HD5980: is it a real LBV? Leonid Georgiev 1, Gloria Koenigsberger 2, John Hillier 3 1. Instituto de Astronomía, Universidad Nacional Autonoma de México; 2. Instituto de Ciencias Físicas, Universidad Nacional Autonoma de México; 3. Department of Physics and Astronomy, University of Pittsburgh, USA. Introduction: It is believed that the LBV’s represents a short evolution phase between the O stars and the WR stars. They experience a violent increase of their mass loss rate and brightness. In some cases the changes in the brightness are related only to changes in the temperature at constant luminosity. This points to a variability which is not related to changes of the energy generation processes but rather to adjustments of the exterior envelope. In other cases the luminosity of the star increases during the outburst which needs an additional energy source. Therefore the main expected difference between giant LBV eruptions and other types of variability with similar characteristics may be tested by the constancy of the luminosity during the events.. HD5980 is a triple system which contains a binary star with a period of 19 d The star which is closer to the observer at phase 0 is called “A” and the other is called “B”. In addition, the spectrum of the system shows absorption lines which are not related to the orbital motion. We associate these lines to a third star, called “C”. It is not clear yet if this star is bound to the A+B system. Was the 1994 event a mayor eruption? The LBV outburst are divided in two main groups: 1) Mayor eruptions, like in eta Car. 2) A normal ``normal” eruptive events. The primary difference is the changes in the stellar luminosity. The normal variability is related to a change in the mass loss rate, the temperature of the star or a combination of the two, caused by some instability mechanism related to the wind driving force, but not to the structure of the star. On the other hand, the energy released during the mayor o giant eruptions requires an additional source and it is related to changes in the stellar luminosity. Thus, the clue to determining the nature of the 1994 eruption lies on the conservation of the luminosity during the event. We estimate the luminosity at two epochs: on December 1994 and April, The spectra obtained in the two epochs cover the region from 1150A to 7000A (1994) and from 1150A to 10000A (2000). The wide spectral coverage allow us to compare the observed and calculated line profiles and continuum shape. We fixed the contribution of stars B and C as deduced above and fitted the parameters of star A using CMFGEN code. Fig. 3 show the calculated spectra of the two preliminary models compared to the observations and Table 1 shows the used parameters. Is the eruptive event unique? We checked the old observations obtained between 1955 and 1962 in Radcliff Observatory, South Africa. Fig. 2 below shows the comparison between the spectra obtained in 1961 and The similarity is striking. At that time, star A was a WN6 star as it is now. We do not see any absorption lines and the rest of the spectra in the series show only one WN star. The Hbeta radial velocities follow the motion of star A, so we conclude that star A was the one seen in these spectra in its WN6 phase. We do not see any impressive increase of the mass loss during this epoch, but the 2005 spectrum is taken 11 years after the eruption, so if the 1961 spectrum represents the same phase, the eruption may have happened in 1950, before the first spectrum in the collection. This leads us to speculate that Star A may have undergone another mass-loss rate increase about 44 years ago, leading to the question of whether the eruptions may be periodic?. We can’t answer that question now. But if it is the case, in 2009, we may expect to observe a spectrum similar to the one obtained in 1965 (Fig 2). Conclusion 1: The spectral evolution of Star A. In this work we added two new facts to the HD5980 puzzle. Based on spectra obtained in the late 1970s, the system was classified as OB+WNE with Star A believed to be an O supergiant. During the 1980s, Star A gradually increased its mass loss rate and the degree of ionization and was classified as WN6. This phase continued until 1993 when the temperature (degree of ionization) dropped to WN8 and then, in August 1994, the star increased its brightness and again dropped its degree of ionization to WN11 with spectral features also similar to those of a B1.5Ia+. Shortly after that the star returned to the WN6 phase (see Koenigsberger 2004 for a review of the system’s changing parameters). The new piece of information in this list is the ionization state of the star in the late 1950s and early 1960s. The spectra show a star with strong emission lines. The spectrum in 1961 is almost identical to the spectrum in There are no absorption lines seen above the noise level. The strength of the emission lines drops significantly after 1963 with many of them disappearing by The spectra taken in the 1970s show only HeII 4686 and several absorption lines. Based on the above, we conclude that Star A may periodically increase its mass loss rate. The 1994 eruption may have been one of several that have occur every 40 o so years. Given the similarity of the 1961 spectrum to that obtained in 2005, one might expect that by next year the star will have decreased its mass loss rate to its value in the 1970s, thus allowing a clearer view of star B. Contribution of Star B and C to the spectrum of the system. The key point of the interpretation of the system is the contribution of Stars B and C to the system’s spectrum. The signature of Star C in the spectrum is a system of absorption lines observed at the center of the H and He lines. Star B is more complicated. The only clear signature of this star is the NV 4604/4620 doublet. We used a model of a typical O7 supergiant for star A and a WN4 model for star B, made with CMFGEN (Hillier & Miller,1998), as first approximation for these two stars. We determined the relative contribution of the three stars by fixing Star A’s spectrum and added different percentages of Star B and C until the HeI 4471 absorption and NV doublet of the model fit the observations (Fig. 1). The lines are reproduced with star C contribution 15% of the light. Star B is less defined, but is seems that a solution between 1/4 and 1/3 of Star A’s flux at 5000A is reasonable. This solution of the relative fluxes in the system is in agreement with the light curve solution of Breysacher & Perrier (1991). Table 1. Star A parameters at three different epochs Conclusion 2: Luminosity change during the eruption. The wind models of HD 5980 for spectra obtained at two epochs after the eruption suggest that Star A changed its size, mass loss rate and, most importantly, its luminosity. Adding the results of Drissen et al. (2001) we conclude that Star A seems to have increased its luminosity during the eruption, then rapidly (in ~6 months) decreased it by a factor of 3 and then by another 50% during the next 6 years. The increase of the luminosity requires an additional source of energy. One possible mechanism triggering the additional amount of energy might be the tidally induced instability by the interaction between star A and B (Koenigsberger & Moreno, 2007) Fig. 1. Comparison between the observed spectrum and the sum of the three models for stars A, B and C. Star A is the brightest in the visual spectrum. Fig. 2. Comparison between old spectra obtained in early 60ties and the new one taken in It is obvious that in 1961 the star were in the same stage as in Note that emission lines disappeared in Fig. 3. Comparison between the observed and model spectra for two epochs: 1994 (observed in orange, model in green) and 2000 (observed in black, model in blue). Bibliography. Breysacher, J, Perrier, C, 1991, IAU Symp. 143, 229. Drissen, L., Crowther, P. A., Smith, L.J.,Robert, C., Roy, J.-R.,Hillier, D.J., 2001, ApJ, 546, 484. Hillier, D.J., Miller,D, 1998, ApJ, 496,407 Koenigsberger, G., 2004, Rev. Mex. A.A., 40, 107. Koegigsberger, G., Moreno, E., 2008, Rev Mex A.A., 33, 108.