20 , I I , I , I , I , I , [ • 4300 4400 4500 4600 4700 4800 4900 5OOO ANGSTROM

Figure 13-25. Intensity tracings of two spectra obtained on May 5, 1981 during the outburst and in July 21,1988, when the outburst is almost finished (Spectra obtained at the Haute Provence Observatory).

712 ability of the line intensities, line profiles, and a)

radial velocities is described by Hack et al. 3854 3862 3870 3878 ,t (A) (1986, 1988). Figures 11-8, and 13-26 give some examples of line variability.

1 I' ' ' ' ' ' ' ' 'UL51984'1 It should be noted that the M6III photospheric lines and the forbidden and permitted Fe II emis-

sion lines have radial velocities that are 180 ° out 1 __NOV 24 1984 of phase, suggesting that the Fe II emissions are 0 - associated with the companion (Figure 13-27). NOV 27 1984 The relative amplitudes of the two curves sug- gests a mass ratio of the two components close to

1. Also, the absorption cores of the inverse P 0 Cygni profiles of the metallic ions give radial ve- JAN 29 1985 locities 180 ° out of phase with the primary pho- tospheric lines (Figure 13-28). Mikolajewski et al. (1987), by the metallic ion absorption lines, 0 derive 3 < m I/m2 < 4. However, the radial veloci- ties of the inverse P Cygni absorptions are af- 1

fected by the presence of the emission wings, {_liUL 8 1985 which partly mask the absorption cores affecting 0 their measured shifts. NOV 29 1985 1 -

The most characteristic features that distin- 0 guish the second and third outburst are the follow- _N 20 1985 ing:

1) A similar behavior of the blue continuum, however, it was bluer in 1977-84.

0 2) At certain epochs during the outburst, sev- JUL 22 1986 eral metallic ions and Balmer lines showed direct 011 __Ji UN 3 1986 P Cygni profiles in 1967-70 and inverse P Cygni profiles in 1977-84.

3) The Ca II H and K lines, on the contrary, displayed the same behavior during the two out- burst showing one or two, occasionally three, vio- F&ure 13-26. a) The region of 3868 [Ne Ill]from July let-shifted absorption components and always di- 5, 1984, to July 7, 1986. b) The region of 5006.84 [0 rect P Cyg profiles, which clearly indicate the ex- 111], 5015.675 He 1 and 5018.434 Fe H (Spectra istence of one or several expanding envelopes. obtained at the Haute Provence Observatory). The highest observed expansion velocity is - 100 km/s. constant during the 1967-70 outburst, while dur- ing the 1977 outburst, it was variable with a sig- 4) In both outbursts, emission lines of fairly nificant decrease at the end of the outburst. These high excitation were observed, i.e., He I during ratios indicate that the density in the region where the whole outburst and [O III] at some phases. [Ne the permitted and forbidden lines of Fe II are III] 3868,74 was observed only during the latter formed was lower in 1967-70 than during the lat- outburst. ter outburst. The presence of the lines of [O I] at 6300, 6363 A indicate a density of about 106 5) The ratio between Fe II and [Fell] remained cm _, while [Fell] indicates 10" -10 _ cm -_.

713 b) 4990 5000 5010 5020 5030 X ('_)

1 4990 5000 5010 5020 5030 _'(/_)

I I I ! I II I l I I

14 1978

/ _ FEB 9 1980

4990 5000 5010 5020 5030 '_(]_) 4990 5000 5010 5020 5030 _ ('_) ,i ! _ I I I I I 1 I I I

6 1981

0

714 -50

#°oo_ o o o_, "._'.• .• o o rr _. • o .e o x -60 .o "..._..

o. • - 70 I I I I I I I I

J

t

I -60 - e • aQ

x x x x

- 70 x

x _ x x _ x _ xe

-80 1 I I I I I I I 24 39000 40000 41000 42000 43000 44000 45000 46000J D

Figure 13-27. Bottrml." radial velocities q/'the M6 III photo_wheric lines (0) and of Fe Il permitted emission lines (x) vs. J.D. Top: Radial velocities qf the forbidden lines q# Fe I1 (IlL 0 I (o) and 0 111 (x),from Ha_'k et al., 1986). +10 --

rr-

-10 m & o:

x • - 20 x• • 0 x

O -30 B + O 0 O O0 o

- 40 O _- x + + + x x + + + xo -50 O O × x -60 O

- 70 -- I I I 244 3000 4000 5000 6000 J D

Figure 13-28. Radial velo_'ities of the absorption cores of Fe II (o), Ti II (e), Mg I (x) amt mean q/Ti 11, 11, Cr II, Mn II (+) (J)'om Ha_'k et al., 1986).

715 a) I ! I 105 29 Nov '81

20 Jan '84

x 110

J Lt_ © © J

Apt-Dec 80

Jan-Dec '7g 115

t6 Jul '86

- 12.0

Dec '85-Mar '86

19 Oct '86

- 125 May-Oct 85

12 Dec '86

-!30

4 Sept86 I I I I 1500 2000 2500 3000

WAVELENGTH [_\]

Figure 13-29. a) Continmmz energy distribution in CH Cyg, b) Combined UV low-resolution spectra (1200-3200 A) i_CH Cyg, c) Fat UV low-resolution specta of CH Cyg. (from Mikolajewska et al., 1988).

VI.D.3. THE ULTRAVIOLET SPECTRUM strong variation with time. Not only the flux var- ies but also the shape of the continuum, which becomes completely flat in January 1985, re- The ultraviolet spectrum was also observed mains flat to October 1985 and starts increasing during the whole outburst with IUE in the low- toward the longer wavelengths in December resolution mode (6 A) and, when possible, also in 1985. A drop in the flux intensity by a factor of the high-resolution mode (0.2A). The variation of three was observed between January 24, 1985, the UV continuous flux from 1200 to 3200 A dur- and May 27, 1985. Unfortunately, no other IUE ing the period 1978 to 1986 was studied by observations were made during the rest of Mikolajewska et al. (1987, 1988). Figure 13-29 1984. Hence, we cannot decide whether the shows the ultraviolet energy distribution and its drop observed in visual light in August 1984

716 b) T T 6.00 1-t SWP 27216 : LWP 7232 r SWP 5881 + LSR 5136 I _ < • 10:' . "102 July 21, 1979 II ? 300 Eo 400 E o_

Q o 2.00 2.00

100 i O.O0

400

SWP 22056 + LWP 2674 10_ • Januap/ 21, 1984 000

2,00 • 102 SWP 28010 + LWP 7857 I

March 24 1986 300

100

200

0.00 _ L

• i02 SWP 26219 + LWP 6251 June 22, 1985 100

300

000 3.00 2,00

• 102

2.00

1.00

SWP 26720 + LWP 6797 • 102 September 25, 1985

000 300 3.00 SWP 29856 + LWP 9686 • 102 December 12, 1986

200 200

100 100

000 0CO 12000 1700.0 22000 27000 320(3 0 120(} 0 17000 2700 O 32000 was also present in the UV. In December 1985, it is interesting to note that at k 1300 A, there is the far ultraviolet flux started to rise again (in evidence of a slight increase of the flux toward concomitance with the increase of the contin- the shorter wavelengths, suggesting that we are uum observed in the visual) and reached a sec- now observing the Rayleigh-Jeans tail of a hot- ondary maximum in July 1986, while in Sep- body radiation, a tail whose presence had been temper 1986 it declined back to the level of previously excluded on the basis of observa- October 1985 (see also Figure 13-24). In 1988, tions obtained until 1986-1987. the flux in the far UV had decreased considera- bly (to about 20 times weaker than in 1986), but The ultraviolet line spectrum and its variations

717 cj _ 2.oo I I I 3200/_) and on September 1980 in the far UV 4 ,.< • 103 SWP 15591 (1200-2000 /_,). The spectrum was character- Novembel" 19. 1gel ized by the presence of both emission and ab- E sorption lines. We observed the emission lines of O I 1304 and O I] 1641, C III] 1909. 73, Si III] 1892, the multiplet 191 of Fe II at 1785, while all the other Fe II lines in the far UV were in absorption. Other absorption lines present in 0.50 the spectrum of September 1980 were C IV, Si IV, A1 II resonance lines, and a large number of

000 ! I ground level lines of Ni II. The strong Fe 1II

SWP 24956 • 102 lines from excited levels were not present. The January 23, 1985 Mg II resonance doublet presented a P Cygni 3.00 profile with two absorption components, one at almost the rest velocity, probably of interstel- lar origin, and another shortward-shifted by 2.00 about - 110 km/s.

1.00 Hack and Selvelli (1982) discuss the excita- tion mechanisms for the O I and Fe lI lines, which present some intriguing problems. For instance, 0,0O I the semiforbidden line 1641 O I] has about the • 102 SWP 25832 May 2, 1985 same intensity as the strongest of the three per- 3.00 mitted O I lines at 1302-1306 A, and the ob- served intensity ratio within the triplet 1302, 1304, and 1306 in September 1980 was 2,00 1:9.4:6.5 instead of the theoretical one of 5:3:1. The large optical thickness indicated by the

1£0 strength of the permitted multiplet means that the 1304 photons will scatter many times be- fore escaping from the region of neutral oxy- ! • lo2 gen. Substantial reabsorption will occur mostly SWP 28682 from the 0.00 eV level, and, therefore, the 1302 July 165, 1986 3.00 line will be weaker than the two other lines. There is a small but finite chance that at each coherent resonance scattering, decay from the 2.00 upper term 3s 3S° wilt occur through the 1641 line, which shares the upper term with the 1304

1.00 multiplet. Hence, the great optical depth of the 1304 multiplet has the effect of converting the resonantly trapped photons into 1641 photons, 0.00 I I I which will escape easily. This phenomenon is 1400.0 1600.0 1800.0 observed in several emission line stars, like Z And, V 1016 Cyg, RR Tel, HD 45667: all pres- during the outburst are described by Hack ent an anomalous intensity ratio in the 1304 (1979), Hack et al. (1982), Hack and Selvelli triplet and the strong 1641 emission. Another (1982), Boehm et al. (1984), Persic et al. characteristic common to several stars with (1984) and by Selvelli and Hack (1985a,b). extended envelopes is the presence of multiplet High-resolution spectra were obtained for the 191 of Fe II in emission, while the other far first time on March 1979 in the near UV (2000- ultraviolet Fe II lines are in absorption. In this

718 case,wehavea resonancefluorescencemecha- tensity ratio of the two lines indicates a decrease nism,i.e.,absorptionin thefar ultravioletfol- of the electron density to about 5 x 106 cm 3. lowedby reemissionat longerwavelengths: The N V resonance doublet, which was never a6D....>x6p°(UVmult.9,XN1260inabsorp- observed before either in emission or in absorp- tion) tion, and C IV, and Si IV, which were previouly x6P° .... > a6S (UV mult. 191, _.N 1785 in emis- present in absorption, are now in emission. The sion) strong multiplet 34 of Fe III is present in emis- a6S .... > a6D (opt. mult. 7F, 4287-4475 in sion. Fe II, which before January 1985 was one of emission) the principal components of the absorption line spectrum in the far UV, starting with January The spectrum observed in the high-resolution 1985, changed completely to emission. mode at the end of 1981 has about the same gen- eral appearance as it had in September 1980. It is VI.D.4. INFRARED, RADIO AND X-RAY OBSERVATIONS noticeable the presence of practically all the Ni II absorption lines up to multiplet 30 (low EP=2.8 eV). IR (1-20 lam) observations during the out- burst have been made by Ipatov et al. (1984). Unfortunately, no observations were made The near-lR low-resolution spectrum of CH from the end of 1981 to January 1984. The spec- Cygni observed at different epochs in trum in 1984 shows that the multiplet at 1303 of quiescence and in outburst remained nearly O1 1302-6). The high-resolution spectra are domi- unchanged for wavelengths longer than 7000 A nated by numerous and strong emission lines and is very similar to that of the M giants Alpha the same as in 1981. A spectacular change was Her and g Her. detected in Junuary 1985, at about the same epoch of the disappearance of the blue continuum in the The light curve from 1978 to mid-1983 in J visual range (Selvelli and Hack, 1985a). The (I.25 lam) and in H (1.6 lam) shows long-term continuum in January 1985 has become com- variability never exceeding 0.4 mag. pletely flat, and the line spectrum has dramati- cally changed from an absorption-like to an emis- The IR energy distribution observed in 1982 is sion-like spectrum (Figures 13-30 and 13-31); compared with the standard energy distribution see also Figure i i-29a Lye, and Figure 11-29b of an M6 III star. An IR excess is present long- OI 1302-6). The high-resolution spectra are dom- ward of 3.5 p.m up to a factor of 10 at 20 lain. inated by numerous and strong emission lines whose peak intensity rises to about 100 times the The Infrared astronomical satellite (IRAS), continuum. No absorptions are observable, also which operated for 10 months in 1983, has ob- due to the weakness of the continuum. The emis- served CH Cygni (Kenyon et al., 1988), which sions range from neutral species like O 1 and N 1 allows us to extend its energy distribution curve to highly ionized species like C IV, N V, and Si in 1982 to 100 lam. The flux at 12 and 25 lam IV. New remarkable characteristics of the 1985 agrees well with the ground-based observations spectrum are: by Ipatov et al. at 10 and 201am. The infrared ex- cess in the IRAS range remains of the order of I 0

a) The appearance of a strong and wide Ly with respect to the standard M6 I11 energy distri- Alpha emission (Full width at zero intensity = bution (Figure 13-32). IR spectra (1.5-2.5 lam 16.4 A). The emission is cut by an absorption covering the time interval February 1979 to the centered at rest wavelength (possibly of interstel- end of 1984 have been made by Hinkle et al. lar origin) 3.8 A wide. (1985). The spectrum is that of a typical M giant, with the exception of very weak Brackett gamma b) The appearance of the other faint C III] line emission. The velocities have been measured from the CO bands. It varies with a time scale of at 1906.68; all the other spectra only showed the strongest line of the doublet at 1908.73. The in- several hundred days but does not have a single

719 CH Cyg SWP 74956-360 s

I

E

m

t o

X - II 1 Fe,,(191, ii IL °i' s,,,,,

• i , i I , , I , i I , i I , , I , , I 1200 1320 1560 1920 (k) --,

8 CH Cyg SWP 22056.30 s

Fe 11(191) Sl III](1)

i

i ct_ 7 E O I](145)

q} c_

O _v X

U_ 8

1320 1560 1920

Figure 13-30. Far UV low-resolution spectra of CH Cyg taken on January 23, 1985 (top), and on January _0, 1984 (bottom). periodicity, according to these authors. They fects the optical and UV region. Observations in found a median velocity of -63.5 km/s with an the radio range give more exciting results. amplitude of 9 km/s, close to the values obtained from the visual region. Taylor and Seaquist (1985) were monitoring several symbiotic stars. During the period April Hence, the IR observations do not show any 1984 and May 1985, they discovered that CH Cyg clear evidence of the outburst that so strongly af- underwent a strong radio outburst coincident

720 '' '' I'' '' I .... I .... I'' ' '

3C ._

(300 0000_' 0000£ 0000_ O000L

;Or. X{"17:_ 31"11

0000_, 000

zOL. xnq::l BAI

.... I ' ' ' --Ib_ .. 8o t _

i ._- ._

O0 00_. 000

;Ot o XN7-J 3Al

721 10 I I I i I

CO ? E 0 (.3 11 I 03 I LL I / 0 0 I ._.J I I I I I [] 12 I I

[]

13 x+

\

14

\ \

15

_6 - 1 I I 4- 0.1 1 10 100

,[, _m

Figure 13-32. Composite energy distribution of CH Cygni corrected for an interstellar extinction of E(B- V ) = 0.07 compared with a standard M6 I11 distribution (dashed line), o September 1980," [] June 1981; • December 1981 ; x May-June 1982, • IR observations fi'om Ipamv et al., 1984; + IR observations from IRAS (fi'om Kenyon et al., 1988) (adapted from lpatov et al., 1984)

722 with the appearance of a multicomponent jet, source HI926+503, Int.= 1.22 count/s. In fact, it falls near the center of the error box of a hard-X- expanding at a rate of 1.1 arc - s/yr (Taylor et al., 1986; also see Figure 11-27). The onset of ray source measured by the satellite HEAO the radio outburst coincided with the observed A-2 on November 1977 at 2-6 keV (about 3 _) drop in visual light in July 1984. The radio light (Marshall et al., 1979). The doubt with this curve at 2-cm wavelength indicated a flux in- identification is that the other symbiotic stars crease by a factor of about 50 from April 1984 that have been detected in the X-ray range are to May 1985. The flux increased with increas- soft-X-ray emitters; none is known to emit in ing frequency, indicating a thermal origin. The the hard-X-ray range. expansional velocity was of the order of 2500 km/s, of the same order as the values given by VI.E. TOWARD A MODEL FOR CH CYG the full widths of Lyet and the Balmer lines at about the same epochs: 3950 km/s for Ly alpha (Seivelli and Hack, 1985), 1200 km/s for H The long and homogeneous series of spec- alpha and 1100 for H beta (Hack et al. 1986). troscopic observations made from 1965 to 1986 at the Haute Provence Observatory by Hack The radio jet from CH Cyg is an unusual and and collaborators, together with those made by complex event. The only other symbiotic star Deutsch et al. (1974) since 1961, suggest that the known to show a jet-like feature is R Aqr: it was radial velocity variations in the photospheric observed both at optical and radio wavelengths. lines of the M6 III star are due to orbital motion on Also in the case of CH Cyg, there is evidence of which are superposed some erratic variations the presence of a jet emitting in the optical and- commonly observed in giants and supergiants. possibly- in the UV. Solf (1987) obtained high- The presence of a companion seems to be proved resolution spectra of CH Cyg on September i 986. by the behavior of the permitted and forbidden His data reveal a very compact nebulosity located emissions, showing radial velocity varying in an- about 1" northwest of the star and emitting in the tiphase with that of the M6 photospheric lines. light of [O III] at 5007 A, i.e., in the same direc- The evidence for the presence of this companion tion as the radio jet. An attempt to observe it in the is reinforced by the minimum in the U magnitude UV with the IUE satellite was made by Selvelli et and in the UV flux observed in 1969 and in 1985, al. (1987, IAU Circ. 4491) in November 1987, suggesting the occurrence of an eclipse of the when the 20" slit was oriented in the same direc- companion by the cool star. Moreover, the detec- tion of the jet. By placing the star at one end of the tion of a slight increase of the flux toward slit, a stellar spectrum with P Cyg features was wavelengths shorter than 1300 A observed in observed at one side, while at the other (at about 1988 suggest that the accretion disk has be- 19 arcsec from the star), few emission features come very thin, and we can observe the con- were observed (SilII 1892, NIII 1750, OIII tinuous spectrum of a faint hot companion. 1663). An attempt to observe the jet again in May 1988, when the slit had the same orientation, gave A period of 16 years is not in disagreement negative results. with the values indicated by the radial velocity curve. However, the large scatter of the radial ve- Attempts to detect X-ray emission from CH locity data do not permit one to derive reliable Cygni were made with the X-ray satellite EIN- parameters for the orbit, but just to estimate a STEIN with negative results. The european satel- mass ratio of about unity. lite EXOSAT observed again CH Cygni on May 24, 1985 (Leahy and Taylor, 1987), and at this The observed flickering, with time scales of a time soft-X-ray flux was detected of 1.3xi0 tt erg few minutes, is a typical phenomenon observed in cm-2 s-L" dwarf novae and quiescent novae and believed to arise in the hot spot where the mass flux from the Actually, there is some suspicion that CH cool star impinges on the accretion disk. The Cygni is the optical counterpart of a hard X-ray presence of flickering in CH Cyg and its higher

723 amplitudein theU magnitudeareadditional Thelargewidthsof theLyAlpha(about4000 proofsthatCHCygisabinarysystemandthatthe krn/s),H alpha(about1200km/s)andH beta flickeringoccursin thehotcomponentof the (about1100km/s)atthebeginningof 1985sug- system. gestthatthediskbecomeslargerbyexpanding, andthishastheeffectof increasingtheRVgra- dientin thediskand,hence,thebroadeningof The ultravioletcontinuum may be ex- thestronglines,andof decreasingthedensity. plainedassuperpositionof astellarcontinuum TheEXOSATdetectionof X-rayemission, withT eff rangingbetween8500K (atmaxi- while previousobservationswith EINSTEIN mumUVbrightness)and15,000K (attheend gavenegativeresults,mayindicatethat the of theoutburst)andb-f+f-fhydrogenemission vanishingof theouterpartsofthediskmakesit (Mikolajewskaet al., 1987;1988). possibletoobservetheinnerhotterpartsand,in 1988,alsotheRayleigh-Jeanstail of thecom- Howevertheindicationobtainedin1988ofthe panion. presenceofaRayleigh-Jeanstailat_,_.1200-1300 suggestthepresenceof ahotterobject. Thedurationof anoutburstandtheintervals betweentwo consecutiveoutburstsmayhave Assumingasreasonablevaluesforthemasses very differentlengthsdependingontheellip- m(M6)=m(comp.)= 1solarmass,andP= 15.7 ticityoftheorbit.If thecriticalmassforoutburst yrs,theresultingdistanceofthetwostarsisofthe inthediskisreachedinthevicinityofperiastron, orderof 7.8A.U.or1.2x 10_4cm,i.e.,about10 theactivitymaybelongerthanif it isreachednear timestheradiusoftheredgiant.Hence,thesys- apoastron,becauseofalargeraccretionofmatter temisdetached,andif whatweobserveduring throughthestellarwind.Theintervalsbetween theoutburstisthe\spectrumofanaccretiondisk, thetwooutburstsmaybeshorterif, attheendof itmustbeformedbyaccumulationofmatterfrom oneoutburst,thestarisneartheperiastronand theredgiantwind.Duringtherisingpartof the maymorerapidlyreplenishthediskorlongerif, outburstthediskbecomesthickerandmoreex- attheendofoneoutburst,thestarisattheapoas- tendedasindicated:a)bytheincreasingintensity tron. oftheblueandUVcontinuum;b)bytheappear- anceofabsorptionlinesof once-ionizedmetals, Althoughtheorbitalparametersderivedby andalsoof multi-ionizedatoms(e.g.,C IV, Si YamashitaandMaehara(1979)areveryuncer- IV); c) by theincreasingwidthof theBalmer tainbecauseof thescatteroftheobservedradial lines.Attheendoftheoutburst,thediskbecomes velocities,wehavecomputedthephasesof the lessdense,asindicatedbytheratioof forbidden epochsofthethreeobservedoutbursts.Theout- to permittedlinesof FeII, theappearanceof C burstof 1963occurredatphase0.7(countedfrom III] 1906andof[OIII] 4959and5007,andbythe theepochofperiastron),andtheoutburstof 1967- diminutionordisappearanceoftheUV-bluecon- 70startedatphase0.94.The1977outburstoc- tinuum.Anindicationofthedecreaseindensityis curredatphase0.58;it wasnearto theendin alsogivenby thetransitionof UVabsorption- January1985,atphase0.05.Hence,theshorter dominatedspectrum(becausetheUVcontinuum timeintervalbetweenthefirstandthesecondout- isstrong,thediskisopticallythickin theUV) to bursts(4years)andthelongerintervalbetween an emission-dominatedUV spectrum(because theendof thesecondandthebeginningof the thediskbecomeopticallythininthecontinuum). thirdoutburst(7 years)maybejustifiedbythe TheappearanceofLyAlphaemissioninJanuary abovehypothesis.However,it isnoteasytojus- 85probablyhasthesameorigin.TheLy Alpha tify thestrengthandlengthofthethirdoutburst, absorptionwasnotobservableatearlierdates, startedwhenthecompanionwasnearto the probablybecausethecontinuumis verylowat apoastron. 1215AonaccountofthelowsensitivityofIUEat thatwavelengthandbecauseof interstellarab- Tohaveabetterexplanationfortheoriginof sorption. theoutburst,weneedmuchlongerseriesofradial 724 Gm velocitiesfor computingmorereliableorbital L= -_-- --dm/dt, 13.3 data,and,hopefully,newoutburstobserva- 2(V2e.+ V2,-)3t2Vo,,rt2a-" tions.

ThemodeladoptedbyWarner(1972)toex- where Vr__is the velocity of the companion mov- plaintheluminosityofMiraCetiB asduetoac- ing at orbital velocity, relative to the wind of the cretionfromthewindofMiraCetiA canbeap- companion, V is the sound velocity in the out- pliedto CHCyg.Assumingm_= m2= I solar flowing envelope, V _ I km/s, m_ is the mass, and masses,P = 15.74yrs,thesemiaxisa of the r__the radius of the companion. orbitisequalto 7.75A.U.or 1.1610_4cm.For m_=1andm2= 0.25(assuggestedbyMikolaj- For m2 = 1, r2= 0. I (both in solar units), V, = ewskietal., 1987),a = 6.6A.U. 8km/s, Vre_= Voo, and dm/dt = 10.5 m o/yr, it fol- lows L = 9.6 103_ erg/s. For m 2 = 0.25, Vorb =32 Themasslostbytheprimaryisgivenbydm/dt km/s, Vre_= 59 km/s, assuming r 2 = 0.01, it fol- =4rta2pV o,whereVoutistheexpansionalve- lows L = 1.75 1033erg/s. locityobservedduringtheoutburst.Foranob- servedparticledensityof l06(asindicatedby Now, assuming that the flux due to the M6 theforbiddenemissionsappearingattheendof giant is completely negligible below 1600 A, we theoutburst)orof 109(asindicatedbythefor- observe that the flux at the of the compan- biddenandpermittedemissionsof Fe I1 at ion in the period of maximum activity in the inter- maximumoutburst)andfor V u,= 50 km/s,it val of maximum emission, 1200-1600 A, is F = follows1.510_ m° /yr < dm/dt < 2.4 10.5me / 4 10-9 erg cm 2 s-_. Hence, the luminosity L, as- yr. suming for the distance of CH Cyg about 250 or 300 , results equal to about IW a erg/s. Hence, a mass loss of 10.5 m /yr, as indicated The radio structure observed by Taylor et al. ® (1986) from April 84 to January 85 indicates dm/ by the optical and radio observations, is suffi- dt = 7 10.6m o /yr. The luminosity, due to accre- cient to explain the observed luminosity of the tion only, of the object which is gaining mass, companion of mass about 1 and is given by radius about 1/10 the solar radius.

725

14

SUMMARY OF OUR PRESENT KNOWLEDGE ABOUT SYMBIOTIC STARS

M. Friedjung and R. Viotti

At the end of these Chapters 11, 12, and 13, observations can be interpreted if there is a main where the different observational and theoreti- sequence accretor surrounded by a disk. How- cal aspects of the symbiotic stars have been dis- ever, such a situation does not explain what is cussed, with special attention to a few well- seen for many (probably most) S-type symbiot- studied objects, it is now necessary to clarify to ics. For these, wind accretion on to a while dwarf the reader the present status of our research, appears more probable. In the latter cases a high and to find out the lines for future work on the mass-loss rate from the cool giant might be sug- field. In the following, we shall summarize the gested; actually, Kenyon and Fernandez-Castro main points concerning the symbiotic phe- (1987), Kenyon et al. ( 1988), and Kenyon (1988) nomenon. found evidence of enhanced mass-loss rates com- pared with those of normal cool giants of the same I. SYMBIOTIC STARS AS INTERACTIVE types. The last conclusion is clearly dependent on BINARIES the accuracy of the spectral classification, and especially on mass-loss rates determined from continuum radio emission and masses of dust There is now strong evidence that most, if not present. Indeed, the accuracy of mass-loss rates is all, symbiotic stars are interactive binary systems. not so good: radio emission, in particular, needs Indeed, this conclusion is so much believed by to be modeled taking into account ionization of specialists, that if a star previously classified as the cool-giant wind by radiation from the hot symbiotic had been found not to be binary it component (see Chapter 12). It can be noted that would then have been classified as something the mass-loss rate for the D-type symbiotic H 1-36 else! One component is thought to accrete from found from the model of Taylor and Seaquisl the other. One, the mass loser, is supposed to be a (1984), which is obtained by assuming the same cool giant, while the mass gainer should be a wind velocity as those authors give, but with the main sequence star or a white dwarf or possibly a distance from Kenyon el al. (1988), is two to eight neutron star. Accretion is supposed to occur ei- times larger than that derived by Kenyon et al. if ther from Roche lobe overflow via a disk, or one assumes spherical symmetry. from the wind of the cool component. When the accretor is a white dwarf, the accreted hydro- gen can be burned almost continuously or in For the D-type symbiotics such as RR Tel and shell flashes. However, even accepting this V1016 Cyg, it has not been possible to derive an basic picture, many "details" are not under- orbital period, although there are many reasons to stood, while it has to be related to the behavior state that they also are binary (e.g., from energy of particular symbiotic stars. balance considerations). Their periods are often believed to be much larger than those of S-type In the case of S-type symbiotic systems, the systems (10 t to 10_ years), thus implying larger evidence for binarity has grown. The orbital peri- separations and less interaction phenomena un- ods seem to be of the order of several hundred less the components are more "active." In fact, days. In some cases (e.g., CI Cyg and AR Pav), the cool giant in such systems appears to be a

727

PREGEDING PAGE BLANK NOT FILMED Mira variable, having a larger mass-loss rate needs to be checked in much more detail in the than that of other cool giants. It appears that it future. The claim of an increased mass-loss rate is the presence of a Mira which explains the for the cool components of S-type symbiotics different properties of such systems, and it now should be verified using proper seems best to call them symbiotic Miras modeling for infrared dust and radio emission. (Whitelock, 1988). Their large mass-loss rate The atmospheric structure needs to be studied is considered to be associated with the conden- using high resolution infrared spectra (e.g., sation of a large amount of dust. Indeed, the from Fourier Transform Spectroscopy), which infrared properties of symbiotic Miras might be can also be used to measure stellar surface understood by supposing a much larger extinc- gravity, turbulence, and rotation velocity. tion of the light of the Mira than that of its companion, as the result of the presence of a It appears, however, that binarity is the main massive circumstellar dust envelope (Kenyon cause of the properties of symbiotic stars; it now et al., 1988). is very likely that they can be ascribed to the pe- culiar properties of the cool component. Indeed, When speaking of interaction in symbiotic bi- the normality of this component is largely used to naries, we need to consider quite a number of dif- derive the distances of symbiotic systems. These ferent processes. These include (I) mass ex- distances need to be refined, but almost certainly change following mass loss from one component will not be substantially changed in the future. leading to transfer and accretion by the compan- ion, perhaps modulated by varying stellar separa- The exact nature of the hot component is still tion during an orbital period; (2) possible nuclear more open and controversial. In some cases, it burning of hydrogen accreted by a white dwarf; appears to be like a hot subdwarf, and can be (3) interaction of radiation of one component understood as being an expanded white dwarf with the environment of the other; (4) collision undergoing shell burning. Then it often is similar between the winds from each component, etc. to the nucleus of a planetary nebula, and tempera- The interaction of radiation of one component ture and radius estimates are derived as for the with the environment of the other can take many latter from the ultraviolet energy distribution and forms, including heating of a disk and/or of a emission line fluxes. The temperature can be so companion's atmosphere, ionization of the wind high that the observed ultraviolet continuum is and upper atmosphere of the cool component (this quite insensitive to its exact value. The properties can well vary with time), heating and possible of certain hot components during outbursts (e.g., destruction of grains, and acceleration of the cool CH Cyg and PU Vul) somewhat resembling a star's wind by in the lines, etc. supergiant of intermediate temperature, might Up to now, studies have concentrated on one then be explained by expansion of the outer lay- process at a time, so that a proper global picture ers of the white dwarf to supergiant dimensions, of interaction still does not exist. following a shell flash. However, in this case, the luminosity would have to be - 59250 II. NATURE OF THE COMPONENTS OF (Mcore/M o -0.522) L o (Kenyon, 1986). Ac- THE SYMBIOTIC SYSTEMS cording to Mikolajewska et al. (1988), CH Cyg did not have a maximum luminosity of more

To make consistent models of the phenom- than 10-_ L o . Therefore, in order for this expla- ena here studied, we especially need to know nation to work for CH Cyg, the white dwarf the nature of the components. It is now clear core mass would have to be quite low, of the that the cool component very much resembles a order of 0.54 M o . normal cool giant or bright giant, for S-type symbiotics, and a Mira for D-type symbiotics. Here by "normal" we mean that the cool com- The hot component can sometimes be under- ponent appears very much to resemble the stood as an accretion disk plus boundary layer corresponding type of single star. This result around a main sequence star. Such a disk may be

728 formedfromRochelobe overflow and sometimes formation near the hot component. In fact, the from wind accretion, but in the latter case, a disk situation is more complex, as material does not does not seem easy to be formed. In any case, the appear to be distributed uniformly. The fairly detailed interpretation of the hot component as a narrow emission lines seen very often are most disk still involves a number of uncertainties be- easily understood as formed in regions con- cause of the gaps in the present theory of disks. nected with the cool component (wind, outer Studies assuming a disk to radiate as a sum of atmosphere). Line narrowness implies forma- blackbodies or even as a sum of normal stellar tion in a low-velocity region far from any atmospheres cannot lead to highly reliable re- compact accreting object. In addition, a study sults. Other physical effects can also be present. of inter-combination lines of Z And and RR Tel Among these, one should mention heating of the indicated line formation where radiation com- disk by a central stellar component, which can be ing from the hot component at 1176 and 772 A an expanded white dwarf undergoing shell burn- (Altamore et al., 1981) was diluted. ing. The heated disk could reradiate as a result of the heating at a rate much larger than that due to A similar kind of situation is believed to exist gravitational dissipation in the disk. If the hot for zeta Aur binaries where a hot main sequence component is a disk, one can envisage an explana- star appears to be immersed in the wind of a cool tion of active phases of symbiotic stars, by supergiant companion (Shroeder, 1988). The mechanisms similar to those invoked for dwarf wind scatters photons from the hot main se- nova outbursts. In addition, it is not certain quence star in resonance lines according to suc- whether a disk, if formed, would always be pres- cessful models for such stars, and information on ent. For instance, it might temporarily appear dur- wind velocities and mass-loss rates can be ob- ing activity, if it were formed by accretion from tained. Such stars provide lessons for the study of the wind of the companion. The properties of the symbiotic systems. latter or the separation of the stellar components would have to change. Symbiotic stars, however, are more complex. As discussed in Chapter 12, not only is the ge- The presence of a disk should lead to observa- ometry of regions of the cool star wind ionized by tional consequences other than those exhibited if the hot component complex, but other physical only an object similar to a hot star were present. processes determine line formation regions also. Not only is the continuum energy distribution Evidence of a high-velocity wind from the hot changed, but also the emisssion and absorption component, in particular, can be seen for AG Peg, line profiles (or at least contributions to the pro- this producing the broad component of line pro- files), as well as all the variations in eclipse, can files. Collision between the two winds, when be expected to have characteristic properties. important, can also lead to another line forma- High-quality observations, combined with better tion region, whose properties really have not theory, should help to eliminate these problems. been studied up to now. Emission-line forma- In any case, the correct identification of the na- tion can occur near an accretion disk (plus per- ture of the hot component for a particular system haps a bright spot formed where a current from is important for deducing which processes domi- a companion losing mass by Roche lobe over- nate. in addition, such information is essential for flow strikes the disk). It may be noted that such understanding the system's evolution. a disk might also produce a wind. The influ- ences of these and other effects on emission- 11I. EMISSION LINES line formation still need to be fully elucidated.

The picture just described obviously has a lV. CHEMICAL COMPOSITION bearing on the formation of the emission lines. If formed by electron collisions or by cascade Some of the models discussed above imply following recombination in an ionized region, that the matter in the symbiotic system should be one might, in the simplest situation, expect line processed. Abundances different from cosmic

729 valuesarealsoexpectedforthehigh-velocityob- close to those for M giants, suggesting that the jects,suchasAGDra.Therefore,thedetermina- line-emitting material come originally from tionofthechemicalcompositionof allthecom- the cool-giant companion. The latter would ponentsof thesymbioticsystemsisacrucialpa- have its abundances somewhat modified by rameter.Obviously,nothingcanbedirectlysaid CNO cycling. The only symbiotic system that aboutthehotcomponent,sincenotrulyphoto- shows clear signs of deviation is HM Sge, for sphericlinesarevisible.Thespectrumofthecool which C/N/O ratios were found to be similar to componentcanbe studiedwith theclassical that of novae. It would be important to extend curve-of-growthmethod,providedthathigh- these results to other ions, and to check to what resolution,high-S/Nspectrogramsareavailable, extent the abundance estimates are dependent whichisnotthecasefor thelargemajorityof on the assumptions and on the adopted model. symbioticstars.In addition,in orderto avoid Future work on both (and simultaneously) the errorsintroducedbytheveilingof thevariable emission-and absorption-line spectra using blue continuum(see Chapter11, Section high quality material are urgently required to IV.A), theanalysisshouldbe madeon high- make any progress in this field. resolutionspectrogramstakenin thered,or evenin thenear-IR,whichis notsoeasyat V. VARIABILITY present.Sofar,theabundanceanalysesof the brightestobjectsareveryfew.Froma curve- of-growthanalysisof theopticalspectrumof The variability of symbiotic stars is a funda- thesymbioticnovaPUVul, Belyakinaet al. mental property of these objects. It is the result of (1984)foundsomechemicalanomalies,such many mechanisms, which in many cases are far asCaandFedeficiency,andexcessof theother from being well-understood. Various time scales irongroupelementsandof somerareearths(cf. are involved, and we mention here (nearly in GershbergandShakhovskoj,1988).Lutzet al. order of increasing time scale): (1987)foundthatin AGDra,theBail andSrlI linesareprobablyenhanced.Unfortunately, - The flickering of CH Cyg with time scales of theseauthorswereunabletoperformadetailed 5 and 15-20 min seen during activity, which abundanceanalysisfor lackof a goodcalibra- might be physically related to the flickering seen tion of theirechellespectra. for cataclysmic variables.

- Variations during the orbital cycle, which Concerningthe "nebular"emission-line may not be only geometrical (eclipses, reflection spectrum,the abundancedeterminationsare effects, etc.), but also physical, associated with a stronglymodel-dependent.Severalestimates varying separation of the stellar components. In havebeenmadebasedontheopticalandUV the latter case, the orbit must be eccentric'. It linefluxes.Theresultsof CNOabundancede- should be noted that accretion variations due to terminationsfrom IUE are summarized by varying separation cannot have much effect on Nussbaumer et al. (1988), who used the emis- the brightnesses of AG Dra, AX Per and AG Peg; sion-line fluxes of CIII, CIV, NIII, NIV and otherwise, the determination of orbital ele- OIII, supposed to be formed in a common re- ments from the reflection effect as discussed by gion. They also supposed the ionizing hot Leibowitz and Formiggini (1988) would not source to have an effective temperature equal work. to or larger than 105K, and the nebular re- gions to be uniform with an electron tempera- - When the cool companion is a Mira, vari- ture of 12,000 K and an electron density of 10+_ ations occur over its pulsationai period, which cm 3. Line emissivities were found to be insen- are of the same order as the orbital period, if it sitive to the assumed electron temperature, as is an S-type system. Accretion and dust con- well as to the assumed electron density, at least densation might sometimes be modulated. for N e below 10÷9 cm -3. Nussbaumer et al. found that in symbiotic stars the abundance ratios are - Active phases can occur over time-scales of

730 decades.Withineach,oscillationsofactivitycan betweenejectaandstellarwinds.Theformer occur.Some,butprobablynotall,of theactive areclumpyandassociatedwith a knownout- phases,mightbeexplainablebyaccretionevents. burstof thesystem,whilethelatteraresmooth Theactivephasesofothersymbiotics(e.g.,AG andfeaturelesswith an angularsizethatin- Dra)maybehardtounderstandwithoutinvoking creaseswithfrequency.It is theejectafollow- shellburningof awhitedwarf. ing outburststhatshowbipolarflow (or jet- like) structures. - D-typesymbioticscanhavefaintphaseslast- ingonetoseveralyears.Thesearepossiblydue Thephysicsbehindthe originof nebular todustobscuration,and/ortophenomenaassoci- structureis notreallyknown.Forwind-pro- atedwiththeperiastronpassage. ducednebulae,theionizationof thecoolcom- ponent'swindby thehotcomponentandcolli- - Onlyone outbursthasbeenobservedfor sionbetweenwindsfrombothcomponentsmay eachsymbioticnova.Sucheventshavebeenex- playmajorroles.A systemseenin theplaneof plainedbyshellt]ashesofwhitedwarfs. itsorbitcouldshowapparentlylinearstructure in certaincasesbecauseof thesemechanisms. However,theexplanationsinvokedfordiffer- However+suchan explanationis notexpected enteventsarenottobebelieveddogmatically. to begenerallytrue.Bipolarflowsandjetsare Certaindividinglinesbetweendifferentclasses commonin otherastrophysicalsituationssuch ofeventmayturnouttobeartificial. asradiogalaxies,activegalacticnuclei,young stellarobjects,etc.Their existencemaybe VI. NEBULAE linkedto theexistenceof disks,butit wouldbe dangerousto extrapolatethistypeof "'explana- Smallnebulaehavebeendiscoveredaround tion" to symbiotic systems, and to conclude manysymbiotics,especiallyat radiowave- that disks are very often present. Conversely, lengths.Imagedeviationsfromcircular(there- the study of bipolar nebulae and ejecta in forefrom spherical)symmetryareobserved, symbiotic systems might be useful to understand with indicationsof the presenceof bipolar their nature. A large progress in this field is ex- flow'sandjets (Taylor,1988,Solf, 1988). pected from the new astronomical technologies Accordingto Taylor,oneshoulddistinguish for imagery and polarimetry+ and front HST.

731

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747

15

PERSPECTIVES AND UNSOLVED PROBLEMS

M. Friedjung, M. Hack, C. la Dous and R. Viotti

749 PRECEDING PAGE BLANK NOT FILMED In this bookwe reviewedtheobserva- spectra, in conjunction with the extremely tionsof dwarf novae, nova-like stars, novae, short orbital periods, suggest that these systems recurrent novae and symbiotic stars, and the consist of two white dwarfs and, therefore, that current state of their interpretation. We tried to their evolutionary history must be different demonstrate the immense variety and variabil- than that of the other systems which are be- ity of phenomena found in these objects. lieved to consist of a white dwarf and a . As is probably true in all other fields of sci- ence, we are facing a dilemma. On one hand the Although nova-like stars are commonly be- availability of new observational material- from lieved to be dwarf novae in a permanent state of satellite telescopes as well as from ever more ad- outburst, statistical differences in the photom- vanced and sophisticated ground-based devices - etric and spectroscopic appearance of members is opening our eyes to ever new phenomena in of both classes suggest that the physical differ- cataclysmic variables which help clarifying some ences between them might lie beyond mere questions. On the other hand, however, unex- outburst behavior. pected new questions and problems arise. The general picture becomes clearer; for instance In spite of the very large number of observa- there is little doubt left that in principle the Ro- tions of dwarf novae, the very start of rise to an che model is rather well suited for explaining outburst, due to the unpredictability of the the basic physics of cataclysmic variables event, has so far escaped detailed observations. (with the possible exception of symbiotic However, this particular phase in the activity stars), but curious assumptions and concepts cycle of dwarf novae is very likely to contain flourish if details of individual observations valuable clues to the structure and dynamics of are to be explained. And, not surprisingly, the accretion discs. Thus, what is called for in conceptual understanding has evolved much order to help the situation is continuing moni- further than our ability to carry out detailed, toring of a few selected objects, both photom- meaningful computations. etrically and spectroscopically, in as large a spectral range and over as long a time interval Let us consider dwarf novae and nova-like as possible. stars. A few intriguing statistical differences between the different sub-classes of these sys- In both dwarf novae and nova-like systems tems seem to exist, with respect to the distribu- the companion stars are cool main-sequence tion of orbital periods and the masses of the stars which in all probability undergo solar- stellar components. Furthermore, judging from type activity cycles. These activities, in turn, their spectroscopic and photometric appear- are likely to influence the brightness variabil- ance, nova-like stars belonging to the sub-class ity of cataclysmic variables. Moreover, accord- of UX Ursae Majoris stars, anti-dwarf novae ing to theoretical considerations, the secondary and dwarf novae appear to be basically the stars are forced to co-rotate with the binary same kind of objects. In most respects also the orbit at a considerably higher velocity than DQ Herculis stars are very similar to them. normal for stars of this spectral type. It is not They are suspected, however, to possess a known yet what effect this has on the star, nor moderately strong magnetic white dwarf. And is it known what effect it must have for it to be finally, the appearance and behavior of AM confined to the non-spherical shape of the Herculis stars can be understood, if it is as- Roche lobe. All these problems should be taken sumed that they also are basically the same into account by a new generation of theoretical kind of objects, but that their white dwarfs models. possess a very strong magnetic field. AM Ca- num Venaticorum stars, on the other hand, A weak point in the computation of models is represent a different kind of system. The com- also the structure of the accretion discs. So far plete absence of hydrogen lines from their only very simplified models have been consid-

750 ered.Two-dimensionalhydrodynamicscompu- observe that only 1/1000 binary systems are tationsdependstronglyontheassumptionsabout formed of a close white dwarf plus a red dwarf, theviscosity.Moreover,theverticalstratifica- able to produce a nova outburst, the second possi- tionhasbeenincludedinthecomputationsonlyin bility is the only one which can be accepted. The a rathercrudeway,yieldingcorrespondingly hibernation theory follows, according to which vagueresults.Afurthercomplicationispresented each nova will erupt thousand of times and stay in bythehotspot.Thegeometricalstructureandthe a low state for centuries between one outburst and positionofit in thedisc,whichareatheoretical the following one. Although many details are not ratherthanacontroversialissue,areboundtoin- completely explained, the hibernation theory fluencethestructureanddynamicsof thedisc. gives a unifying picture of dwarf novae, nova-like Similarly,spectrumcomputations,too,severely stars and novae, explaining the change from one sufferfromthepoorlyknownphysicalstructure class of cataclysmic variables to another with the of theaccretiondiscs. change of mass transfer over the millennia.

Classicalandrecurrentnovae,bothin quies- Another question is why novae are so similar to each other at minimum, and so different from cenceandin outburst,presentplentyof open each other in outburst? Actually it is difficult to problems. We summarize here some of them, and say if an old nova behaved as a fast nova or a slow will try to indicate which course to pursue for nova, just looking at its characteristics in quies- tackling them. cence.

One main question is: are dwarf novae, Fundamental parameters like the masses of nova-like stars and old novae the same class of the two components, or even as basic as the objects, just seen in different stages of activity? orbital periods, of a cataclysmic variable sys- Actually the only known statistical difference tem are badly known. There is a need for simul- among these three groups, besides the outburst taneous as well as long-term observations of activity, is their : the mag- light curves and radial velocity curves, so far nitudes at minimum of old classical novae available for just a few individuals. cluster around +4.5, those of quiescent dwarf novae around +7, and the nova-like stars have Still, careful analysis of observations indicates magnitudes clustering around +5. This differ- that the problem of deriving system parameters is ence might be due to a weakening of the accre- not as straightforward as one might naively as- tion disc in very old novae. Furthermore, there sume. Besides the technical problems of suitable is some indication that very old novae, like CK data analysis, such as properly taking into ac- Vul, are several magnitudes fainter than old count irradiation effects in the system, some dif- novae just a few decades after the outburst. ficulties seem to be even more basic, reflecting Also two of them, CK Vul, and WY Sge, are that the underlying physics probably is not quite reminiscent in their photometric behavior of as simple as would be desirable. For instance, in dwarf novae. Both these observations can be some systems the photometric periods seem to be taken as a sugestion that dwarf novae merely variable on unreasonably short time-scales. In are very old novae. It is not clear, however, other systems the spectroscopic and photometric where the apparently closely related nova-like periods are different from each other. variables fit into this picture. It is very important to get better determina- Actually we observe 50 novae per year in M tions of the chemical composition of the ejecta of 31. In 15 billion years there should occur novae, both for obtaining information on the 7.5x10 _t novae. Hence we have two possibilities. phenomena occurring at the surface of the hot de- Either practically all stars in M 31 will become or generate star when it accretes matter from the have been novae, or each nova must sufferfout- boundary layer of the accretion disc, and also for bursts (withfthe ratio of the total number of stars obtaining information on the characteristics of in M 31 to the number of nova systems). Since we the white dwarf itself.

751 Theimpactof ultraviolet observations on our stars, and finally, T CrB and RS Oph have a red knowledge of cataclysmic variables has been giant in their system (rather than a red dwarf as very important. Not only because it has given us the other cataclysmic variables). T Pyx, unlike the possibility to obtain abundance determina- other recurrent novae, presents a very steep rise of tions from ions not observable in the optical the ultraviolet flux toward the shortest wave- range, but also for having given us proofs of the lengths accessible to IUE. reality of the existence of accretion discs from the

shape of the continua, which are generally fitted Also, the distinction between novae and sym- by a power law. Only in few cases can a Rayleigh- biotic stars is not as neat as one would like it to be. Jeans tail, imputable to the hot member, be de- Symbiotic stars are known, like V 1016 Cyg and tected. The importance of extending the observa- V 1329 Cyg, which behave much like classical tions to shorter wavelengths than those accessible novae. On the other hand RT Ser and RR Tel, both to IUE is clear. However the hope of detecting the commonly classified as old novae, behave spectrum of the hot companion at shorter wave- much like symbiotic stars. lengths than those accessible with IUE will not necessarily be satisfied. In fact the few data from Since the basic structures of novae and sym- Voyager actually indicate that the same power biotic stars might be rather different from each law explaining the IUE range is not valid for the other, but since recurrent novae bear character- shorter wavelengths, but neither is it explicable istics of both, their investigation might yield by a hot black body emission. The spectrum in- valuable clues to the nature of both, old novae stead is generally very fiat. and symbiotic stars.

Few nova outbursts have been observed in the Because of the length of the orbital periods near infrared. From these data it seems that a dif- of those symbiotics that have been ascertained ferent behavior characterizes the fast novae (no to be binary systems, the interaction between appreciable formation of a dust shell), the inter- the cool giant and its companion must occur, in mediate novae (formation of an optically thin the majority of cases, by accretion through dust shell) and the slow novae (formation of a winds rather than by overflow of the Roche thick dust shell, although the few data for two lobe. It is possible that many characteristics of very slow novae indicate absence of a dust shell). the outbursts are affected by the properties of It is not yet clear how general these relations are. the , by their eccentricity, etc. It should A larger sample is certainly needed, in order to be, therefore, necessary to obtain better deter- firmly support the notion that these behaviors re- minations of the radial velocity curve, over ally are correlated to the speed class of a nova. long periods of time, covering more than one The mechanism of dust formation is not clear ei- orbital period. In fact one main difficulty in ther. Dust may be preexisting to the outburst or it obtaining reliable radial velocity curves is that may form in the envelope when it cools off. the amplitudes are generally small and compa- rable with the irregular fluctuations often pres- ent in the atmospheres of cool giants and super- The foremost problem with recurrent novae is giants. one of classification, thus one of what one imag- ines their basic structure to be. The problems of line blending of the wide Doppler broadened lines mean that outburst spec- tra in different spectral regions must be studied It is not obvious that they can be addressed as using spectral synthesis. The use of spectral syn- a reasonably homogeneous class at all. Only five thesis requires the finding of parameters which individuals (*) are known which can be divided are sensitive to unknown quantities connected into at least three sub-classes: T Pyx and U Sco (*) A sixth recurrent nova has been recently discovered: are similar to classical novae, V 1017 Sgr has a V394 CrA 1949 which had a second outburst in 1987. it behavior more reminiscent of that of symbiotic behaves similarly to T Pyx and U Sco. (Liller, 1987)

752 with the model, abundances, etc. A vague re- Observations of high spatial resolution in vari- semblance of the results of a calculation to ous spectral bands can give more information what is observed is not sufficient. about the geometry of the ejected envelope, in- cluding deviations from spherical symmetry. It Besides better observations we certainly need will be particularly interesting to resolve the en- better and more realistic models. velope as early as possible after the outburst of a nova and to follow the evolution of its structure Parameters which affect the absolute magni- with time. Study of the spectrum of different parts tudes at maximum, the shape of the outburst light of the nebula will not only give information about curve, and the ejection velocities, etc. can be the differences in physical conditions, but also about mass and chemical composition of the white the velocity distribution in the line of sight dwarf, the mass of the accreted material and the (from line profiles), leading to the possibility accretion rate, the degree of mixing of white of the three-dimensional reconstruction of dwarf material into accreted hydrogen-rich shell envelope structure. material, the orbital elements and the inclination of the orbit, and the presence and strength of mag- It must be emphasized that the study of netic fields, etc. Unfortunately the majority of nebular expansion, combined with knowledge these data are only poorly known or altogether of the expansion velocity, is the best method unknown. Through extensive comparison be- for determining the distance of a nova. tween theoretical models and observations it The theory of classical-nova outbursts needs should be possible to get a better handle on physi- to be further developed taking into account de- cally reasonable parameter ranges for more so- viations from spherical symmetry and the in- phisticated models. Such new models then clearly should take into account non-LTE effects, cir- fluence of the companion. cumstellar absorption and dust, and the hydrody- In addition magnetic fields appear to be impor- namics and thermodynamics of shocks, etc., be- tant for some old novae. The old nova V 1500 Cyg fore spectral synthesis is attempted. seems to be a polar object and GK Per an interme- diate polar one. Strong fields should have a major As long as such a theory is not available or not effect on the outburst. usable, semi-empirical methods of analysis need to be further developed, among these the self-ab- The causes of recurrent-nova outbursts are still sorption curve method by Friedjung and Murato- not clear, and it is possible that some of the impor- rio (1987) for studying spectra of ions where there tant physical processes involved are completely are very many lines, such as Fe II, should be unknown. More work on accretion events and mentioned. In such methods the fit is made to as- thermonuclear runaways at very short intervals of sumed physical situations which, because of time may help to solve these problems. physical uncertainties involved, have more free parameters than a self-consistent theory. A very large amount of work needs to be done before one can claim to understand novae rea- Simultaneous or nearly simultaneous multi- sonably well. frequency observations are required for compari- son with spectral synthesis leading to determina- In the chapters devoted to the observation and tion of the total bolometric luminosity, the mass- modeling of symbiotic stars, attention has been loss rate and velocity distribution of the continu- focused on the symbiotic phenomenon, rather ously ejected wind, and the disc structure, etc. than on a group of stars. The reason was that (1) it is not clear whether symbiotic stars really Novae often go through post-maximum oscil- represent a well-defined category of stars, and lations; and to understand their nature, multi- (2) our main interests are the physical proc- frequency observations closely spaced in time esses occurring in the atmospheres of stars, are needed. rather than the nature of the stellar objects.

753 Actually,symbioticstarsrepresentonly a stars has been collected in many different fields, comparatively small number of objects, and there from radio to X-rays, which in principle could be should be no reason to dedicate so much time to very useful for a statistical approach. But some their study, unless we expect from it - as is in fact concern should be expressed on the criteria which the case- results of much wider interest. have been used in the selection of the targets in these surveys. It seems to us that these data have We have shown in the previous chapters that not yet been satisfactorily analyzed with the stan- symbiotic stars, as other cataclysmic variables, dard methods of statistical analysis. Thus new are characterized by non-LTE phenomena that surveys, and extensive use of statistical methods, have been observed in many different categories are aspects which require more work in the fu- of astrophysical objects. These phenomena in- ture. Selection effects will, however, have to be clude: stellar chromospheres, mass loss and stel- properly studied, and fully taken into account. lar winds, circumstellar nebulae, accretion phe- Another major problem is associated with the nomena in close binaries, superionization and X- ray emission, etc. These phenomena characterize techniques of determination of the physical para- several differents fields of astrophysics such as: meters from the observations. To develop an em- formation of planetary nebulae, cool-and hot-star pirical model, we need to use diagnostics of the environment of the symbiotic stars derived from winds, circumstellar dust, jets, maser sources, pulsation, X-ray sources, accretion disks, ther- observational data. For instance, radio observa- tions can provide information about ionized cir- monuclear runaway, and close binary evolu- tion. Therefore, symbiotic stars are linked with cumstellar envelopes, while the infrared rather gives information about the cool giant and cir- many astrophysical categories of stars, including M-giants, Miras, Zeta Aur/VV Cep stars, hot cumstellar dust. Emission-line ratios provide first-approximation estimates of electron densi- subdwarfs, planetary nebulae, and obviously close binary systems, besides active galactic nu- ties and electron temperatures in the line emitting regions. However, such regions are not uniform, clei too. This is why symbiotic stars have always which can explain differences in the parameters attracted the interest of so many investigators from different fields. Their hopes were, and are, determined using different emission-line ratios; moreover the physical constants involved in the to have a better insight into the physics of their calculations are not too certain. own objects from the study of the symbiotic phe- nomenon, where the same mechanisms are proba- To overcome these type of difficulties, we bly present but with a larger strength, or on a dif- need to calculate observable quantities from de- ferent spatial scale. tailed models. The latter must take into account many different physical effects, including ioni- On the other hand, we have found that the di- zation of the wind of one stellar component by agnostics of the symbiotic phenomenon are not the other component, collision between winds, very straightforward. First we need as complete etc. These calculations use many unknowns, an observational description of the phenomenon thus future work will have to involve conver- as possible. However, because of their long time- gence between models and observations. The scale variability, symbiotic stars must be studied ideas behind the former will have to be con- for several decades, before a reasonably complete fronted by the facts provided by the latter. picture can be drawn. In addition coordinated ob-

servations in a broad frequency range are re- In order to make progress from the observa- quired, since, as discussed in previous chapters, tional point of view we need better orbital data different spectral regions involve different physi- about the binary systems. This requires the accu- cal phenomena. In practice, this requirement was rate determination of the radial-velocity vari- (at least partly) fullfilled in only a few cases. ations of the cool component. High resolution infrared spectra will not only yield this, but will On the other hand, a large amount of observa- also provide information about the "normality" or tional data on very broad samples of symbiotic "non normality" of the cool component. In addi-

754 tion,it maybenotedthatin thisspectralregion exist, and they must be carefully examined, in onecandetectabsorptionbetweentheobserver particular for wind accretion. andthecoolcomponent,sotheenvelopecanbe studiedusingabsorptionlinesproducedbyit in It is only when ideas based on much more suchpositions. careful research are confronted with better ob-

Diagnostic methods need to be enlarged. For servations spanning many years that signifi- instance, symbiotic stars are rich of Fe II emission cant progress will be made about the symbiotic lines which have proved to be very useful for the phenomenon. investigation of emission line objects. But still a large amount of data on this and other ions have Finally, along which lines should further not yet been studied. It should also be emphasized reasearch be developed? We realize that new so- that the profiles of spectral lines formed by phisticated observations really raise more ques- circumstellar material, rather than their radial ve- tions than they yield answers. Each of us has locities, needs to be examined. The radial veloc- made the claim "We need more observations". ity, in fact, gives only a kind of"integrated" infor- We surely do. But we need certain particular mation about the region of line formation, while kinds of observations. We need observations in- the study of the line profile at high spectral reso- volving a somewhat different approach than lution is basic for a detailed modeling. Finally, that which has been useful and appropriate so those systems showing eclipses can give a lot of far. The traditional tendency to obtain bits and information about the geometry and physical pieces of observations on as many objects as conditions of different regions. possible needs to be supplemented. There are many obvious gaps in the records of observa- Observations of spatial structure at very high tion; many are even relatively easy to fill in. It resolution are essential. In this regard, we expect is sobering to realize how little is actually very exciting results from the Hubble Space Tele- known. For instance, we know little about scope as well as from the large ground-based tele- spectra of dwarf novae during the outburst scopes equipped with new sophisticated instru- state, or about how they develop on time-scales mentation. In particular, information about jets of minutes or hours (rather than days), or what and bipolar flows should be obtainable. Polariza- orbital photometric variations occur during out- tion studies will provide precious additional in- bursts, or whether they are of orbital (geomet- formation on the geometry. rical) origin or only occur on orbital time- scales. And we still don't know what triggers a On the theroretical side, much is still to be dwarf nova outburst, or what causes many understood about accretion discs in symbiotic nova-like stars to drop in brightness. We be- systems, and in particular those formed following lieve the mass-transfer rate is responsible, but wind accretion. The reprocessing of radiation it is not clear why it changes in the first place; from a central object by a disc needs to be treated i.e., what the trigger is. in a much more rigorous way than previously. Furthermore, we now know that whatever ob- Effects of symbiotic binarity on the cool-star servations of a system we obtain, these are not upper atmosphere and wind have not yet been likely to be entirely characteristic of the particu- fully studied. This aspect is obviously also linked lar brightness state we happened to monitor, but to the uncertainties one still has about single cool are likely to be merely one possible pattern. Ob- giants. Perhaps the symbiotic phenomenon will servations at the same brightness level well may be of help for a better understanding of the outer yield another result at some later time. Currently, atmosphere of normal cool stars. we have only glimpses of the actual life of indi- Colliding wind models need to be worked out vidual systems, without being able to appreciate in much more detail, in order, for instance, to bet- how the most intricate interactions of all kinds of ter predict emission-line fluxes. Other situ- phenomena finally result in, for instance, the ations where shock formation should occur also violent outburst events.

755 Sowhatactullayis calledfor is long-term tionsof thissortarestill largelylacking,l_rom semi-continuousmonitoring,in asmanywave- comparisonof featuresof such statistical lengthrangesandbyasmanymeansaspossible. sampleswiththetheoreticalresultsobtainedfrom Thishasbeentriedfor veryfewobjects,with variationsof parametersoverwidelimits(after ratheramazing,andquiteunexpectedresults. parametersof onlysecondaryimportancehave Campaignsliketheseneedto berepeated,and beenruledoutbeforehand)hopefullypossible observingplansandschedulesof thevarious valuescanbelimitedtoreasonablynarrowranges workinggroupsshouldbecoordinatedfor ob- thatthenmightyieldrealisticresults;whilemost servingtimesandfor objectsto bemonitored. theoreticallypossiblevaluescanthenberuled Theprioritiesofcommitteesfortheallocationof out.Oncesuchnarrowerlimitswillhavebeenes- telescopetimeneedtobechangedsothattheyfi- tablished,thereishopethatcomputationscould nallyrealizethatthereisafterallalotofpotential beusedtoactuallydeterminesystemicparame- scientificmeritin re-observingthesameobject tersfromcomparisonwithobservations. overandoveragain.Resultsshouldbecommuni- catedquickly,ratherthanaccumulatingdustin Onefurtheraspectof researchin cataclysmic officesandarchives.Alreadyexistingmaterial, variables(butthisisvalidalsoforseveralother belongingtoindividualsaswellashiddenaway fieldsofresearch)shouldbechanged,soworkin in archives,needsto bemadewidelyavailable lateryearscanremainfruitful.Thenumberand andstudied.And,lastbutnotleast,thecount- qualityof publicationsshould,in thefirst case lessamateurastronomerscouldandshouldbe decrease,in thesecondcaseincrease.Overthe involvedin systematicmonitoringactivities. last20 yearsthenumberof papers,onlyre- Theyhavethetelescopetimeandtheenthusi- strictedtothenarrowfieldof cataclysmicvari- asmto observe,bothof whichmanyprofes- ablesandrelatedstars,hasincreasedalmost sionalastronomersarelacking. exponentially,sothatprobablyalreadybynow astagehasbeenreachedwhereit ispractically Astronomicalsatellitesdedicatedtothemoni- impossibleto readandfully assimilateall that toringof stellarvariabilityfromX-raytooptical is beingpublished.Manyfractionalresultsare wavelengthsarenowbeingprojected,andwe beingturnedout asindividualpublications, hopetheywill soonberealized. ratherthanastheresultsof completeinvestiga- tionsmadeavailableonceall theworkisdone. Asforthetheory,it seemsthatanypossible Furthermore,thereclearlyis no needfor pub- progressheavilydependsonobservationstoset lishingthesameresultsseveraltimesin differ- limitsto thenearlyinfinitenumberof possible entjournalsor proceedings. combinationsof parametersthatgoverncompu- tationsof all aspectsof cataclysmicvariables. Ourhopeisthatthemonographathandwill Andhere,atthecurrent,stillverygeneral,stateof makeiteasyfortheoristsaswellforobserversto thetheory,it isnecessarytoresorttostatistically findtheirwayintoandthroughthejungleof re- significantsamplesof observations,in orderto suitsandpapers.Wetried to point out relations notmistakeidiosynchrasiesofindividualsystems and gaps of knowledge. We hope this will be- forgeneralpatternsoftheentireclass.Investiga- come a useful reference book.

756 SUBJECT INDEX

abundances (see chemical composition) - CNO 330-331,341,359-361,401-402, 407, 424, absolute magnitude 5, 10-11 522, 564 - dwarf novae and nova-like 4-5, 27-28, 161 - 162, - 331,334, 361,404,562 176, 216 - O/H 360, 564 - novae 5,275-280, 294, 373 classical novae 11,261-369, 371-410, 413-510 - symbiotics 5 classification 1 accretion disk 3-4, 11, 35 I, 362 - dwarf novae and nova like 15-17, 19, 234-235 - column 188-192 - novae 262, 298-300 - funnel 188 - symbiotics 583,607-609 - hibernation model 294,408-410 chromosphere-transition region-corona (see spectrum) - inclination 294, 320. 373 CN Orionis stars 16 - instability 3, 171-179, 181-184, 449-450 colliding winds 656-661 - luminosity 294, 320-321,373, 41 I color 32-35, 38-39, 41-42, 44-45, 47-48, 52-53, 69, 96, - mass accretion rate 320-321,406-408, 411-412 98, 103-104, 112, 119, 121, 126-127, 134, 140, 147, - models 320-321,408 150, 160-161, 164, 171, 176, 178, 195, 199, 211-212 - pole 188-192 - of flickering 55 - power law 316-320 - of oscillations 58 - shock 188-192 - of superhump 52-53 - symbotic 654-657 - of X-rays (see hardness ratio) - ultraviolet spectrum 316-317, 319-324 combination spectrum 1 - X-rays 342-345 common envelope 217-219, 227-229 Algol variables 146, 149, 216 cool (late-type) component 592-596, 638-640 alpha disk 154, 169-178 corona 65, 153, 155, 194, 199-200, 213 Alfven radius 153, 188, 191, 195 cycle-amplitude relation 511-512 AM Canum Venaticorum stars 19-20, 95-96, 140-143, cyclotron radiation 139, 188-190 178, 222,229-230, 235 decline 346-355,587-588 AM Herculis stars 19-20, 95-96, 125-140, 153, 188-191 degenerate star 222, 226-227,229 anti-dwarf novae 19-20, 95-96, 102-112, 125, 172, 234 descendants 215,227-229 Bailey relation 27 disc instability (see instability) Balmer continuum 593,596 distance 7, 161-163,212,214, 351-352, 485-486, 518. Balmer decrement 73, 135 521 Balmer jump 73, 81, 84, 110 DQ Herculis stars 19-20, 95-96, 102, 112-124, 190, Balmer lines (see spectrum) 234-235,263,317,320, 485-487 banana diagram 56-57 D-type symbiotics 607-609 (see also classification) beat period 50, 52, 123, 192 D'-type symbiotics 607-609 (see also classification) beat phenomena 50 dust 300-303,307-313,403-404 Be stars 4, 599 - emission 404, 595 binarity 6, 21, 145, 148, 150-151,235 - formation 307, 408-404, 564 binary model 8-9, 589 - opacity 394 binary period 9 dwarf novae 11, 15-94 binary system 6, 8-9, 418, 435-437,442-443,448-453, eclipse 587, 589, 596, 601 485-487,490-497,523,546-547, 556-558, 642, 651 - eclipse mapping 163-164. 203-206, 209-212 661,727-728 eclipse, secondary 40-41, 44, 98, 102-103, 130-13 I bipolar flow 658 eclipsing binary 261,485-487 black body disc 192-194 Eddington limit 278-279, 329,341,371,378,385,397- black dwarf 221-222 398,400, 564-565 blue continuum 1,593,596 (see satellites) bolometric correction 164, 216 electron density 38 I, 554 boundary layer 46, 73. 152, 154-155,185-188, 194-195, electron temperature 381,554 199, 207,209, 212-214 energy BQ stars 4 - distribution 420-421,423,446, 452, 465,515,517, B stars 207 605 bremsstrahlung 71-72, 188-189 - radiative 5 bright spot (see hot spot) - mechanical 5 carbon stars 594 envelope 5,351-369, 372-373,405-406, 424,426, 428, Chandrasekhar limit 218,220. 228, 410-411 430, 440, 446-448, 467-470, 479-485. 501,521-523, chemical composition 170, 194, 207, 217-218,220, 13/18, 13/32-35,731 226-227,315,331-334, 336, 340, 341,359-362, 381, evolutionary state 194, 214-231 398, 400-403,424, 426, 479-485,510, 517,562, EXOSAT (see satellites) 729-730 expansion velocity 5,298-300, 414-415

757 extinction(reddening),interstellar289,320,397,623 maseremission618-621 extremeultraviolet(EUV)317,320,343 mass9 flare42-43,47,97,113,117,119,121,126,128-129, -accretion155,163,207,213-214,222,224,227, 138 406 flickering16,35,45,50,54-55,60,62,78,96-99,103, -determination156 106,113,116,118-119,121,126,128-131,140, -ejection376-378 18I,262,589 - loss(seealsoPCygniprofiles)94,206-209,218, free-free 220,223,415,517-518,555 -emission535,554 -ratio!/9,20-21,157-158,160-161,184,192,199, -opacity535,554 219-220,222,228 gammavelocity86-88,92-94,111,138,157,184 -stellar21,156-158,160,214,216-219,229 gap(seeperiodgap) -throughput163-164,172,193 globularcluster216 -transfer1,5,102,154,162,164,176,178-179, graingrowth(seealsodust)403-404 184,186,191,193,195,199-201,206,211,214, gravitationalradiation163,220,222-231 222-225,227-228,230 hardnessration7I, 106 Mira-typesymbioticstars693 hibernation230-231,371,408-410,565 Miravariables95,694 hotcomponent556-558,622 models hotspot38,46,55,64,149,152,155-157,164,166, -causesofoutbursts374-392 168,174,187,206 -centralstardominantmodel388-389 HubbleSpaceTelescope(seesatellites) -combinationofmodels390-391 hydrogenbound-free,free-freeemission554 -continuedejection342,384-388,564 imagery -hibernationmodel408,565 -optical424-425,440,475 -hiddenparameters374 -radio617 - instantaneousejection382-384 individuality3 -non-sphericalmodels408 instability -novaeclassical378-391 -disc171-179,181-184 -novaerecurrent391-392,410 -transfer172,178-181,183-184 -nuclearburning348 intermediatehump35-36,48 -snowplow351 intermediatepolar119-124(seealsoDQHerculisstars) -symbiotic647-661,672-674,679-683,723-725 interstellarabsorption65,161-162,216,537 novae interstellarextinction(orinterstellarreddening,seealso - activephase275-281,296-300,327-341 extinction)278,320,397,623 -catalogue262 ionizationmechanism596 -distance280,351-352,485 IRAS(seesatellites) -generalproperties261,561-565 IUE(seesatellites) -infrared299-314,403-404,419-423 jets693-694,700 -luminosity268-270,342 Keplerianvelocity(seevelocity) -masses449,502 Kukarkin-Parenagorelation26-27 -nebularphase342-369,466-467 K-velocity79,86-87,94,119,138,157-159,161 -photometry262-281 lategiant4 -quiescentphase262-275,281-296,316-327 late-typespectrum(seespectrumofthecoolstar) -remnant561 lightcurve371,414,427,443-446,460-461,469,491, -spacedensity409 500,513,542,544,547,585-592,664,666,674- -spectroscopy281-299 675,679,704-706,710 -ultraviolet315-342,418-419 limbdarkening194-195,209 -variabilityonshorttime-scale262-263,266-268,418 limit-cycleinstability(seeinstabilitydisc) -X-ray261,342,345-346,405-406,452 lineprofiles417,524-525,598-599 novaenvelope(orshell)351-356,405-406,424,426, -broadcomponent,wings593-594 446-448,467-470,501 -linewidth599 - chemicalcomposition359-362,381,400-403,424, - multiplestructure600 426,479-485 - PCygni2,81-82,84,88,90-91,93,99,110-124, -coolants361,489 155,206-209,213,598-600 -deceleration351-355,373, -WRfeatures596-599 -dustenvelope300-303 LTE,non-LTE195,198,206 -expansionvelocity351-355 luminosity -fillingfactor402 -bolometric329-330,341-342,346-347,371,384, -masses354-355,403,405,410 397,422,475-476,540,547,561,564 -models388,405-406,501 -kinetic398 -morphology356-359,361,372-373,425,428,430, magneticaccretion124,140,188-192 440,475 magneticbraking163,219-220,222-225,228-229,231 - radioemission366-369,405,423-424 magneticfields96,124,139,153-155,165,169,184- -recurrentnovaecircumstellarenvelope391-392 186,188-192,199,225,230,263,593-595 -snowplowmodel353,390

758 -spectra360-363,365-366,439-440,442 -orbital9,20-21,193,221-222,226-227,229 -structure361,364-365,405-406 -outburst17,27-28 -temperature361,480 -jitter140 nova-likestars1,3,95-144,292 photosphere(seequasi-photosphereorpseudo- novaoutburst photosphere) -absolutemagnitude275-281,438,459,488,491 planet227,229 -causes(classicalnovae)406-410,564-565 planetarynebulae228,583-584,596 -causes(recurrentnovae)410-412,565 polars(seeAMHerculisstars) -classification275-279,373 polarsintermediate(seeDQherculisstars) -energy(mechanical)5,297,398 polarization64-65,96,98,106,113,125,127,129-135, -energy(radiative)5,297,408 140,188,190,263,461,603-605 -expansionvelocity298-300,351-354,374,564 population216 -extremeultravioletradiation(EUV)343 pre-cataclysmicvariables(seealsoprogenitors)228 -infraredemission300-314,467-470,561 pre-nova262,281 -lightcurves275-280,414,427,450,491 primarystar(seealsowhitedwarf) -lightcurvesbolometric330,332,371,468-469 progenitors215-216,219,227-229 -lightcurvesinfrared302-304,306,308,371,468- pulsation(seealsooscillations)591 469 -whitedwarf(seecoherentoscillations) -lightcurvesoptical275-280,348,371,383-384, -X-ray60-61,106-107,124,130,213,269 460-462,471,475,500 quasi-periodicoscillation(seeoscillations) -lightcurvesradio371,473-474 quasi-photosphere(orpseudo-photosphere)384,394 -lightcurvesultraviolet330,371,474 -density373 -lightoscillations393,444 -mass-loss394,397 -speedclass(seeclassification) -pressure373 -ultravioletradiation327-342.474 -radii385-386.388 -X-rayradiation346-347 -temperature386,394 novaoutburstspectra296-300,372-373,417,426-428, quiescentnova(seealsonova)262-275,281-296, 438-439,460-461,465-466,470-475,490-492, 373-374 5OO-510 -absolutemagnitude280,374 -continuum394-400 -discinstabilities41I,457 -coronallines298,391-392,510 -lightcurve264-265,269-275,435-437,444 -diffuse-enhanced297,373,380,391,393 -orbitalinclination293-294,501 -nebular297,373,380,391,393,400 -orbitalperiods442,448,486,502,562 -orion297,373,380-381,391,393.509-510 -oscillations458-459,476 -post-nova298 -photometricproperties262-275,435-438,449-453, -pre-maximum297,373,380,390,393,401,502 457-459,486.487-488,490-491,562 -principal297,373,380,391-394,401 - radialvelocity435,438,442,449 -recurrentnova337,339-341 -spectra281-296,373,427,441.452,455,489-492 -stratification393 -ultravioletlinespectrum324-325,432-433,490, -ultraviolet327-342,464 497 OAO-2(seesatellites) -ultravioletradiation315-320,430-433 OH/IRsource(seemaser) -ultravioletspectralvariability324-326,430-433, olddiskpopulation600-601 495,498 oldnova(seequiescentnova) -X-rays342-346,442,457,484,498,563 orbit(seebinaryperiod) quiescentrecurrentnova513-514,519,525-531,542, oscillations(seealsopulsation)35,54,56-64,113,185, 546-548 192,213,234-235,372 -photometry547 - coherent54,56-64,98,114,116-119,122,124, -radialvelocity546 130,185,234,262-263 -spectrum326-327,515-516,519-521,538-540,546 -colors58-59 -UVspectrum519,521,539-540,549-554 -dwarfnova(seecoherentoscillations) -X-ray556-557 - quasi-periodic54,58-64,106-107,119,126,129- radialvelocity35,76,78-81,86-88,92-93,102,109- 130,140-141,185-186,265-268 110,117,119,123-125,134,138-139,142,148, - whitedwarf(seecoherentoscillations) 150,156-161,187,200,202,209,212,218,600-601 Ostars207 radioemission4,65,73,119,134 outburst5,21,586-588,600 - fromnovae366-369,405,473-474 -anomalous22-23,34 - novae:emissionmechanism366-369 -duration5 - novae:extendedradiosources366,368 -intervalbetween5 - novae:radioenvelopemasses366 period263,589-592 -novae:radiolightcurve367,405 -binary9 -novae:radiospectra369 -changes45-46,97-98,111,115,119,121,124, radius,disc160 186-188,264 radiusprimarystar153-160,187,199,214,216 -gap9,20,24,126,163,219,222-226,235 -secondarystar156,160-161,219

759 RCoronaeBorealisstars147 recurrentnova1-3,511 - emission line I/3,297-298,441,514-516, 519-520, -amplitude-cyclelength511-512 523- 531,545, 550-552, 596-600, 627-630, 729 -chemicalcomposition517 - high ionization 298, 466, 467, 490, 53 I, 66 I -companion518,556-558 infrared 3,299-314,419-422, 531,536, 605-613, -definition512-513 719-720, 722 -distance518,521 - molecular 588, 592-595 infraredspectrum533-535 - optical 665 -mass-loss517-518,555 - quiescent 4 - radio 3. 366-369, 424, 473-474, 519, 538, 613-618, -model547-548 719-720, 722-723 -nebularspectrum521-522 -orbitalradialvelocity546 - ultraviolet 2-3, 315-320, 418-419, 430-432,466, -outburstcause410 495-499, 516-517, 539-540. 550, 622-634, 664- 671,675,678-683. 697-699, 703-704, 716-721 -radioemission369,405,537-538 - variations 2, 415-417, 454-457,470-475,630-631, -ultravioletspectra327-329,334,337,339-340, 710-711,730-731 539-540,550-552 - X-ray 2-3, 342-346, 498-499, 541-542,634-636, -X-ray392-393,541-542 681-682, 723 revurrentnovaenvelope392,521-523 SS Cygni stars 16 recurrentnovaoutburst523-532 standstill 16-18, 29, 31,48, 56, 66. 84, 88, 234-235 -coronallines519,524,531,544 Stark effect 147, 202, 206 -expansionvelocity513,543 statistics 7, 19 -lightcurve513,519,542 supercycle 29-3 I -outburstmechanism543 superhump 16-18, 42, 49-54, 84, 86, 88. 113, 185 - spectrum514-515,519-531,523,536,539-540, - color (see color) 541,543-544,546 - late 48-50, 184 - UVspectrum516-517,539-540 - period 49-50 reddening(seeextinction) supermaximum (see superoutburst) rejuvinatedwhitedwarf652-654 15, 216-218, 229 reversedX-raymode129,131,140 superoutburst 16-19, 21.24, 29-32, 34, 48-54, 56, 60, Rochemodel3,7,38,67,92,95,145,148,150-153, 62, 68, 71-72, 94-88, 90-91,235 156,163,192-193,209,233,235,376 - models 184-185 rotationalbraking(seemagneticbraking) rotationaldisturbance102,l 11,117,122,296 - period 30-32 RWAurigaestars95 surface brightness 162, 166 satellites surface density 166, 168, 170-172 - ANS=astronomicalNetherlandssatellite329 SU Ursae Majoris stars 16-21, 24, 29-32, 42, 48-54, -ArielV344 157, 184, 234-235 S-wave 80, 86, 91,94, 102, 110, 122, 124, 137, 142 -COPERNICUS=OAO-3329 symbiotic novae 4. 296, 587-589, 596, 635-636, 683- -EINSTEIN=HEAO-2261,342-345,484,563 693 -EXOSAT261,269,340,344-346,532.563 symbotic stars 1, 3,583-585, 728-729 -HEAO-I=HighEnergyAstronomicalObserva- - accretion processes (winds) 654-661 tory-I 342 - activity phase 586, 665-667, 678-679 - HST=HubbleSpacetelescope299 - binary models 651-654 IRAS=InfraredAstronomicalSatellite310-314, - blue spectrum 596-598, 675, 706 532-534,558 - emission lines 596-600, 627-630, 697-699 - IUE=InternationalUltravioletExplorer261,263, - infrared observations 605-613, 719-720, 722 316-317,323-331,333-336,341,361,521,532, - light curves 585-592,664, 666, 674-677, 704-706, 539-540,561-562 710 -OAO-2=OrbitingAstronomicalObservatory-2 - line profiles 598-600, 627-630, 697-699 263,329 - magnetic fields 601-603 -Voyager320 - maser emission 618-62 I S-curve170,172-178 - models 647-651,672-674, 679-683,723-725 shell(seeenvelope) - polarization 603-605 singlestarmodel649-65! - quiescent phase 586, 664-665 Sobolevapproximation202-203,209 - radial velocity 600-601 solardynamo(seemagneticbraking) - radio imagery 616-618 solar-typechromosphere598 - radio observations 613-616 spacedensity7,215,228 - red component 592-596, 638-640 spacedistribution7-8,214-216 UV spectrum 622-634, 664-671,675,697-699, spectrum281-298,416-417,427-428,461-466,467, 703-704, 716-721 490,503-510,540 - X-ray 634-636, 681-682, 723 -absorption1/3,73,75,78-83,86-87,117,119,134, symbiotic phenomenon 583-584, 638 139,151.157,161,546,550-552,583,592-595, temperature 381,392-393, 402 600 - color 381

760 -excitation38I -radiuscorrelation398-400,428,429 -ionization381 -stratification393-394 -Zanstra381,394,555-556 viscosity154-155,165-166,168-169,170-183,200- thermonuclearrunaway(TNR)3,341,371-373,376, 210,214 406-412,564 VYSculptorisstar(seeanti-dwarfnova) tidalforces168,184,225,228 whitedwarf4,374,556-558,601 transferinstability(seeinstability) -CO341,562 two-colordiagram32-34 -luminosity406 typeCV's(seeclassification) -mass374,409 UGeminorumstars16-21,23(I,234,235 -O-Ne-Mg341,404,408,562 ultravioletdelay67-68 -pulsation(seecoherentoscillations) ultravioletflare(seeflare) wind385-387,393,397 UVPerstars16 Wolf-Rayetstars207,373 UXUrsaeMajorisstars19-21,95-102,172,230,234 WSagittaestars17 variabilityI-2,6,261-262,415-417,583,586 WUrsaeMajorisstars42,44-45,149,216.228 - irregular586-587,589 XLeonisstars16 - longterm458-459,586-589 X-raysource268,317,336,340-341,342-346.391, -periodic415-417,589-592 634-638 -quasi-periodic591 - novaemission342-346 - secular592 - recurrentnovainoutburst346.392-393 - shortterm262-263,266-268,415-417,457,587. - symbiotics634-636.13/I9-20 589,598 yellowsymbioticstars608 veiling593,596 Zanslratemperature(seetemperature) velocity Zcamelopardalisstar16-18,20,29,48,56,146,148. -expansion298-299 172,230,234-235 -gradient380 Zeemansplitting134,139-140,188,190 -Keplerian152,154-155,165,169 Z-wave79 -radial(seeradialvelocity)

761

STAR INDEX

YZ Cnc 55, 91,274-275 Z And 4. 290. 584-587. 589. 592-593. 596-598.607. 612-613. 621. 623. 626. 628. 631-632.641-644. 648- HL CMa 77, 85-86 649. 651. 660. 663-674. 729

BG CMi 112 RX And 55.64-65.68.82. 89.285.287

AM CVn 140-143,227 AR And 222.287

TX CVn 601-602, 641,643 EG And 592. 598-599. 601-602.608.624. 641. 643- 644.651 OY Car 36-37, 39-43, 47-48, 50-53, 58, 77, 80, 82, 84, 90-91 R Aqr 598. 604. 606. 608-612.616. 619. 632.634. 635. 637-638.641. 644. 648.650-651. 656.693-703 HT Cas 32, 37, 55, 58, 70, 76, 79-80, 185,274, 288

AE Aqr 64.80. 95. 112-113.117-119. 145. 148. 200. BV Cen 36, 55.78 234. 274-275.290

V 436 Cen 34, 36, 49-50, 73, 77,274 FO Aqr 112. 123-124. 191

V 442 Cen 274 UU Aql 287 V 803 Cen 140-141 V 603 Aql = N Aql 1918 112. 193. 200. 263.266. 277. 281-282. 284. 286. 299. 312. 316. 318-320. 322. V 834 Cen 130, 138-139 324-325.327.343-344. 348. 351-354. 356. 372.381- 382. 389. 393-398.413.426-438 W Cep 602 V 605 Aql = N Aql 1919 310. 312-313 IV Cep = N Cep 1971 366 V 1301 Aql = N Aql 1975 304,404 VV Cep 648 V 1370 Aql = N Aql 1982 304, 308-311,314, 332-335, WW Cet 26, 34, 76, 288 342, 354, 366, 394, 562

Omicron Ceti 619, 634, 638 TT Ari 35 103-108, 110, 112, 124, 187,234, 274-275, 290 Z Cha 37-38, 46-48, 50-52, 67, 76, 80, 84-87, 93, 185, 203,203-206, 211-212,274 FS Aur 287

SY Cnc 288 KR Aur 108, 110, 290

YZ Cnc 288 SS Aur 26, 28, 287

TV Col 112-113. 124 T Aur = N Aur 1891 277,285-286, 314, 316, 320, 352, 357,359,375, 414, 487-490 GP Corn 140-142. 227 UV Aur 593-595,641,643-644 T Cor 314 V 363 Aur 98, 100 V 394 CrA = N CrA 1949, 1987 328. 752

UZ Boo 73 V 693 CrA = N CrA 1981 328. 332. 334. 562

AF Cam 222 T CrB = N CrB 1886, 1946 4. 230. 277. 285.287.296. 298-299.311-313.316. 319. 322. 326. 343.348. Z Cam 18, 25-26, 28, 31, 34, 39, 64, 66, 90-91,274, 366. 368-369. 375. 391-392.411.511. 543-558.561. 287 563.565.589. 598. 601-602. 609. 641. 644.648. HT Cas 288 65 I. 752 AC Cnc 37 BI Cru 598-599. 610-611. 640. 641. 644 SY Cnc 57-58, 66, 274 BF Cyg 583. 587-588.596. 599. 639.641. 644. 648

763 PRECEDING PAGE BLANK NOT [:lIMED CHCyg587,589,593,598,601-604,617-618,620-DQHer = N Her 1934 112-117, 145, 148,187, 191, 621,623,625,627,631-632,634,637,641,643- 205,234,262,267-268,274-275,277,286,295, 644,650-651,663,703-725,728,730 299,311-312,314,316,320,323,345,352,354- 355,359,361,366,373,376,378-379,381,393, CI Cyg64,583,587,589-590,592-593,596,598-599, 400-403,413,475-487 600-605, 607,619,626,632,641,643-644,648, 660,727 V 433Her = N Her 1960696, 641

EMCyg24-25,40,67,69,78,80,161,217,219,274,288 V446 Her = 262,352

EYCyg 8, 150 V 533 Her =Her 1963 112-113,117, 262, 274-275, 277,281-282,286,298,316,352,354,366 SS Cyg 15,18-19,22-26,30,32-34,40,42-43,59,60- 61,64,66,68-72,74,76,79-83,85-87,128,145- YY Her 641,644 148,150,157,159,161,176,193,213,274-275, 288,343 EXHya 20,112-113,119-122,124,134,289

V 476 Cyg = N Cyg 1920 352,354 R Hya 619

V 503 Cyg 74 RW Hya 583.592, 601-602, 641-644

V 1016 Cyg =MH328-116 585,588,595-596,599, BL Hyi 126, 131,138 601,604,606-607,610,612-617,623,628,634-635, 641,644,651,654,683-684, 688-691,693,727,752 VW Hyi 16-18, 24-26, 29-31, 33-36, 40-41,47-53, 60- 64,66-68,71-72,74-75,84-85,88-89,176,181-182. V 1329 Cyg =HBV 475 584, 588,593,596-597,599- 213,274-275,282 600,607,613-614,632,641,644,654,658,683- 684,688-691,693,752 WX Hyi50,68-69,74-75,88

V 1500Cyg = N Cyg 1975 263-264,266,277,279, CP Lac = N Lac 1936 348, 351-352, 355, 381-382 297-300,304-307,347,349-350,352,354,366-368, 372.392,394,401,403,405,413-426,753 DI Lac = N Lac 1910 285-286,348,353

V 1668 Cyg = N Cyg 1978 265-267,294,300,303-305, DKLac = N Lac 1950 353,381,391 330-331,334-335,349,354,372,403-400,413,460- 470 T Leo 80,157-158,289

CM Del 74 X Leo25-26,274,289

HR Del = N De11967 112,261,266,269,277-278, DP Leo 125-126,128,130 280-281,298,300,312-314,316-320,322-323,325, 343,350.352.354,360,366-368,389.401,404- ST LMi130,132-135,139 405,414,500-510,561,563 AY Lyr 26 AB Dra 70,288 CY Lyr 82 AGDra 584,587,591-594,596,598-602,612,622- 628,630,632,634,636-637,641,643-645,651, HR Lyr 353 655,663,674-683,730-731 MV Lyr 95,103-104,106-108 YY Dra 290 EF Eri 126-128, 130, 133-135, 138 TUMen 20,24,26,32,52,85-87,93,222,235

DN Gem = N Gem 1912 277, 285-286, 299, 354, 389 BTMon =NMon 1939 202, 230, 262, 285-286, 294, 316,320-321,352,409 IR Gem 288 BXMon 591-593,599,602, 641,644 U Gem15,19,24-26,30,37,40-42,59-61,64,66,68, 70-72,74,76,78,80,145-147,161,187,205,213, GQMus = N Mus 1983 311-312, 314, 328, 337-338, 217,219,274-275,288 341-342,345-346

AH Her 17,34,57-58,64,88-89,176,274,289 SYMus 591-592,596,641,644

AM Her95,125-126,128-135,137-140,290 ILNor =NNor 1893 348, 375

764 RS Oph = N Oph 1898, 1933, 1958, 1967, 1985 4, 16, CP Pup = N Pup 1942 9, 264-265,277, 282, 284, 316, 277,296, 285,287,300, 312-313,316, 322, 327- 343,353,359, 361,392,413-414,438-443,562 328,337,339-341,346. 355,368-369, 391-392, 411- 412.511-512. 523-542,563,565,589,598, 602, RX Pup 598-599, 607,609-612,615,632-633,641, 609, 641,644, 752 644

V 426 Oph 64, 290 UY Pup 26

V 442 Oph 290 VV Pup 125-127, 129-132, 135-136, 138, 189, 290

V 841 Oph = N Oph 1848 269, 284, 286, 316, 323, T Pyx = N Pyx 1890, 1902, 1920, 1944, 1966 277, 284- 343,353 285,287,298-300, 316, 326-327,355,359, 362, 412, 511,519-523,563,565,752 V 2051 Oph 37, 39, 55,274, 29(I CL Sco 291 V 2110 Oph = AS 239 290, 588,684, 688 FQ Sco 26 BI Ori 26 HK Sco 291 CN Ori 26, 35, 42-43, 46-48, 58, 68, 79, 82, 111, 124, 181-182,235,274, 289 KP Sco 312-313

CZ Ori 26, 289 T Sco 216, 352

AR Pay 587, 592, 601,609, 641,644, 727 U Sco = N Sco 1863, 1906, 1936, 1979, 1987 4, 300, 328,334-335,412.511,513-519, 563,752 BD Pav 40, 42, 44-45 V 711 Sco 277 RU Peg 26, 54, 57, 59, 78, 161,274, 289 VZ Scl 103, 107 AG Peg 588-589, 591-592, 596, 598-599, 601,603, 605,613,617,619-620,623,628,630, 641,643- EU Sct 312 644. 649, 651,654,683-684, 688-690, 693,729-730 N Sct 1991 368 IP Peg 37 V 368 Sct = N Sct 1970 353,366 TZ Per 73 V 373 Scl = N Scl 1975 298, 353 UV Per 26, 74 CT Ser 284 AX Per 583-584, 591,599, 601,609, 619, 631,641, FH Ser 277,279, 285, 287, 300, 302,304-308.310- 643-644, 648-649, 663,730 314, 328,348,352-353.355,366-368,384, 389, 397-399, 403-405,413,470-475,561,564 GH Per 296, 343

GK Per = N Per 1901 28, 176, 222, 230, 261,268-269, LW Set = N Ser 1978 302, 304, 404 274-275,277,285-286, 294,316-317,320, 344, 348, 352-353,355,359, 361,368-369, 312,375,382, LX Ser 98, 103-104. 107, 110, 201 389, 394-395,397-399, 413,443-459, 562-563. 753 MR Ser 126, 135 KT Per 274, 289 MU Ser = N Set 1983 310 V 471 Per 608 RT Ser = N Ser 1909 296, 588, 601,641,644,654, RR Pic = N Pic 1925 193,231,262,269-275,277-278, 684, 688-690, 693,752 284, 299,314,322,325, 343-344, 348, 350, 352- 353,355,357-363,375,382,396,398, 400, 414. UX Ser 26, 69 490-499, 561 UZ Ser 282 AO Psc 113, 122-123,191 X Ser 287 TY Psc 50 RW Sex 274

765 SWSex97-98,102,201 CK Vul = N Vul 1670 231,269, 294, 310, 314, 362, 364-366, 374, 414, 561, 751 HM Sge 585, 588, 596, 599, 604, 606-607, 609, 612- 614, 616-617, 628, 634-636, 641,644, 654, 683-684, LV Vul = N Vul 1968/1 277,279, 352, 355 688, 690-691,693,730 NQ Vul = N Vui 1976 277,279, 303-305,312-313, RZ Sge 76-77 355, 366, 368, 372,403

V Sge 16, 95, 291 PU Vul 588, 605, 612, 630, 632, 641,644, 663,684, 688-693, 728, 730 WY Sge = N Sge 1783 23 l, 269, 274, 294-296, 353, 374, 414, 561,751 PW Vul = N Vul 1984/1 328, 336, 345-346

WZ Sge = N Sge 1913 1, 17, 24, 29, 40, 42, 44-45, 50, QQ Vul 126-127, 131-132, 136-137, 140 52-53, 63, 67, 91-94, 112-113, 122, 134, 268, 274, 289, 512 QU Vui = N Vul 1984/2 312, 328, 339, 345-346, 372, 404, 562 LG Sgr = N Sgr 1897 312-313 2A 0311-227 290 V 949 Sgr = N Sgr 1914 277, 312 AS 201 608, 614 V 1016 Sgr 312, 313 AS 239 see V 2110 Oph V 1017 Sgr = N Sgr 1919 (1901, 1973) 4, 8, 150, 285, 343,369, 51 i, 542-543, 65 l, 752 AS 289 584

V 1059 Sgr 343 AS 296 601

V 1148 Sgr = N Sgr 1943 297 AS 338 604,606

V 1223 Sgr 112-113, 123-124, 191 BD -7 3007 = RW Sex

V 3885 Sgr 98, 100, 274, 282 CD -42 14462 = X 3885 Sgr

V 3890 Sgr 291 CDP -48 1577 = IX Vel

V 4077 Sgr = N Sgr 1982 311-312, 314, 328, 335-336, E 2000 +223 314, 351 366, 368 IE 0643.0 288 RR Tel 4, 296, 312-313, 588, 591-592, 596-599, 601, 607, 608-611,613-614, 621,624, 627-629, 632, 634- G 61-29 = GP Corn 290 635, 639, 641-642, 644, 654, 683-691,693, 727, 729, 752 GD 1662 = VY Scl

EK TrA 67,86 GX 1+4 (4U1728-24) 634

RW Tri 55, 64, 80, 97-99, 101-102, 128,207, 212 2H 2215-086 291 291

AN UMa 125, 130, 291 HI-36 613-615,617, 727

SU UMa 17, 26, 32, 66-68, 73 H2-38 596, 610

SW UMa 148, 289 He 2-17 614

UX UMa 97-98, 100, 114-115, 149-150, 187, 192, 207, He 2-34 610, 612 274, 291 He 2-38 596, 611,614, 641 CQ Vel 313 He 2-106 610-612 IX Ve| 97, 99-100, 102 He 2-390 614 KM Vel 611 HD 149427 594, 608,641 TW Vir 40, 74, 289

766 X0643-1648 (IE) = HL CMa HD 330036 594, 641

X1729+103 (3A) = BG CMi HZ 29 = AM CVn

X1012-029 (PG} = SW Sex Lanning 10 = V 363 Aur X1013-477 (Hi = KO Vel LMC $63 595-596, 641,644

X1103-254 (CW) = ST LMi LMC SN 1987a 384

X1114+!82 (IE) = DP Leo MI-2 594,608,641

X1140+719 (PG) = DO Dra PS 74 274

X1346+082 (PG) 140-142 PS 141 = VY Sci

X1405-451 (H) = V 834 Cen UKS Ce-I 595

X1550+191 (PG) = MR Ser Weaver star 595

X2003+225 (HI = QQ Vul X0139-608 = BL Hyi

X2215-086 IH) = FO Aqr X0311-277 (2A, 3A) = EF Eri

X2252-035 (H) = AO Psc X0527-328 (2A, 3A) = TV Col

0623 +71 291,351 X0538+608 (H) 126, 135

767

LIST OF CONTRIBUTING AUTHORS

Antonio Bianchini Osservatorio Astrofisico Via Dell Osservatorio 8 136012 Asiago, Italy

Astronomishes lnstitut HilmarW. Duerbeck Universitat Muenster Domagkstrasse 7_5 D 4400 Muenster, Germany

lnstitut d'Astrophysique Michael Friedjung 98 his Boulevard Arago F 75014 Paris, France

Osservatorio Astronomico Margherita Hack Via G. B. Tiepolo 11 ! 34131 Trieste, Italy

Constanze la Dous (preparation of Part !) Osservatorio Astronomico Trieste, and ... Institute of Astronomy University of Cambridge (current) EAS/IUE Observatory Villafranca del Castillo Apartado 5072 7 E 28080 Madrid, Spain

OAT Pierluigi Selvelli Box Succ Trieste 5 Via G. B. Tiepolo 11 ! 34131 Trieste, Italy

IAS Roberto Viotti CNR CP 67 1 00044 Frascati, Italy

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