ABSTRACT. GX 339-4 has long been known as a black hole candidate because of its rapid variability and high/low. X-ray states, which are similar to those of ...
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The Astronomical Journal, 123:1741–1749, 2002 March # 2002. The American Astronomical Society. All rights reserved. Printed in U.S.A.
OPTICAL OBSERVATIONS OF THE BLACK HOLE CANDIDATE GX 339 4 (V821 ARAE) A. P. Cowley1 and P. C. Schmidtke1 Department of Physics and Astronomy, Arizona State University, Tempe, AZ, 85287-1504; anne.cowley@asu.edu, paul.schmidtke@asu.edu
and J. B. Hutchings1 and David Crampton1 Dominion Astrophysical Observatory, Herzberg Institute of Astrophysics, National Research Council of Canada, 5071 West Saanich Road, Victoria, BC V9E 2E7, Canada; john.hutchings@nrc.ca, david.crampton@nrc.ca Received 2001 November 13; accepted 2001 December 4
ABSTRACT GX 339-4 has long been known as a black hole candidate because of its rapid variability and high/low X-ray states, which are similar to those of Cyg X-1. Although GX 339-4 is assumed to be a binary, the orbital period has not yet been convincingly determined, and hence little is known about the nature of the component stars. In this study simultaneous photometric and spectroscopic observations have been made in an effort to address these problems. Although a definitive period has not been found, we present evidence that it lies near 0.7 days. The I-band light curve has an amplitude of only 0.2 mag, with flickering of >0.1 mag superposed, making it difficult to determine the orbital period from light variations alone. The emission-line velocity amplitudes are small, suggesting that GX 339-4 is seen at a low orbital inclination angle. Even though the period is still not determined, the relative phasing of the velocities and light minima are known, since the photometric and spectroscopic observing runs overlapped. Assuming the orbital period is near 0.7 days, we can set limits on the mass ratio and stellar masses. For reasonable assumptions about the inclination angle, the mass ratio appears to lie between q 3 and 15, where q = MX/Mdonor. If the donor star mass is greater than 0.3 M , the X-ray source is likely to be a black hole. Key words: binaries: close — stars: individual (GX 339 4) — X-rays On-line material: machine-readable table
B > 21 (Hutchings, Cowley, & Crampton 1981; Ilovaisky & Chevalier 1981). From earlier simultaneous X-ray and optical observations, it has been suggested that when the source is in its soft-high state the optical brightness is V 16.5–17, while in the low-hard state the source is brighter at V 15.4 (Motch et al. 1982; Makishima et al. 1986; Corbet et al. 1987). However, the data taken in 1999 March for this paper reveal the mean V magnitude was V 16.9, while simultaneous RXTE All Sky Monitor (ASM) data show that GX 339-4 was in a low X-ray state. Figure 1 displays the ASM daily averages from 1998 to 2000, with dates of our optical observations marked. Thus, as Corbet et al. state, ‘‘ the correlation between optical brightness and X-ray state is. . .poorly defined.’’ In spite of the considerable amount of observational data obtained to date for GX 339-4, we still have no clear evidence of the properties of the component stars or even the orbital period. Measurements of low-resolution spectra (Cowley, Crampton, & Hutchings 1987) suggested the velocity varies by about 250 km s 1, but these authors were only able to limit the orbital period to lie in the range 0.2 to 2.0 days. Corbet et al. (1987) observed the system photometrically and found no periodic modulation for semiamplitudes greater than 0.03 mag for periods less than 0.5 days and greater than 0.07 mag for longer periods. They concluded the period must be greater than 0.5 days. Part of the difficulty in determining the orbital period has been that the source shows so much irregular variability at all wavelengths that any periodic phenomena are easily obscured. Callanan et al. (1992) announced an orbital period of 14.8 hr (0.62 days) based on photometric data obtained during a low X-ray state when irregular variability was lower than in
1. INTRODUCTION
GX 339-4 (V821 Ara; X1659 487) has long been described as a probable black hole binary (or ‘‘ black hole candidate ’’). It displays a wide variety of unusual behaviors. Even when GX 339-4 was first discovered, its aperiodic X-ray variability on timescales from milliseconds to seconds placed it in a category similar to the black hole binary Cyg X-1 (Samimi et al. 1979). An optical identification was quickly made (Doxsey et al. 1979) with a faint (17–18 mag) star showing He ii 4686 emission, characteristic of lowmass X-ray binaries and other X-ray binaries not containing a massive early-type star. Rapid variability was also found at optical wavelengths and shown to be anticorrelated with the X-ray variability (Motch, Ilovaisky, & Chevalier 1982; Motch et al. 1983). Although the rapid variability present in the low-hard state is similar to the behavior seen in the black hole systems Cyg X-1 and Cir X-1, such variability is also found in the X-ray source V0332+53 (BQ Cam), which is a Be/neutron star system (Stella et al. 1985; Corbet, Charles, & van der Klis 1986). Radio data show evidence of a jet associated with GX 339-4 (Fender et al. 1997, 1999) whose behavior is strongly coupled with the X-ray state (Corbel et al. 2000). At X-ray wavelengths GX 339-4 displays three distinct states: X-ray–off, hard-low, and high-soft. During the X-ray off state the optical counterpart is extremely faint with
1 Visiting Astronomers, Cerro Tololo Inter-American Observatory, National Optical Astronomy Observatory, which is operated by the Association of Universities for Research in Astronomy, Inc., under cooperative agreement with the National Science Foundation.
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Vol. 123 TABLE 1 Photometry of GX 339-4
HJD 2,450,000+
V
V
R
R
I
I
1258.804................ 1258.812................ 1258.819................
16.845 ... ...
0.009 ... ...
... 16.133 ...
... 0.004 ...
... ... 15.384
... ... 0.005
Note.—Table 1 is presented in its entirety in the electronic edition of the Astronomical Journal. A portion is shown here for guidance regarding its form and content.
Fig. 1.—RXTE-ASM light curve for GX 339-4 showing 1 day averages of the X-ray intensity vs. Julian date. The optical data in this paper were taken during the time interval marked by the vertical dashed lines. The source was in a low X-ray state, but had not yet completely turned off.
the high state observed by Corbet et al. However, there has been no confirmation of the period proposed by Callanan et al. We undertook a photometric-spectroscopic study to independently find the orbital period by studying phaserelated changes in the spectra and luminosity obtained at the same time. As part of this process we also tested the Callanan et al. period using our new data. Our ultimate goal was to determine the masses of the component stars and study their mutual interactions. All of our observations were obtained at Cerro Tololo Inter-American Observatory (CTIO) in 1999 March. The 0.9 m and 4 m telescopes were used for the photometry and spectroscopy, respectively. The two observing runs overlapped, allowing us to examine the relative phasing between these data sets.
2. OBSERVATIONAL DATA: OPTICAL PHOTOMETRY
Three-color photometry of GX 339-4 was obtained with the 0.9 m telescope and Tek 2K No. 3 CCD during eight nights, 1999 March 20–27 UT. In total, 23 VRI sequences were obtained over the first four nights, while an additional 69 images in I alone were taken on the remaining nights. Nearly all images were 5 minute exposures. The seeing was 1>1–1>2 under photometric skies at the beginning of the run, but conditions deteriorated slightly near the end. All of the data were processed with DAOPHOT (Stetson 1987) and calibrated with observations of Landolt’s (1992) standard stars. As first noted by Remillard & McClintock (1987), a faint field star lies only 1>1 northwest of the optical counterpart of GX 339-4. From our best images, we measured the magnitudes of this star to be V = 20.70, R = 19.24, and I = 18.42, similar to values given by Remillard & McClintock. The expected contribution of the star was subtracted from all frames prior to the final imageprofile fitting. Shahbaz, Fender, & Charles (2001; see their Fig. 1) show the field star itself is a blend of at least two components, but we did not resolve the pair and therefore treated the contaminant as a single star. Differential photometry of GX 339-4 was calculated using reference stars within the CCD field of view. The mean errors of an individual measurement were 0.006, 0.008, and
0.007 mag in V, R, and I, respectively, with little variation from night to night. Table 1 contains the photometric data, a sample of which is shown in the printed version of this paper. An extended version of Table 1 is available in the electronic edition and contains all our photometric observations of GX 339-4. Figure 2 shows the I-band data plotted for all eight nights. Within each night the system varied by up to 0.2 mag, and the overall brightness ranged from I 15.2 to I 15.6 during the observing run. The mean visual magnitude of our observations was V 16.9, which is 1.5 mag fainter than previously reported photometric measurements of GX 339-4 when in its X-ray low-hard state. This dimming may be related to the source’s imminent transition to an X-ray–off state, which occurred within 25 days after the conclusion of our photometric observations (see Fig. 1). However, we find no evidence for a systematic optical dimming over the course of the observing run. Clearly, a much longer photometric campaign is needed to firmly establish a relationship between optical brightness and X-ray state. GX 339-4 is quite red for an X-ray binary, with B V 0.85 (Ilovaisky et al. 1986). A more typical value would be +0.01 (van Paradijs 1983). Strong interstellar absorption lines show the color is strongly affected by interstellar reddening. Figure 3 is a plot of V versus V I, showing that the color may become slightly redder as it fades, but we have too few data points for this to be very significant. Our data are consistent with the off-state photometry of Remillard &
Fig. 2.—I-band photometry of GX 339-4 (V821 Ara) over eight nights of observation, 1999 March 20–27, with error bars shown. Spectroscopic observations were taken during the final three nights, plus one additional night following the photometry run.
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Fig. 4.—Normalized mean spectrum of GX 339-4 with principal features identified. The absorption features (marked ‘‘ is ’’) are all interstellar lines ˚ . If its reality is confirmed or bands. Note the weak emission line at 6508 A with additional spectra, it might be a violet-displaced component of H (i.e., a ‘‘ jet ’’ feature) as is seen in some of the supersoft X-ray binaries (e.g., Cowley et al. 1998).
star. Shahbaz et al. (2001) observed GX 339-4 with low-resolution spectroscopy when it was in its very faint, X-ray–off state. The secondary star was not detectable even then, so the donor star must be extremely faint.
Fig. 3.—V magnitude vs. V I color. There may be a weak trend for the source to become slightly redder as it fades, but we have too few data points for this to be very significant.
McClintock (1987), who found V I 1.7 when the source had dimmed to V 20.4.
3. OBSERVATIONAL DATA: SPECTROSCOPY
Spectroscopic observations were also obtained at CTIO, using the 4 m telescope over four consecutive nights from 1999 March 26–29, such that the first three nights were coincident with photometric observations. The spectra cover the ˚ with a resolution of wavelength region from 3800–6800 A ˚ pixel 1. With the 1>5 slit we used, this corresponds to a 1A ˚. spectral resolution of 3 A Eight spectra of GX 339-4 were obtained on the first, third, and fourth nights of the observing run; six spectra were taken on the second night. The spectrum contains broad emission lines of H, He ii (4686 is especially strong), ˚ (see Fig. 4). It is He i, and the C iii–N iii blend at 4640 A likely that the emission lines arise in an accretion disk around the compact star, perhaps at a range of disk radii. Some weaker emission features also appear in our summed spectrum. Wu et al. (2001) identified N ii 6505, but we question whether this feature is due to N ii. There is no evidence of other lines from this multiplet, and other stronger N ii lines are not present. Interstellar absorption features are prominent and include the broad diffuse bands at 4430, 6180 and the narrower features at 5780, 5797, 6203, 6270, 6284, 6614 and the Na i lines at 5890, 5896. There is no evidence of stellar absorption lines which could be attributed to a companion
3.1. Radial Velocities, Equivalent Widths, and Line Profile Variations In Figures 5a–5d we plot the individual spectra in a time series for each of the four nights. In order that some detail in the line profiles can be seen, we have only shown two spec˚ , which includes He ii and H, and tral regions: 4600–4900 A ˚ 6500-6600 A to isolate H. One can easily see the line profiles and strengths change during a single night. Velocities and equivalent widths measured from each spectrum are given in Table 2. In the case of He ii 4686 the profile is usually double peaked, with the short-wavelength peak stronger. It is not clear whether these peaks represent two narrow, blended components arising in different parts of the system, a broad emission with a superposed narrower peak, or a broad line with a central absorption. Therefore, we have measured the line in several ways. We fitted the line by two Gaussian profiles and found their central wavelengths. We also centered on the mid-intensity region of the overall broad line. For He ii 4686 the velocity of the center of the broad emission line is given in Table 2. We have not tabulated the peak velocities since it is difficult to know if the short-wavelength peak crosses over the other one (i.e., has a high velocity amplitude), as in s-wave disturbances common in cataclysmic variables, or if the red and violet peaks are independent. However, we do tabulate the separation of the two peaks. In some of the spectra there is a third, weaker peak, which also contributes to the line profile. Since the measurements of individual features are quite scattered, we have also determined velocities by cross-correlation using both a restricted wavelength region (4400–5100 ˚ ) and using the whole spectrum. The danger in this techniA que is that, if different emission lines have different amplitudes and phasing, the resulting velocities may be meaningless. Each spectrum was correlated against the mean of all the spectra. In Table 2 we give the cross-correla-
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Fig. 5.—In each of the four panels, two selected spectral regions for data obtained during a single night. Time increases from bottom to top, with UT given below each spectrum. The signal of each spectrum has been normalized to a constant mean value over the whole wavelength range, and the spectra are plotted ˚ region has three-point smoothing; the region near H is unsmoothed. Note that the line profiles change within with a constant vertical offset. The 4600–4900 A a single night.
˚ region, which is domition velocities for the 4400–5100 A nated by He ii 4686. H emission is much weaker and contributes little to these calculations. Note that the crosscorrelation velocities have an arbitrary zero point that is about 35 km s 1 below the H velocities. This is because the spectra that were co-added to produce the mean spectrum have different intensities so that some are weighted more heavily. The mean velocity, derived from the strongest lines, are +2 km s 1 for H, +4 km s 1 for H, and +26 km s 1 for He ii 4686. Since the source lies toward the galactic center region (l = 339 , b = 4 ), a low systemic velocity is expected. For the strongest lines we examined their continuum to assess the full velocity range of gas and measured the
FWHM both from nightly mean spectra and from the summed spectrum. We find FWHM(H) = 740 km s 1 and FWHM(4686) = 800 km s 1, with little change from night to night. In addition, equivalent widths of the major emission lines were measured and are included in Table 2. These vary from night to night, but appear to be only weakly correlated with each other or with the system brightness. For example, on the first spectroscopic night (March 26) both He ii 4686 and H decreased in intensity through the night while the system was both bright and rising. However, on subsequent nights there was no clear correlation with brightness, as follows. On March 27, GX 339-4 was faint, while H was strong and 4686 was at a mean brightened at its weakest level, but H weakened through the night. On March 29, the equivalent widths of both H and 4686 were
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Fig. 5.—Continued
approximately constant, but H was at its weakest level and He ii was at an intermediate level. We have no photometric data for the last spectroscopic night.
4. SEARCH FOR PERIODICITIES
Our primary goal in this study was to determine the orbital period of GX 339-4 from our data. This included testing the period proposed by Callanan et al. (1992) of P = 0.619 days (14.86 hr). Periodograms (Horne & Baliunas 1986) were calculated for the I-band photometry and for each of the measured emission-line velocities. Three of these periodograms are shown in Figure 6. The peaks with frequencies less than 1 day 1 (i.e., those implying periods greater than 1 day) appear to arise from the limited number of observing nights. These longer periods just place all the data consecutively or with minimum phase overlap. None of the sug-
gested ‘‘ periods ’’ greater than a day is confirmed by any other data set. The periodograms for the photometry and velocities all show peaks near P 0.7 days, but none of the data has a peak at the Callanan et al. period (dashed vertical line, Fig. 6), and all of the data sets indicate that the orbital period is somewhat longer. It is likely that the orbital period lies near 0.7 days, but additional data are required to refine the period. We note that in the high-state V data shown by Callanan et al. (their Fig. 5), the strongest peak (peak 3) in the Fourier transform corresponds to P 0.7 days, although it is not discussed by these authors. There are at least two reasons why a unique period is so difficult to find in these data: (1) The light amplitude is small, with the range of nightly I means being only 0.2 mag. In addition, on each night there is considerable shortterm flickering, which on some nights amounts to as much as 0.2 mag (see Fig. 2). (2) The velocity amplitude is small.
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TABLE 2 Spectroscopic Measurements of GX 339-4 H UT Date 1999 March (mid-exp) 26.227 ............ 26.242 ............ 26.301 ............ 26.315 ............ 26.369 ............ 26.385 ............ 26.401 ............ 26.416 ............ 27.231 ............ 27.245 ............ 27.301 ............ 27.315 ............ 27.367 ............ 27.381 ............ 28.190 ............ 28.204 ............ 28.260 ............ 28.274 ............ 28.355 ............ 28.369 ............ 28.403 ............ 28.414 ............ 29.186 ............ 29.201 ............ 29.255 ............ 29.269 ............ 29.344 ............ 29.359 ............ 29.395 ............ 29.410 ............
Center (km s 1) 9 6 6 1 0 5 28 1 30 17 29 40 45 45 13 28 38 39 17 14 25 16 8 3 5 12 2 12 2 10
Peak (km s 1) 45 26 35 90 32 31 6 13 10 45 83 70 42 40 72 67 122 104 86 172 172 182 28 150 145 0 49 127 49 1
He ii 4686 A˚a
H W (A˚) 9.3 8.6 8.7 9.0 8.3 7.5 8.0 7.8 9.0 8.6 9.0 8.6 9.2 9.5 9.5 9.5 8.1 7.8 8.1 7.8 7.4 7.7 7.8 7.6 8.0 7.5 7.7 7.7 7.9 8.1
Center (km s 1) 34 7 27 34 17 34 57 5 3 34 34 15 50 0 46 52 33 26 3 57 34 26 32 95 27 15 55 37 27 34
W (A˚) 2.6 2.5 2.7 2.8 2.4 2.2 2.5 2.9 2.5 3.1 2.9 2.6 2.4 2.6 2.3 2.7 2.0 2.3 2.3 2.3 2.5 3.2 1.9 2.0 2.1 2.2 2.8 3.0 2.4 2.6
Center (km s 1)
Peak Sep. (km s 1)
39 25 31 26 42 26 12 85 12: 31 44 83 83 83 20 44 6 25 31 38 25 44: 140 89 6 127 19 25 44 1
499 v 524 v 506 v 493 v 499 v 486 v 531 v 531 v 473 v, 3w 313: 3w 460 v 473 v 461 310: 3w 474 v 524 v . . . v, 3w . . . v, 3w 410 v 403 365 397: 339: 396: r? 339: v 397: v 361: ... 384: v 390 v
W (A˚) 5.4 5.6 5.8 4.9 4.9 4.4 4.3 4.0 4.3 4.2 4.3 4.3 3.8 3.9 3.6 3.6 4.1 4.4 3.1 3.4 3.6 4.1 4.5 4.5 4.1 3.8 4.0 4.5 4.6 4.3
Cross-corr.b Rel. Vel. (km s 1) 103 72 53 108 78 80 42 80 4 8 40 6 19 12 95 51 46 73 40 56 9 33 7 69 42 20 9 4 15 48
a A semicolon indicates peaks not clearly separated; ‘‘ v ’’ or ‘‘ r ’’ indicates whether the violet or the red peak is stronger; ‘‘ 3w ’’ indicates the presence of a weaker third peak on the red side of the line. b Cross-correlation velocity of the 4400–5100 A ˚ region relative to mean of all spectra; the zero point is arbitrary.
The broad emission lines of H and He ii show a semiamplitude of K 20–30 km s 1 when fitted to their ‘‘ best ’’ period near 0.7 days. The velocities also show considerable scatter, partially due to the low signal-to-noise ratio of the spectra, but probably also due to additional components contributing to the emission lines at various phases. Both the low velocity amplitude and small range in brightness suggest the orbital inclination is low. Since the system does not eclipse, the inclination must be less than i 60 , but the velocity data suggest it could be much lower than that upper limit. Since our photometric and spectroscopic data were obtained on overlapping nights, phasing the data on an approximate period allows one to see how the light and velocity curves are correlated. In Figure 7a we show the Iband light curve and the velocity curves from the cross-cor˚ region and H center phased relation of the 4400–5100 A on a period of 0.7 days with the same value of T0. Data from each night are shown by different symbols. As previously pointed out, the cross-correlation velocities have an arbitrary zero point, but this is not of concern since we are mainly interested in the amplitude and phasing of the velocity curves. The light curve, based on a period of 0.7 days, is not entirely symmetrical in shape; the rise time is longer than the decline. It is not known whether the light curve results
from partial occultation of the bright accretion disk by the donor star, from disk height variations causing occultations of the inner disk, from variations in the visibility of the X-ray heated face on the donor star, or from some combination of these. Adopting minimum light as phase = 0.0, the maximum velocity for the strong emission lines occurs at 0.8, and the minimum at 0.3, with both being poorly determined. This is approximately the phasing expected if minimum light occurs at conjunction when the compact star is behind and if the emission lines are predominantly formed in an accretion disk near the compact star. Hence, with a period of 0.7 days the relative phasing of the velocity curves and light curve are consistent with the proposed model. However, the velocity curves are not very sinusoidal, showing that the period used here is only approximate. An unlikely possibility is that the orbit is noncircular, but until a better value of the period is known this cannot be tested. Velocity semiamplitudes from these curves lie in the range of K 20–30 km s 1. Some lines may also be blended with emission arising in regions such as a disk hot spot or in a gas streams, further complicating the velocity curves. To test the Callanan et al. period, even though it was not preferred in our period analysis, in Figure 7b we have plotted each of our data sets on their value of P = 0.619 days. In
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masses. For a velocity semiamplitude K = 25 km s 1 and P = 0.7 days, the mass function is f(M) = 0.0011 M . In Figure 8 we plot the masses of the two stars for various values of the orbital inclination using this mass function. Since there are no eclipses in either the X-ray or optical data, the inclination must lie below i = 60 . The high X-ray luminosity tells us that the compact star must be either a neutron star or a black hole. Hence, the lower limit for the compact star (MX) has been drawn at 1.4 M . Shahbaz et al. (2001) have set an upper limit on the luminosity of the donor star from its lack of visibility when GX 339-4 was in its faint state. Hence, one can infer that an upper limit on its mass of Mdonor < 0.6 M . These limits are shown as a bold outline in Figure 8. It can been seen that the mass ratio is high and must lie between q 3 and 15. If the compact object is a neutron star, the orbital inclination is very low, with i = 15 –30 . Most of the outlined region implies that the compact star is indeed a black hole, as has been supposed. We point out that, if the donor is a main-sequence star, it does not come close to filling its Roche lobe for any of the solutions shown and hence must be an evolved star. 6. BEHAVIOR OF THE EMISSION PEAKS
Fig. 6.—Periodograms for (top) I-band photometry, (middle) cross˚ ), and (bottom) central H velocities. correlation velocities (4400–5100 A The dashed vertical line corresponds to P = 0.619 days, suggested by Callanan et al. (1992). All of our data sets imply a slightly longer period.
the top panel one can see that photometry from one of the nights does not fit the folded light curve. More importantly, the phasing of the velocity curves is inconsistent with the light curve. Photometric phase zero (presumably when the compact is at superior conjunction) occurs on the rising branch of the velocity curve rather than at the expected phase half a cycle later. This would imply that the emission lines are formed on the donor’s side of the center of mass. In such a case one would expect very narrow emissions and a much higher velocity amplitude. This phasing clearly does not make physical sense, and hence P = 0.619 days does not fit our data. In a further effort to refine the orbital period, we have also looked at the variations of the equivalent widths. While no convincing periods were detected, a plot of W(H) on P = 0.7 days does show a small systematic variation with the emission line being strongest near photometric phase 0.1 and weakest near 0.6.
5. POSSIBLE ORBITAL PARAMETERS AND STELLAR MASSES
If the observed velocity variations are mainly due to motion of the disk that directly surrounds the compact star, we can draw some inferences about the components’
The behavior of the emission peaks is puzzling. Throughout our observations the short-wavelength component of He ii 4686 was the stronger peak in all but a few spectra ˚ peaks separately, (see Table 2). Analyzing the two 4686 A the violet component shows no clear periodicities, while the red peak shows a modulation with a period 0.7 days, similar to that of the broad emission lines. During the 4 days we observed, the separation of these two peaks decreased from 500 to 370 km s 1. At H we were only able to measure a single peak, although with better resolution the line could probably be fitted by two components. The violet peak was predominantly the stronger one, except during part of the second night. Like the He ii red peak, the H peak shows some modulation on a 0.7 day period, with phasing similar to the broad emission lines’ velocities. We can compare our observations with others in the literature. In 1996 May, Smith, Filippenko, & Leonard (1999) found H to be double peaked with a component separation of 370 km s 1. In their data, like ours, the violet peak was stronger. From data taken during three nights in 1998 April, Soria, Wu, & Johnston (1999) found the strong emission lines were double, with separations increasing with excitation. He ii 4686 showed a separation of 480 km s 1, with the red peak stronger in most spectra. They measured double peaks at H, with an average separation of 250 km s 1. Wu et al. (2001) obtained spectra in 1999 April, which show H with a narrow central peak and 4686 double with the stronger violet peak separated from the red peak by 600 km s 1. Hence it is clear that the separation and relative strengths of the emission peaks change considerably, but it is unclear whether this occurs on an orbital timescale or something much longer. Furthermore, it is not yet understood whether these changes are related to periodic effects in the binary (such as orbital motion or possible disk precession), to high or low X-ray states, or to other phenomena. Determination of the orbital period of this perplexing binary may help to sort out these complex variations. There is no evidence in our data that we have observed any line emission from the heated face of the donor star. The implied mass ratio from our analysis above should
Fig. 7.—Left: Velocity and light curves phased on a sample period of 0.7 days for (top) I-band photometry, (middle) cross-correlation velocity from 4400– ˚ region (note this velocity scale is relative rather than true heliocentric velocity), and (bottom) central H velocity. Data from a single night are all shown 5100 A by the same symbol. While we have not determined the actual value of the orbital period, this example demonstrates that periods near 0.7 days show minimum velocity approximately a quarter of a cycle after minimum light, as would be expected for emission lines formed in an accretion disk close to the X-ray source. Right: Similar to left, but showing velocity and light curves on P = 0.619 days, the period proposed by Callanan et al. The relative phasing shows the unlikely condition of having minimum light occur on the rising branch of the velocity curve.
Fig. 8.—Possible masses for GX 339-4 shown for a sample orbital period of 0.7 days and velocity semiamplitude K = 25 km s 1. Lines of constant orbital inclination and mass ratio (q = MX/Mdonor) are shown. The lack of visibility of the donor star, even when the system is in an X-ray–off state, indicates that it must be very faint and hence low mass, probably below 0.6 M . Since no eclipses are seen, an inclination of i = 60 partially defines a lower bound for the donor star. If the X-ray source is a neutron star, the donor lies between 0.3–0.6 M . If the X-ray emitting object is a black hole, its mass could be up to 10 M .
THE BLACK HOLE CANDIDATE GX 339-4 mean that any sharp emissions moving out of phase with the broad lines would be quite obvious.
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sis of the possible masses indicates that the secondary is a very low-mass evolved star and that the X-ray source is probably a black hole.
7. CONCLUSIONS
In summary, although we have not yet been able to determine the orbital period for GX 339-4, the data presented suggest that the binary has a period near 0.7 days. Analy-
We wish to thank the staff at CTIO for their assistance. A. P. C. acknowledges support from the National Science Foundation.
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