Publications of the Astronomical Society of the Pacific 97:138-150, February 1985

IUE AND OPTICAL SPECTRAL SCANS OF U SAGITTAL: AN ANALYSIS AND COMPARISON WITH U CEPHEI

JAN J. DOBIAS AND MIREK J. PLAVEC0 Department of Astronomy, University of California, Los Angeles, California 90024 Received 1984 October 2

We have determined the effective temperatures and surface gravities of the components of U Sge, using a combination of optical spectral scans and IUE spectra. The primary component is found to be a B7.5 V star with Te{{ = 12,250 (± 250) Κ and logg = 3.9 ± 0.1. The secondary component is of spectral type G4 III-IV. The best value for distance is 295 ± 20 pc, and the color excess is E{B— V) = 0^06. Our determination of the slope of the radial curve of the secondary star at primary eclipse confirms the previous results (mainly by Tomkin) that the mass ratio is very close to 3. Com- bination with photometric solutions enables us to derive reliable system parameters, given in Table VII. One high-dispersion IUE spectrum has been studied in greater detail. We have been able to fit selected lines of Fell, Sill, and Sim assuming solar abundances. The resonance doublet of C π at 1336 Á could not be fitted and we are in- vestigating several possible explanations. The photospheric lines of the hotter component are broadened by rotation; we -1 find ürot 100 km s . We discovered emission lines of the W Serpentis type in an IUE spectrum taken in the total primary eclipse. Although the observed fluxes of these lines are very low, the powers emitted are half as strong as those seen in U Cep. This is rather surprising, since we did not find any other evidence of circumstellar matter or mass transfer activity. Key words: stars: abundance—stars: circumstellar shells—stars: eclipsing binaries—stars: emission-line—stars: individual (U Sge)—ultraviolet: spectra

I. Importance of the System the Varo image tube attached to the coudé spectrograph The eclipsing binary U Sagittae (HD 181182 = of Lick Observatory. We will attempt to present a com- HR 7326 = BD+1903975) is the brightest semi- prehensive picture of the system and then discuss its re- detached eclipsing binary of the Algol type that displays lation to more active interacting binaries. total primary eclipses. It is also one of the simplest sys- II. A Review of Previous Work tems of this type, since the photometric and spectroscopic perturbations are smaller than in most Algol vari- Although U Sge is rather free of circumstellar mate- ables. There appears to be so little circumstellar material rial, there does not exist yet a clearly defined picture of present in the system that it may be considered as nearly the system, and individual solutions differ. The first pho- dormant, i.e., probably near the end of the mass-transfer toelectric light curves were obtained in the blue and red phase. This gives an excellent opportunity to derive the colors by Irwin, but have never been published. They properties of the component stars with greater reliability were used by Kopal and Shapley (1956) to obtain the than in other Algol variables. On the other hand, an in- first good set of photometric elements. Later, the star teresting problem emerges when we compare U Sge was observed by Cester and Pucillo (1972) on the UBV with U Cephei. While the two systems are rather similar system, and by McNamara and Feltz (1976) on the in a number of characteristics, U Cep is much more ac- Strömgren four-color system. Both groups of observers tive than U Sge. Why? What parameter, or set of param- published solutions based on their observations. More- eters, causes this dichotomy? over, the same observations were subsequently redis- But perhaps the most important aspect of U Sge is cussed by Cester et al. (1977), Tsouroplis (1977), and Al- that it offers the opportunity to study the chemical com- Naimiy (1978). position of its atmosphere and to see if we can find evi- There exists a general consensus that the hotter star, of dence for abundance anomalies related to the past evolu- about the spectral type B7 V, has a fractional radius be- tion of the system, during which a large-scale mass tween or near the values 0.21-0.23, and that the frac- transfer should have occurred, if our current theories of tional radius of the cooler, G-type subgiant, lies between close-binary-star evolution are correct. or near the values 0.28-0.30. The system is totally eclips- For all these reasons, we have observed U Sge with ing and the inclination must be very near 90°; actually, the International Ultraviolet Explorer (IUE) satellite, some of the above solutions simply accept that value. with the ITS scanner of the Lick Observatory, and with The value most deviating from 90° is found in the Cester and Pucillo solution, and is 89? 3. Since the elements obtained by Cester et al. and by McNamara and Feltz "IUE Guest Investigator. agree very closely, we will adopt r(B) = 0.225 for the 138

very nearly spherical primary component, and r(G) = metric investigations of the system by McNamara and 0.295 for the photometric radius of the secondary star. Feltz (1976) seemed to support Naftilan's result. Their The elements in both solutions agree closely except for value of the m1 index, if interpreted in the usual way, the U and u light curves. These near-ultraviolet light suggested that the secondary star in U Sge must be met- curves require systematically larger values of both frac- al-poor. Quantitatively, their result, [Fe/H] = —0.63, tional radii, and most likely indicate the presence of ad- agrees very well with Naftilan's. However, the much ditional weak radiation at the Balmer limit. The small more detailed study by Parthasarathy, Lambert, and uncertainty in inclination is not important for our work. Tomkin (1983) does not corroborate these results. They In spite of total eclipses and relatively high apparent applied spectrum synthesis to Reticon spectra obtained brightness of U Sge, the absolute parameters of the sys- near 630 nm with a much higher dispersion, 0.75 Â tem were poorly known until very recently. The primary mm-1, and found the metal abundances normal within component of U Sge is so much brighter than the sec- the observational errors; [Fe/H] = 0.0 ± 0.3. ondary star that for many years the only information on The important result of the work of Parthasarathy, the radial-velocity amplitude of the secondary star was Lambert, and Tomkin is the detection of carbon defi- the determination of the slope of this radial-velocity ciency and nitrogen overabundance in the secondary star curve at the primary eclipse by Joy (1930). Joy actually in U Sge. They applied spectrum synthesis to their Digi- gave the ratio of the slopes, expressed in terms of the con spectra of the CH lines at 4300 Â and CN lines at mass ratio: ^(B)/sJl)?(G) = 3.3. From a complete orbital 4215 Â and 3888 Â. They found that the secondary star solution for the primary star, he obtained K(B) = 67.9 is deficient in carbon, [C/Fe] = —0.5, and over- km s_1. This latter value was subsequently only slightly abundant in nitrogen, [N/Fe] = +0.5. This result not revised by McNamara (1951a) to 69.7 km s-1 (± 0.7 km only does not contradict the theory of the evolution of s-1). Thus, the implied value of K(G) would be about the Algol variables by mass transfer (Plavec 1968; Ziol- 230 km s-1. Plavec (1967a) rediscussed Joy's slope deter- kowski 1969), but actually strongly supports it. mination and concluded that the mass ratio is closer to While there is every good reason to believe that the 3.0, thereby revising K(G) down to 211 km s_1. Joy's present system of U Sge is a product of mass transfer, the spectrograms were taken with a reciprocal dispersion of present level of such activity in it is very low. In spite of only 75 Â mm-1; nevertheless, Plavec's rediscussion many attempts, emission in the optical spectrum was closely agrees with our new observations presented in seen only once: McNamara (1951b) obtained two spec- section VI, as well as with the recent accurate work of trograms on 1950 August 26, which show violet-dis- Tomkin (1979). placed emission components in the Balmer lines, in Mg π In an important development, Tomkin succeeded in 4481 Â, and Ca π Κ 3934 Â. McNamara (1951α) also no- obtaining a complete radial-velocity curve of the sec- ticed the typical asymmetries of the hydrogen absorption ondary star, using a Reticon. Tomkin measured the Na D lines, which are, however, much weaker than in U Ceph- lines, the infrared Ca π triplet, and the M g ι line at ei. Students of the light curves of U Sagittae point out 8806 Â. He derived K(G) = 209 km s_1 which, when that the secondary eclipse is anomalously deep, probably combined with McNamara's value for X(B), gives a mass because circumstellar material surrounding the primary ratio of almost exactly 3.0. Since our new slope determi- star causes additional loss of light of the secondary (Ces- nation, reported in section VI, corroborates this ratio, we ter and Puccillo 1972). The depth of the primary eclipse are confident that at long last, the absolute parameters is also variable, and McNamara and Feltz (1976) explic- of U Sge are known with an accuracy that is as good as is itly suggest the presence of a weak Balmer emisison. Ol- the standard for a well-studied Algol variable. son (1982b) finds, from his five-color photometry, evi- The problem of the chemical composition of the Algol dence for small brightness changes in the secondary star variables naturally attracted attention to U Sge because and for transient low-level mass-transfer activity. As a of its brightness. Naftilan (1975) studied two high-dis- new discovery, Olson (1982a) reports the detection of persion spectra of the secondary star, one obtained by fairly large light-curve variations which are most simply M. J. Plavec and one by D. M. Popper, both during the explained as radius changes of about ±2% or ±3% in total eclipse and covering the spectral region 350-480 one of the stars. Olson prefers to attribute these varia- nm at an average reciprocal dispersion of about 16 Â tions to the secondary star, although he realizes that in a mm-1. From a crude curve-of-growth analysis and a subgiant filling its Roche lobe, they should trigger fluc- more sophisticated spectrum synthesis, Naftilan derived tuations in the mass-transfer rate much larger than those a mild metal deficiency, [Fe/H] = —0.7. This was puz- observed. He suggests that the subgiant may not nor- zling, since the theory of a large-scale mass-transfer in mally fill the Roche lobe exactly. But is it also quite pos- Algol variables can hardly satisfactorily account for sible that the variability of the radius occurs in the pri- anomalies in metal abundances. Nevertheless, the photo- mary component. Very recently, McCluskey and Kondo

(1984) reported seeing moderately strong C iv and Si iv tegration times were quite short, on the order of two to ultraviolet resonance lines. These lines cannot be formed four minutes for the "R" setting, and four to eight min- in the relatively cool photosphere of the primary com- utes for the "B" settings. Dates and phases of the ITS ponent and are probably formed somewhere above it, in scanner data are listed in Table I. a region producing the weak emission lines of these ions The observed flux distribution in the optical region, all discussed in section IX. the way to 680 nm, is entirely dominated by the hotter We have studied U Sge as part of our ongoing in- component, except at the time of the primary eclipse. vestigations of spectral flux distributions in Algol-type We found no definite evidence of a contribution from variables. The hotter primary components are studied by the secondary star. Also, we found no evidence for any combining low-dispersion IUE spectra with optical scans continuous absorption due to the circumstellar material, taken with the ITS scanners at Lick Observatory. These nor were any optical emission lines present. The ob- flux distributions are then matched by model atmo- served energy distribution is that of a normal late-type Β spheres calculated by Kurucz (1979). The cooler com- star, and it is easily matched by a normal Kurucz ponents have been optically scanned during the primary atmosphere. eclipses. As a rule, these subgiants have effective tem- Ultraviolet observations of U Sge, made by M. J. peratures below the lower limit of Kurucz's model grid. Plavec, are listed in Table II. The low-dispersion data, Their spectral types are then obtained by matching them obtained in both the short-wavelength (SWP) camera by standard stars. The entire procedure and the results and in the long-wavelength (LWR) camera, have approx- for U Sge are discussed in section III for the hotter com- imately the same spectral resolution as the Lick scans, on ponent, and in section V for the cooler component. Radi- the order of 7 Â to 9 Â. They can therefore be combined al velocities of the cooler component, obtained with the with the optical data in order to obtain a complete flux Varo image tube of Lick Observatory, are discussed in distribution between 110 and 680 nm. There remains a section VI. small gap in the vicinity of 320 nm, where neither in- The important problem of chemical abundances in U strument provides a really realiable coverage. Sge was studied on high-dispersion IUE ultraviolet spec- Our program FITMEX matches the observed flux dis- tra of the hotter component. The procedure and results tribution to the Kurucz model-atmosphere fluxes, and are discussed in section VIII. Our search for emission uses also the color excess E{B — V) as an independent pa- lines due to the circumstellar material is described in rameter, determined by least-squares solution. The red- section IX, and section X summarizes and discusses the dening of U Sge must be rather small, since the IUE results of the entire work. spectra show no obvious 220 nm interstellar feature. This The orbital phases quoted in this paper are based on agrees with all previous determinations. Hall and van the ephemeris published by Olson (1982a), slightly modi- Landingham (1970) studied the open cluster Collinder fied again by Olson (private communication, 1984 No- 390 and concluded that U Sge is not its member, being vember 3): farther away with a color excess of 0^1. Olson (1975) adopted E{b — y) — 0^4, so that E(B — V) = 0^7. Kon- ( . min.) = HJD 2440774.4853 + 3.3805942 E. (1) Γ ρή do, McCluskey, and Wu (1981) determined E{B—V) — III. Spectral Type of the Primary Component TABLE I Optical scans of U Sge have been obtained with the Optical Scans of U Sagittae and Standard Stars ITS scanners of Lick Observatory, attached to the Casse- grain focus of three different telescopes: the 3-m, 1-m, DATE (UT) JD SPECTRAL EXP and 0.6-m reflectors. The ITS scanners at the 3-m Shane YEAR MON DAY 2,440,000+ PHASE ECLIPSE STAR TEL REGION (MIN) telescope was equipped with a green-sensitive image tube which gives good spectral coverage down to about 1979 JUL 31.267 4085.767 0.497 SEC,DESC U Sge .6m B, R 4, 2 325 nm. The two smaller telescopes were equipped with 1979 SEP 13.283 4129.783 0.519 SEC.ASC U Sge 3m B, R 4, 2 a red-sensitive tube which becomes insensitive shortward 1980 JUN 11.353 4401.8528 0.996 PRI,TOTAL U Sge Im B, R 8, 4 of 360 nm. Both scanners used similar gratings, which 1980 JUN 11.381 4401.8806 0.005 PRI,TOTAL U Sge Im B, R 8, 4 yield about the same spectral resolution, on the order of 1980 JUL 8.421 4428.9208 0.004 PRI,TOTAL U Sge Im B, R 8, 4 7 Â to 9 Â. In order to cover the entire optical spectrum, 1980 JUL 8.445 4428.9451 0.011 PRI,TOTAL U Sge Im B, R 8, 4 two different settings were used for the gratings. The blue (B) setting on the 3-m telescope covered about 200 SPECTRAL TYPE nm and was centered at 420 nm. The red (R) setting was 1979 SEP 23.456 4139.956 G4III HR1327 .6m B, R 2 centered at 580 nm, and the coverage extended just 1982 AUG 6.486 5187.9860 G8III o Psc 3m B, R 2, 2 beyond Ha. For the two smaller telescopes, the "B" set- 1982 OCT 2.450 5243.9521 G4III HR1327 3m B, R 2, 2 ting was shifted longward for about 40 nm. The in-

Owe. Our FITMEX solutions converged to the same val- ure 1 compares the observed flux with the selected Ku- ue; thus, E{B — V) = 0^6 was adopted for the color ex- rucz model, represented by open circles. The magnitude cess of the system. scale used here, and in the rest of the article, is based on The optical observations were made at nearly the the calibration by Hayes, Latham, and Hayes (1975) in same orbital phase as the ultraviolet observations, near such a way that we convert monochromatic fluxes Fx to the secondary eclipse, but they differ in dates by two monochromatic magnitudes mx by means of the formula years. Although secular variation is not likely in U Sge, = -2.5 logF - 21^17 we took the precaution of treating the two sets of data x (2) separately at first. The IUE data are safer from potential Figure 1 shows that the overall fit between the Ku- contamination by the G-type component. They were rucz model and the observed flux distribution in U Sge is matched to the Kurucz model fluxes at 82 points be- excellent. In fact, the overall fit is better than for any tween 110 and 320 nm. The best fit is obtained for a bilin- other Algol variable we have studied. It could be ob- early interpolated model atmosphere with Teff = jected that cases among Be stars are known when a 12,500 Κ and logg = 3.85 - 4.00, and E{B-V) = simple Kurucz model provides a good fit and yet there 0^6. The optical flux distribution (100 points between exists a continuum-contaminating circumstellar enve- 325 and 680 nm) is best matched by a similar model with lope. However, in our experience, such continuous con- Teff = 12,250 K, logg = 3.90, and E{B-V) = 0^05 to tamination does show at certain places, in particular at 0^8. These two results are so close that one fit to the the Balmer limit. Moreover, in U Sge even the strongest whole combined spectrum is justified. Then the best fit is shell lines observed in the ultraviolet are quite weak; Teff = 12,250 Κ and logg = 3.90. We estimate that the therefore, the continuous absorption must be negligible. uncertainty in the determination of the effective temper- An additional important argument is that our results ature is about ± 250 K, and about ± 0.1 in logg. Thus, agree perfectly with those obtained by Olson. For the ef- the spectral type of the hotter component is B7.5 V. Fig- fective temperature, Olson (1982a) derived 12,500 Κ from five-color intermediate-band photometry, and for TABLE II surface gravity, he found logg = 3.88 from hydrogen- IUE Observations of U Sagittae line profiles (Olson 1975). Date (UT) EXP IV. Comparison of U Sagittae and U Cephei (mid-exposure) 2,440,000+ PHASE ECLIPSE IMAGE DISP (SEC) There has not been much unanimity about the spectral 1981 Jul 30, 1812 4816.2583 0.580 SWP 14595 LG 15 types of the hotter components of U Sge and U Cep. 1981 Jul 30, 1815 4816.2604 0.575 LWR 11192 L0 20 They both fall into the range where few lines are avail- 1981 Aug 09, 1640 4826.1944 0.519 sec, asc SWP 14702 HI 1400 able for a good spectral classification in the optical re- 1983 Nov 07, 1020 0.002 pri, tot SWP 21468 LO 1500 gion, and those lines that are strong enough are strongly affected by rotational broadening, particularly in U Cep.

U SAGITTAE PRIMARY STAR

230 290 350 410 470 530 590 650 WAVELENGTH (nm)

Fig. 1—Flux distribution of the primary component of U Sge (continuous line), matched by an interpolated Kurucz model with Teff = 12,250 K, logg = 3.90 (open circles). The downward spike near 180 nm and the upward spike near 220 nm are spurious.

Photometric solutions are also not without difficulties, so V. Spectral Type of the Cooler Component that either way of determining the spectral class leads to Optical scans of the pure spectrum of the cooler com- uncertainties. The spectral types given in literature lie ponent were obtained during the total phase of the pri- between B6 and AO, and U Cep has typically been classi- mary eclipse on 1980 June 11 and July 8 (see Table I). fied as of a slightly earlier spectral type than U Sge. Our All these spectra show the typical energy distribution of spectral scans, obtained by a combination of IUE obser- a mid-G-type star. There are no emission lines visible, vations with Lick scanner data, provide a much broader nor is there any trace of an additional radiation source, basis for spectral classification. The data for U Cep have such as hydrogen bound-free radiation at the Balmer been published by Plavec (1983). Our result is that the limit, which we have found in many Algol variables. The primary star of U Sge is somewhat hotter than that of U observed spectral distribution of the cooler component Cep, contrary to what most previous determinations say. has been dereddened and then compared with scans of The spectral types are B7.5 V for U Sge and B8 V for U standard stars, since Kurucz models are not available for Cep. In terms of effective temperatures, U Sge appears stars of effective temperature lower than 5,500 K. to be hotter by about 1000 K. The overall distribution of the optical flux of the sec- This result is not an artifact of the application of the ondary star, between 360 nm and 680 nm, is very well correction for interstellar reddening. Figures 2 and 3 matched by the G4 III star HR 1327, if the magnitude of show the observed flux distributions: Figure 2 in the op- that star is dimmed by 3^67. The final match is shown in tical region. Figure 3 in the ultraviolet. U Sge is some- Figure 4. The only conspicuous discrepancy occurs near what brighter at all wavelengths. The slope of its Pas- 390 nm, where the standard star has deeper absorption chen continuum is somewhat steeper than that of U Cep, lines than the secondary of U Sge. indicating a higher effective temperature. This conclusion is confirmed by IUE observations, which give VI. Determination of Masses in U Sagittae systematically higher fluxes for U Sge as we go to shorter We decided to determine the masses of the com- wavelengths (Fig. 3). After the correction for interstellar ponents by using the same method as Joy (1930), namely, reddening is applied, this difference is strengthened, obtaining the slope of the radial-velocity curve of the since for U Sge we find E{B — V) = 0^06, while for U cooler star from spectrograms taken during the primary Cep the color excess is E{B — V) = 0^3. A somewhat eclipse, when its absorption lines can be measured. Our greater reddening for U Sge is to be expected, since this new spectroscopic observations used the Varo image star is located closer to the galactic plane (b = +2°) tube attached to the coudé spectrograph of the 3-m than is U Cep (¾ = +17°). Shane reflector of Lick Observatory. The Varo tube We are not sure, though, that this conclusion has gen- shortens the exposure so significantly that good time res- eral validity, for U Cep seems to be variable in almost all olution can be accomplished with spectrograms of fairly respects, and it may appear hotter at other phases or at high dispersion, namely, about 16 Â mm-1. Thus, our other epochs. data are far superior to those of Joy, who used a recipro-

5·5" M fa a I I |N Λ u SAGITTAE

1 I I I I I I I 320 360 400 440 480 520 560 600 640 680 WAVELENGTH (nm) FIG. 2—A comparison of the observed optical fluxes of U Sge and U Cep.

Fig. 3—A comparison of the observed ultraviolet (WE) fluxes of U Sge and U Cep.

8.5 SECONDARY STAR OF U SAGITTAE

ÛÜJ 2 9.0 o< S HR 1327( +3.67 mag) S 9.5 QÜJ ÛLü crLü Q 10.0

10.5

360 400 440 480 520 560 600 640 680 WAVELENGTH (nm) FlG. 4—The optical scan of U Sge in total eclipse shows the flux distribution of the secondary component, which is very well matched by the standard star HR 1327 (G4 III) (represented by open circles). cal dispersion of 75 Â mm-1. The dates and exposures for individual lines and the final corrected heliocentric times are listed in Table III. radial velocities for each plate are listed, respectively, in Our spectrograms cover two spectral regions: one is Table IV (the 520 nm plates) and Table V (Ha plates). centered around 520 nm, the other at Ha. As expected in Table VI then lists the final adopted heliocentric radial late-type stars, absorption lines of neutral metals, in par- velocities of the secondary component of U Sge. ticular Fe i, dominate both regions. No emission lines are A least-square fit to the radial velocities of the sec- visible even at this higher resolution. Because of short ondary has then been obtained, giving the slope of its exposures, on the order of 20 minutes, there is not much radial velocity at primary conjunction. The correspond- line smearing due to the orbital motion. Moreover, line ing amplitude of the radial-velocity curve of the second- blending is not severe, in particular in the Ha region. All ary is K(G) = 206.5 ± 6.5 km s-1. Figure 5 shows this plates have been measured on a Grant machine at fit, the individual observations, and also the slope of the UCLA. About 20 reliable lines could be identified and radial-velocity curve of the primary component, which measured on the 520 nm spectrograms, while on the Ha was calculated from McNamara's (1951a) value of K(B) plates, the number was typically about 10. Each mea- = 70 km s-1. From the ratio of these two slopes we get sured line was weighted according to its profile, the mass ratio sIU(B)/sJlU(G) = 2.95 ± 0.05. As already strength, and distortion. The measured radial velocities mentioned, Tomkin (1979) obtained K(G) = 209 ± 5

TABLE III TABLE IV Coudé Observations of U Sagittae Radial Velocities of U Sge Secondary 5200 Â Region Plates

Date (UT) SPECTRAL EXP lab ID EC 13386 EC 14355 EC 14356 (mid-exposure) 2,440,000+ PHASE PLATE REGION (MIN) 5139.2600 Fe I -7.29 -2.17 1974 Aug 09, 0405 2268.6701 0.991 EC 12589 Ho 27 5142.5431 Fe I -1.98 -6.08 -0.91 1974 Aug 09, 0431 2268.6882 0.996 EC 12590 Ha 17 5151.9182 Fe I -1.51 1974 Aug 09, 0452 2268.7028 0.001 EC 12591 Ha 16 5167.6168 Fe I -0.82 -6.22 -1.00 1974 Aug 09, 0513 2268.7174 0.005 EC 12592 Ha 15 5172.6933 Mg I -2.89 -7.74 -1.53 1974 Aug 09, 0531 2268.7299 0.009 EC 12593 Ha 13 5183.6512 1974 Aug 09, 0547 2268.7410 0.012 EC 12594 Ha 11 Mg I -1.17 -6.05 +0.37 1974 Aug 09, 0600 2268.7500 0.015 EC 12595 Ha 10 5192.3520 Fe I -2.65 -8.19 -0.05 1975 Jul 23, 1019 2616.9299 0.007 EC 13386 5200 Ä 13 5206.0461 Gr I -2.10 -5.71 1976 Aug 08, 0942 2998.9042 0.996 EC 14355 5200 Ä 26 5208.4351 Gr I +0.50 -2.77 1976 Aug 08, 1012 2998.9250 0.003 EC 14356 5200 Â 26 5226.8714 Fe I -2.71 -9.16 +1.22 5250.6527 Fe I -2.63 km s-1, in excellent agreement with our results. Tom- 5265.5605 Ca I -0.85 kin's result probably carries greater weight, so we will 5269.5458 Fe I +2.82 -1.98 +1.40 adopt his mass ratio of 2.99 ± 0.08. The masses of the 5283.6291 Fe I -8.79 +3.05 components are then 5.7 and 1.9 respectively. 5328.5376 Fe I -5.81 -10.78 +0.12 From the photometric studies it follows that the orbital inclination of the system must be quite close to 90°. If 5341.0308 Fe I -1.15 -19.63 -18.61? we adopt this value, then the separation of the com- 5371.4962 Fe I -5.68 -12.64 -7.87 ponents is A = 18.7 fí©. In order to complete the set of 5397.1336 Fe I +2.10 +5.34 parameters for the system, we will now use our spectro- 5405.7838 Fe I -9.06 -19.47 -4.86 photometric scans. 5429.7065 Fe I +0.31 -10.54 VII. The System Parameters: 5446.9238 Fe I -3.74 -8.10 -1.10 Distance, Radii, Luminosities 5455.6176 Fe I +2.73 When matching the observed but dereddened mono-

chromatic magnitudes mx to the model values Μλ = WEIGHTED MEAN VELOCITY -1.02 -8.57 -0.30 — 2.5 logFx, obtained from Kurucz's monochromatic HELIOCENTRIC CORRECTION +2.40 +8.47 +8.50 fluxes, we obtain the ratio R/d of the star's radius to the star's distance CORRECTED RADIAL VERLOCITY mx — Mx — 33.321 5 log{R/d) (3) -3.42±0.50 -17.04±0.60 -8.80±0.40 where R is expressed in solar radii and d in kpc. The procedure is described and the constant derived in Plavec, Weiland, and Koch (1982). Our least-square solution tual uncertainty in the distance determination is larger, gives Rjd — 14.33 ± 0.03 for the hotter component. but we hope that we are within about 10% of the correct We must now determine its radius. value. We will adopt d — 290 pc. Tomkin (1979) derived R(B) = 4.06 R@ based on Ces- The radius of the secondary star follows again from its ter and Pucillo's (1972) fractional radius r(B) = 0.218, fractional radius, for which Cester and Pucillo (1972) ob- and on the orbital radius A = 18.7 R@. This leads to a tained r(G) = 0.288, and McNamara and Feltz (1976) distance of 283 pc. If we adopt the fractional radius r(B) gave r(G) = 0.296. Thus, we get R{G) between 5.4 and = 0.225 according to McNamara and Feltz (1976), we 5.5 Rq. The fractional radii mentioned here are actually get R{B) = 4.15 fí®, and the distance will be 290 pc. Fi- the photometric radii, or essentially the length of the h- nally, using our results from atmospheric model fitting, axis in a distorted star. If we assume that the secondary lögg = 3.90, and the mass of the primary star of 5.5 star fills its critical Roche lobe, then for the given mass W©, we obtain R(B) = 4.21 ± 0.3 fí©, and a distance d ratio the tables by Plavec and Kratochvil (1964) give = 295 ± 20 pc. The mean errors shown here express R{G) = 5.3 Rq. Thus we conclude that the secondary only the uncertainty in log g and ^(B). No doubt the ac- star in U Sge does indeed fill its critical Roche lobe.

TABLE V +40 Radial Velocities of U Sge Secondary Ha Region Plates Alab ID EC12589 EC12590 EC12591 EC12592 EC12593 EC12594 EC12595 6393.6113 Fe I -15.64 +0.98 +2.34 +6.53 +10.45 +12.68 6400.0101 Fe I -4.62 2.58 -1.17 +1.57 +12.19 +12.29 +23.52 6408.0262 Fe I -13.08 5.99 -1.52 +1.23 +4.40 +9.54 6420.6659 Fe I 6.58 -2.18 +1.60 +11.29 +12.15 +21.31 6439.0851 Ca I -20.11 0.12 +1.13 +1.96 +9.91 +15.15 +15.72 6449.8836 Ca I -17.21 9.42 +1.39 +0.67 +7.91 +12.00 +17.79 6456.3760 Fe 11 1.85 +0.89 +0.21 +7.51 +12.26 +18.09 6462.6222 Fe I -8.76 9.59 +0.65 +7.41 +18.51 6475.6318 Fe I 7.83 +1.06 +13.81 6482.8075 Fe I 5.97 +2.24 +10.50 6574.2325 Fe I +2.55 +10.98 6592.9221 Fe I -10.19 3.07 +2.87 +10.26 6609.1189 Fe I 8.26 6663.3900 Fe I 2.65 +1.48 +11.23 6677.9958 Fe I -14.18 2.07 +3.19 +10.33 WEIGHTED 0.01 0.02 MEAN VELOCITY -13.26 -7.61 +0.02 +1.67 +9.31 +11.96 +17.86 HELIOCEN. CORR. +8.09 +8.13 +8.16 +8.20 +8.23 +8.26 +8.28 PHASE —► CORRECTED Fig. 5—Radial velocities of both components of U Sge in the vicinity RADIAL VELOCITY -21.35 -15.74 -8.14 -6.53 +1.08 +3.70 +9.58 of the primary eclipse, as measured on spectrograms obtained at Lick -1.50 ±1.40 -0.50 ±0.40 ±0.70 ±0.90 ±1.40 Observatory. ue of V0(G) = 8^7. Using our adopted distance of 290 pc, we find that the TABLE VI absolute visual magnitudes of the two stars are M {B) — Velocities of U Sge Secondary Summary of Results V — IW, My(G) = +1¾. There appears to be no diffi- JD SPECTRAL RADIAL culty in classifying the cooler star as G4 III-IV. Thus the PLATE 2,440,000+ PHASE REGION VELOCITY (km/sec) cooler component lies between genuine subgiants and gi- ants in luminosity. Our collection of scans of standard EC 12589 2268.6701 0.991 Ha -21.4 ± 1.5 stars does not contain a genuine subgiant at this spectral EC 12590 2268.6882 0.996 Ha -15.7 ± 1.4 type, so that a small degree of uncertainty in both pa- EC 12591 2268.7028 0.001 Ha -8.1 ± 0.5 rameters remains. The hotter star appears to be somewhat too luminous for a main-sequence star of spectral Ha -6.5 ± 0.4 EC 12592 2268.7174 0.005 type B7.5. In the calibration by Blaauw (1963) its lumi- EC 15293 2268.7299 0.009 Ha +1.1 ± 0.7 nosity class is IV, while according to Straizys and Kuri- EC 12594 2268.7410 0.012 Ha +3.7 ± 0.9 liene (1981), it would be III. The primary star in U Sge EC 12595 2268.7500 0.015 Ha +9.6 ± 1.4 may indeed be somewhat evolved within the main-sequence band. This is also indicated by the value log g = EC 13386 2616.9299 0.006 5200 Ä -3.4 ± 0.5 3.9 we obtained from model-atmosphere fitting. A lumi- EC 14355 2998.9042 0.996 5200 Â -17.0 ± 0.6 nosity class V star is expected to have logg = 4.05 at EC 14356 2998.9250 0.003 5200 Ä -8.8 ± 0.5 this spectral type (after Straizys and Kuriliene). In spite of these and similar remaining uncertainties, we conclude that at the present time, U Sge is a well- Our optical scans made at Lick Observatory can also known system, and we list all its parameters in Table be used to obtain the UBV magnitudes of the two com- VII. All the above values are in excellent agreement ponents. The dereddened V magnitudes are found to be with those listed by Popper (1980, Table 12). V0{B) = 6^31, V0(G) = 8^86. The V magnitude for the primary star is, within the observational errors, indepen- VIII. Photospheric Line Profiles dent of orbital phase. The V magnitude for the second- Chemical abundances in the system of U Sge have ary star has been derived for phase 0.0; rotation of this been studied by applying spectrum synthesis to a high- distorted star will change this value within a range of dispersion IUE spectrum, taken with the short-wave- about 0^25. Since we are unable to represent this varia- length (SWP) camera. Several spectral regions with well- tion satisfactorily, we will adopt an estimated mean val- defined absorption lines were selected for this purpose.

TABLE VII TABLE VIII Parameters of the System of U Sagittae Spectrum Synthesis Regions

SYSTEM ECHELLE ORDER WIDTH (Â) MAJOR LINES γ (km/sec) -10.1 a e 0.03 a 17 1256.6 - 1268.3 Si II sin 1 1.00 a 20 1292.2 - 1306.5 Si II, Si III f(m) 0.1187a m(B)/m(G) 2.99 23 1330.0 - 1340.0 C II 18.7 34 1490.0 - 1501.0 Ν I Distance (pcs) 290 36 1521.3 - 1537.0 Si II, Fe II E(B-V) 0.06 44 1672.0 - 1679.0 Fe II Κ (km/sec) 70 a 206 46 1713.0 - 1732.0 Fe II, Al II V 6.50 9.05 Vo 6.31 8.86 (B-V) -0.06 0.86 broadening parameter. It is known that it is practically (B-V)o -0.12 0.80 impossible to distinguish between macroturbulance and Τ (0K) 12,250 5,250 axial rotation in spectra with fairly high but not ex- log g 3.90 3.25 tremely high dispersion. In the case of binary stars, the dominant broadening agent is rotation. Mv -1.04 +1.52 The rotational velocity of the primary component of Mass (Μ^η) 5.70 1.90 U Sge has been studied several times in the past. Struve Radius CRg\in^ 4.21 5.30 (1949) observed the rotation effect on the radial-velocity Spectral Type B7.5 IV-V G4 III-IV curve of the hotter star, and from his data, Plavec _1 a (1967b) derived ^ 90 km s . From the line profiles Values adopted from McNatnara (1951a) νΓ0ΐ in the optical region Horn (cf. Plavec 1967¾) obtained The selected regions, their widths, and major absorption 80 km s-1, while Koch, Olson, and Yoss (1965) found 76 lines identified in these regions are listed in Table VIII. ± 12 km s_1. Olson (1984) reports that his spectrograms For spectrum synthesis, we used the Kurucz model at- taken in 1968 and 1970 give a rotational velocity of 69 _1 mosphere (bilinearly interpolated) with Teff = 12,250 Κ ± 9 km s . and logg = 3.90, obtained from matching the observed The UCLA spectrum synthesis program employs flux distribution as described in section III. The micro- macroturbulence as the parameter which broadens a line turbulent velocity parameter was set by the requirement without changing its equivalent width. Only for un- that ionic abundances should show no dependence on blended lines is it possible at the present to calculate line equivalent widths. Very suitable for this purpose are rotational broadening. A comparison of the two ap- the numerous strong and weak Fe n lines in orders 36, proaches shows that the broadened line profiles are es- _1 44, and 46 of the echelle spectrogram. Their solution sentially identical in U Sge if either Vturb = 80 km s _1 _1 converged to a value of 8 km s . This result was then or Vrot ^ 100 km s . Thus, our data give us a somewhat independently checked on the Si π and Si m lines in or- higher rotational velocity of U Sge than the sources men- ders 17 and 20, with identical results. Thus, 8 km s-1 was tioned above. adopted for the microturbulent velocity parameter. Surprisingly, Olson (1984) very recently reported that When this value is used, the abundances of iron and sili- his new spectrograms, obtained in 1982, show that the con equal their solar values: [Fe/H] = 0.00 ± 0.15 and rotational velocity of U Sge is now higher than before: [Si/H] = 0.00 ± 0.1. Adopting a lower microturbulence his new value is 99 ± 8 km s-1. One could not wish for parameter would lead to an overabundance of Si and Fe, a better agreement. while a higher microturbulence parameter would re- The synchronized velocity would be 62.7 km s-1, so quire an underabundance of these two elements, com- the primary component rotates somewhat faster than pared to the sun. synchronism would require, but not by a large factor While the choice of the microturbulence parameter (Vrot/Vgyn = 1.58). This is a typical situation in Algol satisfactorily adjusts the theoretical equivalent widths so variables; in fact, in some of them, the hotter component that they match the observed ones, the fitting of the ac- rotates much faster than would correspond to synchro- tual line profiles requires the choice of a dominant line- nism (such is the case of U Cep).

With all the atmospheric parameters fixed as de- IX. Discovery of Ultraviolet Emission Lines scribed above, the spectrum was synthesized in seven re- in U Sagittae gions of the high-dispersion SWP spectrum. The line The reversed abundance ratio C/N in the photo- identification lists, excitation potentials, and values of spheres of the components of U Sge, as well as the gen- logig/) were all taken from the compilation by Kurucz eral character of the system and the fact that the less and Peytremann (1975). The only adjustable parameter massive star fills its critical Roche lobe, all corroborate left is the abundance of the element in question. At first, the idea that the system underwent a large-scale mass we used a version of the program that predicts the line transfer (and possibly mass loss as well) in the past. Is this strengths for the adopted parameters and solar abun- process still going on (albeit plainly on a much dimin- dances. Only those lines that have equivalent widths Wx ished scale)? Our spectrophotometric scans have not re- > 0.01 Â were then used to generate the synthetic spec- vealed evidence of circumstellar material, but other ob- tra and match them to the observed line profiles. As an servers have been finding various rather subtle illustration of these fits, we show in Figure 6 the match manifestations of it, as already mentioned in section I. It for order 17, with two strong Si π lines, and in Figure 7, is also important to determine the physical state of such the match for order 20, with a mixture of Si π and Si m material. It has recently been recognized that the gainers lines. The only striking discrepancy in the latter figure is in active Algol variables are surrounded by a hot turbu- due to the interstellar Ο ι line at 1302 Â. These illustra- lent region (Plavec 1983; Peters and Polidan 1984). It is tions as well as our other synthesized regions show rather important to check if emission lines of a hot plasma are convincingly that the spectrum of U Sge we used was present in U Sge too. largely free from circumstellar absorption lines. All the For this purpose, U Sge was observed with the IUE matching was done on the assumption that the abun- spectrograph at the time of its total eclipse on 1983 No- dances of elements are the same in U Sge as in the sun. vember 7 by Plavec and Janet L. Weiland. The SWP This assumption works very satisfactorily, except as 21468 spectrum we obtained does show a wealth of noted below, and we conclude that the elemental abun- emission lines, essentially similar to those found in the W dances are, in general, solar: [metals/H] = 0.00 ± 0.2. Serpentis stars such as SX Cassiopeiae (Plavec et al. The only striking exception is the line profile of the 1982) and in the Algol variables U Cep (Plavec 1983), C π resonance doublet at 1336 Â, which we found to be RW Tauri (Plavec and Dobias 1983), TT Hydrae, RY narrower than the synthetic profile based on solar abun- Persei, and UX Monocerotis (for a preliminary an- dance of carbon. It is very tempting to interpret this re- nouncement, see Plavec et al. 1984). In fact, as Figure 8 sult to mean that carbon is underabundant. However, the shows, there appear to be more emission lines present oscillator strengths are currently being revised, so that than have been listed for SX Cas in Plavec et al. (1982). no reliable statement can at present be made about the However, the signal-to-noise ratio is so poor that we pre- carbon abundance in U Sge. fer to postpone the identification of uncertain lines until more and better data are accumulated. Such an im- proved observation is feasible. Our 30-minute exposure covered the phase intervals 0.9993 Ρ through 0.0055 P,

CO Lüζ

< ο CO < 3O

0.0 1257 1259 1261 1263 1265 1267 WAVELENGTH (A) 1292.5 1294.5 1296.5 1298.5 1300.5 1302.5 1304.5 FlG. 6—The observed flux distribution of U Sge between WAVELENGTH (A) 1257 A-1267 A, as obtained from a high-dispersion 1UE spectrum (weak line) is matched hereby a synthesized spectrum (heavy line). Fig. 7—Similar to Figure 6, showing the region 1292 Á-13()4 Á.

1250 1300 1350 1400 1450 1500 1550 1600 1650 1700 1750 1800 1850 1900 1950 WAVELENGTH(Â)

Fig. 8—Far ultraviolet (SWP) spectrum of U Sge observed with the IUE satellite during total eclipse on 1983 November 7.

according to Olson's ephemeris, equation (1). The total TABLE IX Prominent Emission Lines Observed in U Sagittae phase of the primary eclipse lasts, according to the light elements adopted here, about 97 minutes, between χ -13ΜΨ0 obs (10 ν - P(US) P(UC) P(UC)/ Xlab LEP UEP phases 0.990 Ρ and 0.010 P. Thus an exposure time three (A) cm- s- ) (10- Lo) P(US) (A) (eV) (eV) times as long as the one we used is feasible. The ephem- 1240 2.9 Ν V (1) 1238. 0.0 10.0 eris (1) appears to be accurate, and our fears of missing Ν V (1) 1242. 0.0 10.0 the totality phase were not substantiated. Avoiding the 1253 1.1 2.9 7.7 2.7 S II (1) 1250. 0.0 9.8 S II (1) 1253. 0.0 9.8 light of the hotter component is essential for successful 1261 1.8 4.7 7.1 1.5 Si II (4) 1260. 0.0 9.8 detection of the emission lines, since they are quite Si II (4) 1264. 0.0 9.8 Si II (4) 1265. 0.0 9.8 weak. As Figure 8 shows, the peak fluxes in the strongest 1296 2.45 6.4 7.4 Si III (4) 1296. 6.6 16.1 emission lines (with the exception of C iv at 1550 Ä) do Si III (4) 1298, 6.6 16.1 0 I (2) 1302, 0.0 9.5 not even reach 10~13 ergs cm-2 s-1 Â-1. At 1400 Â, the -11 1302 2.8 7.4 15.0 Si II (3) 1304, 0.0 9.5 flux from the hotter star is about 7 χ 10 in the same 0 I (2) 1304, 0.0 9.5 units, i.e., the hotter component radiates about 800 times 1335 4.1 10.9 21.6 II (1) 1334, 0.0 9.3 1335, 0.0 9.3 more flux per angstrom than the strong Si iv emission II (1) lines. Accurate timing is very essential for the detection 1355 2.6 6.7 of the emission lines; this once more stresses the impor- 1394 4.8 12.7 18.4 1.5 Si IV (1) 1393.8 0.0 tance of regular observations of minima of eclipsing 1403 3.4 8.9 25.6 2.9 Si IV (1) 1402.8 0.0 1533 2.4 6.2 16.6 1.5 Si II (2) 1526.7 0.0 8.1 binaries. Si II (2) 1533.4 0.0 8.1 A preliminary list of emission lines found in the 1549 10.4 27.3 67.2 C IV (1) 1548.2 0.0 8.0 C IV (1) 1550.8 0.0 8.0 eclipse spectrum of U Sge is given in Table IX. Next to 1565 3.0 7.8 4.0 0.5 Si 11(10. 1563.8 6.9 14.8 the observed wavelength, Aobs, we give the observed (but 02) dereddened) flux F(US) in units of 10-13 ergs cm-2 s-1 1671 1.8 4.8 13.8 2.9 Al II (2) 1670.8 0.0 7.4 _1 1785 1.7 4.4 5.9 1.3 Ni II (5) 1788.5 0.2 " 7.1 Â . It will be noticed that the observed flux is quite 1811 2.2 5.8 7.0 1.2 Si II (1) 1808.0 0.0 6.9 low. The visual impression of the spectrum also was that Si II (1) 1816.9 0.0 6.9 the emission lines are extremely weak. However, when 1854 4.0 11.0 22.2 Al III (1) 1854.7 0.0 6.6 we apply the reddening correction and take into account 1862 1.4 3.6 16.3 Al III (1) 1862.8 0.0 6.6 the larger distance to U Sge (as compared to U Cep) in 1895 10.2 Si III (1) 1892.0 0.0 6.6 calculating the total emitted power P(US) in Table IX, Fe 111(34) 1895.5 3.7 10.2 we find that the power received from the emission lines 1915 2.87 8.9 Fe 111(34) 1914.1 3.7 10.2 in U Cep under similar circumstances, P(UC) in Table 1926 8.6 Fe 111(34) 1926.3 3.7 10.1 IX, is only 2.0 times stronger, on the average. This is surprising, since all other indicators suggest that U Cep is much more active and its current mass-transfer rate ap- tween the two stars. It appears that the lines of Si π are pears to be definitely higher. relatively stronger in U Sge than in U Cep; in particular, There are certain differences in line intensities be- the lines at 1527 Â and 1533 Â are quite strong in U Sge.

Contrary to that, the usually strong blend of Si m (1) and tems. Wilson and Twigg (1980) have recently argued—in Fe m (34) at 1893 Â to 1895 Â is surprisingly absent in U our opinion rather convincingly—that rapid rotation of Sge. A better spectrum is strongly desirable. The identi- the accreting star plays an important role in the evolu- fications for the more prominent lines, as well as their tion of the mass-transfering systems. In this respect, the lower (LEP) and upper (UEP) excitation potentials are very recent finding by Olson (1984) may be of great im- given in Table IX. portance, namely, that the primary component of U Sge accelerated its rotation from about 70 km s-1 to about X. Discussion 100 km s_1 within only about twelve years. This fact, to- U Sge is in several respects almost a twin system to U gether with the detection of shell absorption lines by Cep. Both objects have about the same apparent bright- McCluskey and Kondo (1984) and our discovery of cir- ness, and consist of rather similar stars. It also follows cumstellar emission lines may herald an increased activi- that they are approximately at the same distance from ty in U Sge. Further observations of U Sge and other Al- us. The period of U Sge is longer than that of U Cep by gol variables are very desirable. a factor of 1.36, but this is about the largest difference The authors are grateful to Janet L. Weiland for coop- between the parameters of the two systems. It is surpris- eration in the observations, and to her as well as to Paul ing, therefore, that their behavior is so different. U Cep B. Etzel, Dr. Charles D. Keyes, and Dr. Zdenka Plavec is renowned for its considerable activity, manifested in for cooperation in the development of various computer the variability of the period, appearance of emission programs. Our thanks are also due to the Directors and lines, formation of transient disks, etc. (see, e.g., Olson Staffs of the NASA IUE Observatory and of Lick Obser- 1980). U Sge actually displays all these phenomena, too, vatory. Dr. D. M. Popper critically read the manuscript, but on a much more modest scale. Why? And why is the and Robert L. O'Daniel assisted in the final editing of primary component in U Sge rotating just a little faster the manuscript. We also wish to thank Dr. Edward C. than at the synchronized velocity, while in U Cep the Olson for his constructive comments and suggestions. primary rotates about five times faster? The answer is, of This project has been supported by grants NSF AST-82- course, that U Cep is in a more active stage of mass 00046 and NASA NSG-5275. transfer. But why? Which of the parameters decides that? We pose this question here without suggesting any definite answer. One possible explanation was offered by REFERENCES Tomkin (1979), who points out that the mass ratios (in Al-Naimiy, H. 1978, Ap. Space Sei. 56, 219. the sense hotter/cooler component) in Algol, U Sge, and Blaauw, A. 1963, in Basie Astronomieal Data, K. Aa. Strand, ed. (Chi- U Cep decrease in this order, while the activity in- cago: University of Chicago Press), p. 383. Cester, B., and Pucillo, M. 1972, Mein. Soc. Astron. Ital. 43, 501. creases. The idea is that the higher the mass ratio, the Cester, B., Fedel, B., Giuricin, G., Mardirossian, F., and Pucillo, M. more advanced is the system in the mass-transfer pro- 1977, Astr. Ap. 61, 469. cess. However, we believe that the dependence is not so Hall, D. S., and van Landingham, F. G. 1970, Pub. A.S.P. 82, 749. simple, since model calculations (Horn, Kfiz, and Plavec Hayes, D. S., Latham, D. W., and Hayes, S. H. 1975, Ap. J. 197, 587. Horn, T., Kfiz, S., and Plavec, M. 1970, Bull. Astr. hist. Czechoslovakia 1970) show that the final parameters depend on several 21, 45. initial parameters of the system. Joy, A. H. 1930, Ap. J. 71, 336. We should keep in mind that the time interval cov- Koch, R. H., Olson, E. C., and Yoss, Κ. M. 1965, Ap. J. 141, 955. ered by observations is so short that there is no guaran- Kondo, Y., McCluskey, G. E., and Wu, C. C. 1981, Ap. J. Suppl. 47, tee that we observe real secular changes. The total mass 333. Kopal, Z., and Shapley, M. B. 1956, Jodrell Bank Ann. 1, 141. of the circumstellar material represents, even in U Cep, Kurucz, R. L. 1979, Ap. J. Suppl. 40, 1. such a negligible fraction of the mass of the system that Kumcz, R. L., and Peytremann, E. 1975, SAO Special Report No. 362. all the activity observed in U Cep may be only an in- McCluskey, G. E., and Kondo, Y. 1984, in The Future of Ultraviolet As- significant temporary perturbation, and that 50 or 100 tronomy, R. D. Chapman, J. Meade, and Y. Kondo, eds., xVASA Con- years hence the comparison between U Cep and U Sge ference Report (in press). McNamara, D. Η. 1951a, Ap. J. 114, 513. might give quite a different picture. McNamara, D. H. 1951^, Pub. A.S.P. 63, 38. Perhaps the very rapid rotation of U Cep is the McNamara, D. H., and Feltz, K. A. 1976, Pub. A.S.P. 88, 688. strongest argument in favor of the hypothesis that U Cep Naftilan, S. A. 1975, Ap. J. 206, 785. is closer to the rapid phase of mass transfer than U Sge. Olson, E. C. 1975, Ap. J. Suppl. 29, 43. The rapid rotation of U Cep is probably something that Olson, E. C. 1980, Ap. J. 237, 496. Olson, E. C. 1982a, Pub. A.S.P. 94, 70. can change only on a longer time scale. If so, the ob- Olson, E. C. 19825, Ap. J. 257, 198. served phenomena may have a more lasting character Olson, E. C. 1984, Pub. A.S.P. 96, 376. and deeper influence on the evolution of the two sys- Parthasaranthy, M., Lambert, D. L., and Tomkin, J. 1983, M.N.R.A.S.

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