The Astronomical Journal, 128:2854–2861, 2004 December # 2004. The American Astronomical Society. All rights reserved. Printed in U.S.A.
HUBBLE SPACE TELESCOPE NICMOS VARIABILITY STUDY OF MASSIVE STARS IN THE YOUNG DENSE GALACTIC STARBURST NGC 3603 A. F. J. Moffat, V. Poitras, and S. V. Marchenko1 De´partement de Physique, Universite´ de Montre´al, C.P. 6128, Succursale Centre-Ville, Montre´al, QC H3C 3J7, Canada;
[email protected],
[email protected],
[email protected]
M. M. Shara and D. R. Zurek American Museum of Natural History, Central Park West at 79th Street, New York, NY 10024;
[email protected],
[email protected]
E. Bergeron Space Telescope Science Institute, 3700 San Martin Drive, Baltimore, MD 21218;
[email protected]
and E. A. Antokhina Sternberg Astronomical Institute, Moscow State University, Universitetskij Prospekt 13, Moscow 119899, Russia;
[email protected] Receivved 2004 April 27; accepted 2004 September 8
ABSTRACT We have used the relatively long data string of the 1997–1999 NICMOS focus tests on NGC 3603 to extract J-band light curves for several hundred stars in the cluster core. Given the relatively modest photometric precision [(J ) 0:05 mag], we were able to isolate only a half-dozen variable candidates with peak-to-valley amplitudes above 0.2 mag. One of the variables is one of the two outstandingly brightest cluster members, A1, located in the very dense cluster center. A1 shows double eclipses on each orbital cycle, with the same period (P ¼ 3:7724 days) as found previously and independently in unresolved ground-based radial velocity variations of the Wolf-Rayet (WR) emission component in the central core of NGC 3603. Very rough best estimates for the masses of the components of A1 are in the range 30–90 M for the brighter and more massive H-rich WR component (WN6ha) and 25–50 M for its assumed O star companion. A more detailed study is urgently needed, given the potential for this extremely luminous system to harbor the most massive main-sequence star ever ‘‘weighed.’’ Another variable, HST 12, escaped the original search, which was based on larger than average standard deviation. It is a probable field-star eclipsing variable with a moderately long period. Key words: binaries: eclipsing — open clusters and associations: individual (NGC 3603) — stars: early-type — stars: variables: other — stars: Wolf-Rayet
the total light output of the cluster. They are of type WR, located within 100 of the center of the cluster, the total radial extent of which is about 10 (2 pc). More precisely, these three stars are of subtype WN6ha in the spectral classification system of Smith et al. (1996). This, combined with their dominating visual brightness, indicates that they may be extremely luminous, relatively H-rich, main-sequence stars that generate strong, hot, WR-like winds (Crowther & Dessart 1998). These ( pseudo-) WR stars are followed at progressively fainter levels by at least a dozen O3–O4 and many other OB-type member stars (Drissen et al. 1995), right down at least to the faintest limit of 0.1 M stars (Brandl et al. 1999). Less frequent images in the H and K bands were also secured with NICMOS during the same interval as the J-band images and could in principle be used along with the J band to construct color-color and color-magnitude diagrams (CMDs). However, this has already been done using large ground-based telescopes and near-IR photometry with adaptive optics (AO), reaching much fainter than these HST test data (Eisenhauer et al. 1998; Brandl et al. 1999), and is not repeated here. In this paper, we concentrate on the analysis of the variability among NGC 3603’s brightest stars in these unique timesampled data. As we show, the minimum (instrumental) scatter in the derived magnitudes is close to (J ) ¼ 0:05 mag, so one can really only probe those stars that show ‘‘nontrivial’’ variability, e.g., deeply eclipsing binaries and nonmicrovariability
1. INTRODUCTION The very young (1 Myr), dense open cluster NGC 3603 (distance 7 kpc, mean extinction AV ¼ 4:5 mag; Moffat 1983; Melnick et al. 1989; Drissen et al. 1995; Brandl et al. 1999) was chosen for approximately weekly measuring and testing of the focus for the NICMOS camera on board the Hubble Space Telescope (HST ) during 1997–1999. This occurred early in the maiden NICMOS mission after a thermal short occurred inside the solid cryogen dewar, resulting in a more rapid warm-up and a change of focus with time. NGC 3603 offers the tremendous advantage for this purpose of containing a large number of resolved IR-bright stars in the NICMOS field, so that focus measurements could be readily made across the whole field. The OB/Wolf-Rayet (WR) stars in this distant Galactic cluster appear IR-bright because of the combination of high intrinsic stellar luminosity and relatively high reddening. The repeated J-band images (we refer to them simply as J-band images from now on; see below) also provide a unique data set to study the variability behavior of some of the most luminous and thus potentially most massive stars known, all of which are in NGC 3603. The visually brightest three member stars (V 11 12, corresponding to J 7 8 mag) dominate 1 Current address: Department of Physics and Astronomy, 1 Big Red Way, Western Kentucky University, Bowling Green, KY 42101-3576.
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Fig. 1a Fig. 1b F i g . Fig. 1.—(a) 1aFig. NICMOS 1b mean superposed J-band chart; (b) a zoom of the central cluster core. Stars listed in Table 1 are identified. The circle of diameter 19B9 (assuming 0B075 pixel1) includes all stars seen on all 43 in-focus J-band images.
in luminous blue variables and other possible unstable stars with ranges in variability greater than 0.2 mag (i.e., 3 ). 2. OBSERVATIONS During its first mission with the unfortunate thermal short, NICMOS was used with the repeated focus tests on NGC 3603 over a span of 671 days, from UT day 63 of 1997 through UT day 4 of 1999, i.e., HJD 2,450,512.2–2,451,183.2. For this study of NGC 3603, we use only the best NIC2 image of each group of focus sweeps, on 43 epochs in the J band (actually with the filter F110W, which is slightly bluer than the true J band). NIC2 images were also obtained in F160W (H band) at 0B075 pixel1, but only at one epoch (along with F110W), as were images in F171M and F180M. These were not used, however, in view of the deeper ground-based AO study of Eisenhauer et al. (1998). Instead, we concentrate on the variability in one filter (J band) in which the most repeated data were obtained. Images were also obtained for NGC 3603 using the NIC1 camera of NICMOS (0B043 pixel1) in a smaller field. The 31 epochs covered in F095N were mostly very close in time to the J-band images obtained with NIC2. Along with the images obtained during the relatively short NIC3 campaigns, they are not considered here. 3. DATA REDUCTION All of the in-focus images were calibrated using the standard CALNICA data pipeline with the best available reference files at the time. Because of the slow temperature increase of the NICMOS detectors over the operational lifetime of the solid cryogen (from 61 to 63 K), a small correction was applied to account for the induced change in quantum efficiency at each epoch. Since the analysis presented here is primarily differential, this would be relatively inconsequential, but the correction was made nonetheless.
Figure 1 shows a superposition of all 43 in-focus J-band images, each rotated and shifted to coincide into one unique orientation and position on the sky. The ragged outer edge is due to the rotation of the approximately square NIC2 field during a yearly cycle. The circle of diameter 19B9 (assuming 0B075 pixel1) shows the limit within which all stars are seen 43 times, whereas the region between the circle and the outer ragged limit includes stars that were imaged only 1–42 times. We restrict this study mainly to stars within the circle for the sake of uniformity and highest precision. Stellar magnitudes on the 43 best-focus J-band images were obtained independently by two of the authors (A by E. B. and B by D. R. Z.), using two variants of the same point-spread function (PSF)-type photometry by each author, respectively: (1) DAOPHOT within IRAF in 1999, with spatially variant aperture corrections using Tiny Tim, calibrated in J-band magnitudes, and (2) P. Stetson’s 2000 version of DAOPHOT, yielding magnitude differences (J ) relative to approximately a dozen isolated stars of constant magnitude. We use B as the prime source here. However, as we see in x 4, the additional source A is very useful in conjunction with B to find likely intrinsically variable sources. 4. ANALYSIS As is usually the case in variability studies like this in which the exposure time is the same for all stars regardless of their brightness, the biggest problem is that of separating the intrinsic and the instrumental (i.e., random Poisson and other systematic) sources of variability, which are magnitude-dependent. The best way to deal with this is to begin by plotting the dispersion from the mean, , for all stars versus mean magnitude difference, J, from the mean reference star level. We show this in Figure 2 for the 264 stars detected in source B. (Source A exhibits a similar overall appearance, although differing in details.) Here we see that generally increases with J, as expected for Poisson-dominated statistics. Assuming also a
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Fig. 2.—Dispersion from the mean () vs. magnitude difference ( J ) for the detected stars from source B, divided by vertical dashed lines into four magnitude ranges. The solid curve follows the expected relation for Poisson+constant noise (see text). The horizontal dashed lines refer to 1.5 in the appropriate box.
constant source of instrumental noise independent of brightness, a simple model for the rms variance is 2 ¼ 20 þ 100:4J : A curve based on this model is shown in Figure 2, with 0 ¼ 0:052 mag (the threshold value for very bright stars) and an arbitrary magnitude zero point. The curve fits well for the brightest and moderately bright stars but fails below J þ3, at which the observations exceed the model, possibly because of an additional instrumental source of noise associated with the readout. In order to limit the ultimate search for ( periodic) variability to the most likely candidates, limits were defined empirically at the 1.5 level for three intervals of magnitude: J < 0, J ¼ 0:0 2:0, and J ¼ 2:0 3:5; in each section, the mean is based on the data themselves (as in Marchenko et al. 1998). Fainter than J ¼ 3:5, the Poisson-dominated errors were larger than 0:2 mag, and these stars were neglected. These limits are shown as dashed horizontal lines in Figure 2. All 37 stars from source B above these lines were then examined for periodic variability, using the CLEAN algorithm as in Antokhin et al. (1995). A similar procedure was undertaken using source A (not shown here). There are thus three types of suspected variable stars above the horizontal dashed lines: (1) 27 stars that were found to be variable in both A and B, (2) nine stars in B only, and (3) 42 stars in A only. Clearly, types (2) and (3) are doubtful candidates; indeed, no convincing light curves were
seen with any periodicity for any of them (with peak-to-valley amplitudes 3 ). We therefore ignore these two subgroups. Among those stars in type (1), it was found that for most stars, the light curves from A and B were not significantly correlated, suggesting that these stars are also probably not intrinsically variable at a detectable level. Setting a probability P < 0:001 for n ¼ 43 values requires a correlation coefficient r > 0:48 to be acceptable as a likely variable. Only six stars of type (1) satisfied this criterion. Table 1 lists these six stars, which includes one of the central WR stars (A1), along with the other two (constant at the current level of precision) WR stars (labeled B and C) for comparison. Five of these six stars show more or less convincing periodic light curves. Figure 3 shows light curves for all eight stars. Of particular note is that the photometric period found for star A1 (1.8867 days from source A and 1.8864 days from B, with a mean of 1.8865 adopted in Table 1) is exactly half (within the errors) of the previously published period [(3:7720 0:0003)=2 ¼ 1:8860 days] found from radial velocity (RV) variations of the global emission-line spectrum of the unresolved central core of NGC 3603, obtained by Moffat & Niemela (1984) and confirmed from additional RV data by Moffat et al. (1985). This would suggest that star A1 is in fact an eclipsing binary system, since its light curve shows two dips per orbit. Assuming this to be the case, we can improve the precision of the period by requiring one of the photometric minima to exactly coincide with the systemic RV passage from negative to positive of the emission-line orbit. In this way we find an improved period of 3.7724 days, with time of primary minimum HJD 2,450,513.520. In the course of obtaining the photometry, source A found one moderately bright star (HST 12 according to Moffat et al. 1994; labeled as such in Fig. 1) to show nearly constant magnitude except for two drops by 0.7 mag at two epochs (t1 ¼ HJD 2,450,553.615 and t2 ¼ HJD 2,451,140.766). We show the light curve in Figure 4. A period search did not yield a unique period better than (t2 t1 )=n, where n is a positive integer. One of the eclipse widths based on focus-sweep monitoring and NIC3 data is 90 minutes, with one ingressing eclipse showing a drop of 0.5 mag in 25 minutes. This star, along with its close neighbor HST 13, composes the unresolved object MTT 30 of Melnick et al. (1989), which they indicated as a nonmember of NGC 3603. However, if HST 12 and HST 13 have different color indices, it is possible that HST 12 could be a member after all. We leave this up to future studies by interested researchers. After the installation of the NCS cooling system in Servicing Mission 3B, the focus monitoring program for NGC 3603 changed filters from the F110W to the F165M (i.e., essentially the H band). Photometry of the F165M data from 2002 May to 2003 September (13 epochs) revealed no additional eclipses for HST 12. The newer H-band data for star A1
TABLE 1 Summary of the Variability Properties in J-Band NICMOS Photometry Parameter
A1
B
C
HST 43
474
481
486
574
r ..................................... P (days)......................... (J ) ............................... J (mag) ......................
0.70 1.8865 0.12 0.30
0.11 ... 0.05 ...
0.32 ... 0.05 ...
0.55 0.965 0.15 0.35
0.76 ... 0.13 ...
0.76 37.3 0.32 0.60
0.49 0.548 0.33 0.8
0.86 333 0.13 0.40
Notes.—Data taken from source B. The value r is the correlation coefficient between the data from sources A and B. The value J is the total peak-to-valley amplitude of the intrinsic variability, if periodic.
Fig. 3a
Fig. 3b
Fig. 3c
Fig. 3d
Fig.Fig. 3aFig. 3.—Differential 3bFig. 3cFig.J-band 3d light curves for (a) star 474 vs. time (no significant periodicity); (b) stars HST 43, 481, 486, and 574 vs. their respective phases (based on the periods in Table 1; origin arbitrarily taken as the date of the first image, i.e., HJD 2,450,512.200); (c) the WR stars A1, B, and C as a function of A1’s phase, calculated on the basis of its RV orbit (see text); and (d ) the WR stars A1, B, and C as in (c), except that photometry from source A was used here, as opposed to source B in all the other panels.
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Fig. 4.—Plot of the J-band light curve from source A of the eclipsing variable HST 12.
basically follow the J-band points but just add noise. Crowding could be a bigger problem in the H-band data because the PSF is significantly larger. 5. DISCUSSION Among the six suspected variable stars detected from larger than average standard deviations by NICMOS in NGC 3603, we merely present the data for five of them (other than the WR star A1) without comment, except for the fact that only star HST 43 is among the brighter O star population. We lack independent information on the other four relatively faint nonOB/ WR stars. 5.1. The Brig ght Central Star A1 However, the brightest variable, A1, is along with star B the (visually) brightest member star in the cluster. Together, the three central WR stars A1, B, and C outshine by far each of the remaining members and are located within 100 (0.03 pc) of each other in the dense core, possibly a result of mergers (by analogy with the predicted behavior of NGC 3603’s clone cluster R136; Portegies Zwart et al. 1999). Alternatively, they could be simply the most luminous and thus massive stars found in this very dense cluster core. This becomes even more plausible with the trend of decreased Teff of the hottest O stars (Herrero 2003) and increased Teff of WR stars (P. Crowther 2004, private communication), thus making a more continuous sequence in the CMD. All three WR stars in NGC 3603 have spectral types WN6ha, implying that they may be extremely luminous, more massive versions of their fainter, nearby O3type cousins in the same cluster. As such, they still contain considerable hydrogen (XH ¼ 0:60; Crowther & Dessart 1998), unlike most WR stars. At the ¼ 0:05 mag level, WR stars B and C remain constant. A1’s clear variability with J-band amplitude of 0.3 mag leads to a periodicity that matches the previous RV periodicity remarkably well. As noted in x 4, A1 is likely an eclipsing variable with two dips per cycle. This is an extremely fortunate and valuable result, since it opens the door to obtaining a reliable Keplerian-based estimate of the masses of the potentially most massive main-sequence star in NGC 3603 and possibly in the whole Galaxy. On the basis of isochrone and spectral line fitting, values of 100–150 M are expected. However, no Keplerian-based mass determination is known for any star above 60 M. Among the most massive main-sequence stars in binaries to date are the WN6ha component of WR22 ¼ HD 92740 in the Carina Nebula, which has 55 7 M
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Fig. 5.—Plot of the J-band light curve of the brightest cluster member, the WN6ha star A1, phased to a period of 3.7724 days with the WR star passing inferior conjunction at phase zero (HJD 2,450,513.520). The data were obtained from source A from an earlier attempt than the data for A1 shown in Fig. 3.
(Schweickhardt et al. 1999), dropping from the original estimate of 72 M by Rauw et al. (1996); the O3 V star in HD 93205 in the same region, which has 56 4 M (Morrell et al. 2001); and the O3 V star R136-38 in the periphery of R136 in the LMC, which has 57 1 M, although two other O3 V stars in binaries in this region have significantly lower masses (Massey et al. 2002). While preparing this paper, we became aware of an even more massive system, WR20a, with two WN6ha or O3 If */ WN6ha stars with minimum masses of 71 and 69 M (Rauw et al. 2004) and actual masses of 83 5 and 82 5 M from its eclipse light curve (Bonanos et al. 2004). Here we can only present a very rough, preliminary estimate of A1’s observed mass. We approach the problem in two steps. Step 1: Starting with the light curve in Figure 5, we have fitted a standard Roche model with circular orbit to the light curve, without an absorbing envelope. In such models, the mass ratio q M2 =M1 is poorly determined. Thus, q was not fitted rigorously; rather, other unknown parameters were 2 fitted for a reasonable range of values for q ¼ 0:2 7:0. The following input parameters were fixed: the temperature of the assumed companion O3 star, T1 ¼ 43 kK (a good compromise between supergiant 40 kK and main-sequence 46 kK; Repolust et al. 2004; the exact value is not critical for the results of this paper); the albedos A1; 2 ¼ 1; the gravity-darkening exponent is taken to be unity for both stars; the limb darkening coefficient u ¼ 0:08 for star 1 (van Hamme 1993) and u ¼ 0:1 for star 2 (the WR star); and the effective monochromatic wavelength, 1.6 m. The WR component was taken to be at inferior conjunction at phase zero (this yields better fits than if this is taken to be at phase 0.5). The fitted parameters as a function of q are listed in Table 2: they are the orbital inclination, i, the Roche-lobe–filling coefficients 1 and 2, and the effective temperature of the WR star, T2. From these, the values of the radii of spheres that have the same surface area as the corresponding Roche lobe surfaces (r1 and r2) are given in terms of the mean orbital separation (a). The best fit for a given q is defined by the minimum of 2 in the multiple parameter space, with taken to be 0.05 mag. The main results of the ‘‘step 1’’ fits are that (1) the O star fills its Roche lobe, (2) the WR star is near to filling its Roche lobe (in view of the model used without extended atmosphere, this represents a very rough representation of the WR star with its wind), (3) the values of q that lead to the best fits lie in the interval 0.5–2.0 (for q ¼ 0:2, 5, and 7, the light curves show poor visual agreement with the observations, although 2 remains relatively small), and (4) values of i and T2 show little dependency on q. Figure 5 shows the best fit for the q ¼ 1
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TABLE 2 Model Parameters of the System A1
q M2 /M1
i (deg)
1
2
r1/a (mean)
r2 /a (mean)
T2 (K)
2/Ndof
0.2.................................. 0.5.................................. 1.0.................................. 2.0.................................. 5.0.................................. 7.0..................................
80.5 72.3 70.9 71.5 76.7 82.3
0.999 0.995 0.997 0.996 0.999 0.990
0.879 0.889 0.878 0.891 0.872 0.807
0.52 0.44 0.38 0.32 0.25 0.22
0.21 0.28 0.32 0.38 0.44 0.43
48800 45600 45600 45400 45100 46500
48/36 45/36 45/36 46/36 49/36 54/36
Notes.—Parameters with subscript ‘‘1’’ refer to the assumed O3 V star, and those with ‘‘2’’ refer to the WR star. All data assume the WR star passing in front at phase zero (HJD 2,450,513.520). The value Ndof is the number of degrees of freedom.
model (the q ¼ 0:5 2:0 models look similar). We adopt i ¼ 71 and find an effective temperature for the WR star T2 close to 46 kK , i.e., slightly higher than that assumed for the O star. Crowther & Dessart (1998) find T2=3 ¼ 36 kK for star A1, assuming it to be single. This low value may be at least partly a consequence of the significantly lower value for the interstellar extinction they adopted for the cluster core (EBV ¼ 1:23) compared with the optically well-observed value EBV ¼ 1:44 (Moffat 1983; Melnick et al. 1989). Such a low value brings the observed B V colors back to intrinsic mid-B stars, rather than the early O stars that one finds in the core of NGC 3603. The only radial velocity information we have is from the unresolved study of the cluster core by Moffat & Niemela (1984), who discovered the 3.7720 days period. For the net He ii k4686 and N iv k4058 emission lines arising in the core, they find an RV amplitude of Kobs ¼ 72 5 km s1, confirmed later by Moffat et al. (1985). Assuming the emission-line RV of star A1 to be diluted by the presence of WR stars B and C in the spectrograph slit according to their respective brightnesses, we can estimate the approximate K-value for the WR component of A1, assuming that WR stars B and C also contribute to the total emission-line flux ( but not to the 3.7720 days RV amplitude) according to their known equivalent widths (Drissen et al. 1995): KA1 WA1 fA1 ¼ Kobs (WA1 fA1 þ WB fB þ WC fC ); where fX and WX are the relative continuum flux (in the B band) and equivalent widths of the prominent He ii k4686 emission line for star X ( fX ¼ 1:00, 1.01 and 0.63 and WX ¼ 27, 46, and 16 8 for stars A1, B, and C, respectively), so that the product WX fX is the relative line flux. We also assume that each WR star has the same systemic RV, although this is not critical here. This leads to KA1 ¼ 223 km s1 and for a circular orbit to a mass function of f2 (M ) (M1 sin i)3 =(M1 þ M2 )2 ¼ 1:04 ; 107 ½KA1 (km s1 ) 3 P(days) ¼ 4:35 M , where index 1 is for the O star and index 2 the WR star. Unfortunately, there is no reliable way to assign a formal uncertainty to KA1 and thus to f2(M ). With i ¼ 71 , we obtain triplets of values (q, M1/M, M2/M) for q in the best range, 0.5–2: (0.5, 12, 6), (1.0, 21, 21), and (2.0, 46, 93). Given that the WR stars in NGC 3603 are optically brighter (and therefore also more massive) than any of the O stars implies that in A1, in all likelihood q > 1; hence, the WR star in A1 likely has a mass in the range 30–90 M and the O companion a mass 25–50 M. Although these values appear reasonable, they are still not very well constrained. This calculation also assumes that all three WR stars are centered similarly on the slit, which can only be
an approximation to the truth. At least it is reasonable to assume that the adjacent brightest WR stars A and B were centered on the slit, so that the approximation is satisfactory. Step 2: Calculations here are more exploratory than restrictive on the parameters of the system. We use the modified SEI code (Sobolev solution with exact integration; see Lamers et al. 1987 for the original approach and St-Louis et al. 2004 for a detailed description of modifications) to calculate the light curve of the system. In its modified form, the SEI code allows one to fit profiles of the emission lines originating in the WR wind while accounting for the presence of a cavity in the WR wind filled in by the O star wind, as well as an additional emission component coming from the wind-wind collision zone. In addition, a light curve comes out as a by-product of the calculations. The large number of fitted parameters demands an adequate data set, which should be composed primarily of phase-resolved spectroscopy plus high-precision photometry. This combination would allow one to place firm restrictions on the fundamental parameters: orbital inclination, sizes and luminosity ratio of the components, and mass-loss rate of the WR star. Since our spectral data are limited to only one spectrum (Drissen et al. 1995) and precision of the photometry is not high, we rely on input from the step 1 calculations. We take i ¼ 71 and use the range of M (WR) þ M(O) masses (hence, the orbital separation) prescribed by step 1. We also take into account that the luminosity of the O component should be comparable to the luminosity of the WR star. We use a lower limit of L(O)=L(WR) ¼ 0:5 in the step 2 simulations; lower values provide inconsistent fits to the photometric data. A quick exploration of the L(O)=L(WR) > 1 domain shows that L(O)=L(WR) 2 can be used as a reasonable formal upper bound, even if we deduced in step 1 that the WR component is likely more luminous than the O companion. Hence, we use two families of solutions, L(O)=L(WR) ¼ 0:5 and 1.5, to highlight the qualitative differences between two limiting cases. As a starting point of simulations, we use the profiles of the He ii kk4340, 4686 emission lines to yield the corresponding set of SEI parameters (Lamers et al. 1987) restricting the density/emissivity distribution in the WR wind. Then ˙ (WR), R(WR), and R(O) to vary. Depending on we allow M the L(O)=L( WR) ratio, this leads to two different solutions. Apparently (see Fig. 6), the quality of the photometric data does not allow one to discriminate between the solutions. Hence, we list two possible scenarios: (1) a ‘‘normal’’ WR star, i.e., fixed L(O)=L(WR) ¼ 1:5 and a ¼ 27 R corresponding to q ¼ 0:5 in step 1, gives R(WR) ¼ 6:5 R , R(O) ¼ 11:0 R ˙ (WR) ¼ 105 M yr1 (50% (with 20% uncertainty), and M
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Fig. 6.—Same data as in Fig. 5, now fitted with the modified SEI model. The solid line shows the ‘‘normal’’ WR solution (a low-luminosity WR component; see text for more details); the dashed line shows the ‘‘high-luminosity WR’’ solution.
uncertainty), comparable with the values derived for shortperiod eclipsing WR+O binaries (Moffat & Marchenko 1996) and in line with the step 1 calculations; (2) predictably, an overluminous [L(O)=L(WR) ¼ 0:5 and a ¼ 53 R, corresponding to q ¼ 2 in step 1] WR star leads to a rather extreme (but nevertheless more likely, given the high luminosity of A1) configuration with R(WR) ¼ 15 R, R(O) ¼ 18 R, and ˙ (WR) ¼ 5 ; 105 M yr1 , but again within the step 1 range M of solutions. Even this latter solution gives a factor of 2 smaller radius for the WR star than does the R( WR) ¼ 30 R found for A1 by Crowther & Dessart (1998). This is in line with Crowther & Dessart’s probably too low estimate for Teff of A1, as mentioned above. 5.2. Comparison with Other Luminous WR Stars It is interesting to note that one of the four luminous WR stars in R136 must also be a short-period binary, given the 4.377 day cyclic behavior of the unresolved emission line of He ii k4686 discovered in the same way as for NGC 3603 by Moffat & Seggewiss (1983) and confirmed by Moffat et al. (1985). Its RV amplitude is Kobs ¼ 37 km s1; this is significantly less than for NGC 3603 and possibly the result of more diluted emission lines than NGC 3603’s WR spectrum and/or the result of a lower orbital inclination angle (the masses are also somewhat lower in R136; Crowther & Dessart 1998). Given that NGC 3603 and R136 are essentially clones out to a physical radius of r ¼ 1 pc (Moffat et al. 1994), their similar diluted WR emission-line RV behavior is quite remarkable.
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Finally, although it does not show large IR variability, star C in NGC 3603 may be the long-period binary (P of hundreds of days?) suspected from the spatially unresolved RV study of Moffat & Niemela (1984). It shows the strongest X-ray flux of any single star in all of NGC 3603 (and is similar to the extremely X-ray–bright WN6ha star WR25 ¼ HD 93162 in the Carina Nebula; Seward & Chlebowski 1982), even though it is about a factor of 2 less luminous in the optical / UV than stars A1 or B (Moffat et al. 2002). Compared with star B with LX =Lbol ¼ 107 , star C has LX =Lbol > 20 ; 107 . Star A1 is confused on Chandra images, mainly with O3 stars A2 and A3, but also shows moderately high combined (diluted?) LX =Lbol 4 ; 107 . This fits well with the fact that shortperiod WR+O binaries tend to show only slightly enhanced X-ray fluxes due to colliding winds, whereas in longer period systems (P on the order of years), the lower intrinsic X-ray fluxes from the colliding winds are more than compensated by the lower absorption of the more diluted winds seen along the line of sight to where the winds collide. On the other hand, it is star A1 (or B or one of A1’s close, bright O3 neighbors A2 or A3) that exhibits significant spatially resolved, nonthermal centimeter radio flux. Time-dependent long-term variability of both X-ray and radio flux of star C might explain this apparent contradiction of star C having very high X-ray luminosity yet low radio flux. It is also interesting that star C has the weakest relative WR-like emission lines in its spectrum, followed by A1 and B with the strongest. This goes in the same rank order as the mass-loss rates for these three stars derived by Crowther & Dessart (1998), assuming each star to be single. Binary dilution could, however, change this order somewhat. 6. CONCLUSION Despite the relatively modest precision of the NICMOS photometry [min (J ) ¼ 0:05 mag], this study has revealed several variable stars in NGC 3603, of which the most important by far is the potentially very massive eclipsing binary A1. The ( pseudo-)WR component of A1 might turn out to be the most massive star ever ‘‘weighed’’ in the Galaxy. A highprecision (but challenging because of the severe crowding) RV study for both components of A1 is urgently needed to check this. An additional study of the phase dependency of colliding wind effects in A1 could also be rewarding, in part to obtain a second independent estimate of the orbital inclination and in part to constrain the winds of the two stars. Unfortunately, the Hubble Space Telescope project that was approved to carry this out using STIS in cycle 13 had to be canceled because of a fatal instrument failure. A. F. J. M. thanks NSERC (Canada) and FQRNT (Que´bec) for financial support.
REFERENCES Antokhin, I., Bertrand, J.-F., Lamontagne, R., Moffat, A. F. J., & Matthews, J. Lamers, H. J. G. L. M., Cerruti-Sola, M., & Perinotto, M. 1987, ApJ, 314, 726 1995, AJ, 109, 817 Marchenko, S. V., et al. 1998, A&A, 331, 1022 Bonanos, A. Z., et al. 2004, ApJ, 611, L33 Massey, P., Penny, L. R., & Vukovich, J. 2002, ApJ, 565, 982 Brandl, B., Brandner, W., Eisenhauer, F., Moffat, A. F. J., Palla, F., & Melnick, J., Tapia, M., & Terlevich, R. 1989, A&A, 213, 89 Zinnecker, H. 1999, A&A, 352, L69 Moffat, A. F. J. 1983, A&A, 124, 273 Crowther, P. A., & Dessart, L. 1998, MNRAS, 296, 622 Moffat, A. F. J., et al. 2002, ApJ, 573, 191 Drissen, L., Moffat, A. F. J., Walborn, N. R., & Shara, M. M. 1995, AJ, 110, 2235 Moffat, A. F. J., Drissen, L., & Shara, M. M. 1994, ApJ, 436, 183 Eisenhauer, F., Quirrenbach, A., Zinnecker, H., & Genzel, R. 1998, ApJ, 498, Moffat, A. F. J., & Marchenko, S. V. 1996, A&A, 305, L29 278 Moffat, A. F. J., & Niemela, V. S. 1984, ApJ, 284, 631 Herrero, A. 2003, in IAU Symp. 212, A Massive Star Odyssey: From Main Moffat, A. F. J., & Seggewiss, W. 1983, A&A, 125, 83 Sequence to Supernova, ed. K. van der Hucht, A. Herrero, & C. Esteban (San Moffat, A. F. J., Seggewiss, W., & Shara, M. M. 1985, ApJ, 295, 109 Francisco: ASP), 3 Morrell, N. I., et al. 2001, MNRAS, 326, 85
No. 6, 2004
VARIABILITY IN MASSIVE STARS IN NGC 3603
Portegies Zwart, S. F., Makino, J., McMillan, S. L. W., & Hut, P. 1999, A&A, 348, 117 Rauw, G., et al. 2004, A&A, 420, L9 Rauw, G., Vreux, J.-M., Gosset, E., Hutsemekers, D., Magain, P., & Rochowicz, K. 1996, A&A, 306, 771 Repolust, T., Puls, J., & Herrero, A. 2004, A&A, 415, 349 Schweickhardt, J., Schmutz, W., Stahl, O., Szeiferth, T., & Wolf, B. 1999, A&A, 347, 127
2861
Seward, F. D., & Chlebowski, T. 1982, ApJ, 256, 530 Smith, L. F., Shara, M. M., & Moffat, A. F. J. 1996, MNRAS, 281, 163 St-Louis, N., Moffat, A. F. J., Marchenko, S. V., & Pittard, J. M. 2004, ApJ, submitted van Hamme, W. 1993, AJ, 106, 2096
