THE ASTROPHYSICAL JOURNAL, 498:278È292, 1998 May 1 ( 1998. The American Astronomical Society. All rights reserved. Printed in U.S.A.

STELLAR CONTENT OF THE GALACTIC STARBURST TEMPLATE NGC 3603 FROM ADAPTIVE OPTICS OBSERVATIONS1 F.EISENHAUER,2 A.QUIRRENBACH,2 H. ZINNECKER,3 AND R. GENZEL2 Received 1997 October 1; accepted 1997 December 12 ABSTRACT We present near-infrared adaptive optics imaging of the Galactic starburst template NGC 3603 and its stellar center HD 97950. There is clear evidence for the presence of stars down to 1M_ or less. No cuto† or turnover in the initial mass function is evident. Applying theoretical models of the preÈmain- sequence evolution of intermediate-mass stars to the observed color-color diagram, the color-magnitude diagram, and the luminosity function, we constrain both the age distribution and the initial mass func- tion. Within the systematic errors, this initial mass function follows a Salpeter power law with index ! ¹ [0.73 down to the observational limit of less than 1M_. The stars with less than 4M_ appear to be younger than 106 yr, in contrast to previous age determinations of the high-mass content of HD 97950. Subject headings: Hertzsprung-Russell diagram È H II regions È ISM: individual (NGC 3603) È stars: formation È stars: luminosity function, mass function

1. INTRODUCTION low-mass stars(Zinnecker, McCaughrean, & Wilking 1993), even down to preÈmain-sequence (PMS) objects near the Apart from W49, the H II region NGC 3603 is the most hydrogen burning limit(Hillenbrand 1997). On the other massive giant H II region in our Galaxy(Goss & Radhak- hand, the jury is still out on the R136 cluster in the 30 rishnan1969). It is located in the Carina spiral arm at a Doradus H II region, which is an order of magnitude more distance of 7.2 kpc (1A + 0.035 pc). A cluster of about 50 O luminous than the NGC 3603 cluster. Knowing the PMS stars dominated by the Trapezium-like system HD 97950 low-mass population in the Orion Nebula cluster, NGC provides the ionizing radiation (Clayton 1986; Melnick, 3603, and R136, with roughly one, 10, and 100 massive O Tapia, & Terlevich1989; Mo†at 1983). The age of the NGC stars, respectively, will go a long way to helping us under- 3603 cluster has been estimated to be 2È3 Myr(Melnick et stand whether the presence of an increasing number of al.1989). With three Wolf-Rayet (W-R) star candidates and massive stars somehow quenches the formation of low-mass six O3 stars in a volume less than a cubic light-year (Drissen stars. If so, and only if this is so, the concept of bimodal star et al.1995), the central core of NGC 3603 is the densest formation(GuŽ sten & Mezger 1982; Larson 1986) has a system of high-mass stars known in the Galaxy. Its total quantitative basis beyond speculation. bolometric luminosity is on the order of 107 L _. NGC 3603 thus can be considered as a template for a massive cluster of 2. OBSERVATIONS young stars typical for the building blocks of starbursts in external galaxies. Although its population of high-mass All observations were carried out with the European stars is fairly well known from Hubble Space Telescope Southern Observatory (ESO) adaptive optics (AO) system (HST ) imaging(Mo†at, Drissen, & Shara 1994) and ADONIS at the 3.6 m telescope at La Silla, Chile, and the speckle-masking observations(Hofmann, Seggewiss, & dedicated near-infrared camera SHARP II(R. Hofmann et Weigelt1995), not much is yet known about intermediate- al.1995), developed at the Max-Planck-Institut fuŽ r extra- and low-mass stars. terrestrische Physik, Germany. ADONIS provides a real- Since it is the closest analog in the Galaxy to a starburst, time correction of the atmospheric wave-front distortions, however, its low-mass stellar content is also of considerable and, in the best cases, the di†raction from the telescope interest. This is because there is observational (e.g., Rieke et limits the spatial resolution. As with every AO system, al.1980, 1993) and theoretical (e.g., Silk 1995) evidence that however, the Ðnal image quality depends on the brightness massive star formation regions in starburst systems are deÐ- of the reference star and its distance from the object of cient in low-mass stars (M ¹ 3M_). It is the purpose of interest. We selected the B1.5 supergiant Sher 25 (Brandner this investigation to use deep star counts as a means of et al.1997a, 1997b) as the wave-front reference star, separat- testing this claim. Even though NGC 3603 is only a small ed by about 20A from the stellar center HD 97950 of NGC starburst by extragalactic standards(Kennicutt 1984), it is a 3603. Its visual magnitude of V \ 12.2(Melnick et al. 1989) major stepping stone for establishing whether starburst is just sufficient for a satisfactory image improvement by initial mass functions (IMFs) are truncated or not. ADONIS(V B 13). The camera SHARP II is based on a It is worth mentioning that the Orion Nebula cluster 256 pixel ] Limit256 pixel NICMOS III detector. The observed (L B 105 L _, two O stars) is certainly not deÐcient in wavelength range covers the atmospheric bands J (1.11È bol 1.40 km), H (1.47È1.82 km), and K (1.99È2.32 km). Com- 1 Based on observations collected at the European Southern Observa- pared with the Johnson K band, we used a somewhat tory, La Silla, Chile, ESO 54.E-1000, ESO 56.D-0573, ESO 57.D-0354. narrower Ðlter to suppress as much as possible the thermal 2 Max-Planck-Institut fuŽ r extraterrestrische Physik, Giessenbach- background. The reference stars SA 100-280 and SA 106- strasse, 85740 Garching, Germany; eisenhau=mpe-garching.mpg.de, genzel=mpe-garching.mpg.de, qui=mpa-garching.mpg.de. 1024 provided the photometric calibration. 3 Astrophysikalisches Institut Potsdam, An der Sternwarte 16, 14482 While we observed the T-shaped mosaic(Fig. 1 [Pl. Potsdam, Germany; hzinnecker=aip.de. 3]) of NGC 3603 with a pixel scale of 100 mas pixel~1, the 278 STELLAR CONTENT OF NGC 3603 279 central region(Fig. 2) of 12A.8]12A.8 was also imaged with could not achieve the same AO correction for the mosaic. In a pixel scale of 50 mas pixel~1. The wave-front reference addition, the coarse sampling with pixels equivalent to 100 star Sher 25 can easily be found in the center of the mosaic; mas increased the FWHM. Depending on the distance to the stellar center of NGC 3603 lies about 20A to the south. the wave-front reference star, the FWHM is between 0A.23 The imaged area is 2190 arcsec2. In the detailed images of and0A.33 in the K band, 0A.32 and 0A.42 in the H band, and the stellar center, the total integration time amounts to 8400 0A.31and 0A.39 in the J band. sintheKband, 9000 s in the H band, and 1500 s in the J As we will explain in° 3, the conservative detection limit band. The mosaic was built from eight di†erent Ðelds with for a point source is about 19.3 (18.4) in K, 19.3 (18.9) in H, integration times of 2400 s (K), 1860 s (H), and 1560 s (J) for and 19.2 (18.4) in J (the numbers enclosed in parentheses the three central Ðelds and integration times of 800 s (K), refer to the detailed images of the central region). 620s(H), and 520 s (J) for the Ðve Ðelds north of Sher 25. DATA REDUCTION Each Ðeld itself is composed of multiple frames with inte- 3. gration times between 1 and 5 s. We could not increase this First we subtracted the sky background from the individ- integration time, because the brightest stars (about 8th ual images, applied a Ñat Ðeld to correct for the variation in magnitude in K band) saturated the detector already within the sensitivity across the Ðeld of view, and corrected for about 3 s. dead pixels. In the mosaic we reconstructed the background Since the reference star was simply too faint and too of every single pixel from its minimal intensity in the eight distant from HD 97950, the AO could not provide fully di†erent Ðelds. In the case of the detailed observations in di†raction-limited images. The applied correction reduced the very center of NGC 3603, we observed separate sky the FWHM of a point source in the detailed images to 0A.18 Ðelds. The Ñat Ðeld comes from sky frames recorded during in the K band,0A.30 in the H band, and 0A.31 in the J band. sunrise and sunset. The large variation in the intensities of These numbers and all of the numbers below refer to the the individual frames allows an accurate determination of stacked images containing all the individual frames of short the Ñat Ðeld in the presence of thermal background and integration times. Close to the reference star, the image straylight. A linear interpolation was used to correct the quality was somewhat better; the FWHM was0A.16 (K), dead and noisy pixels. 0A.27(H), and 0A.26 (J). Unfortunately, the seeing on La Silla Owing to mechanical bending and slightly di†erent Ðeld was not monitored systematically during the observations; positions in the various observations, we had to register the we estimate that the atmospheric seeing was about 1A not individual images before co-adding them. We determined including the turbulence within the dome and the telescope the direction and absolute value of the displacement for itself. Because of somewhat worse seeing conditions, we each frame by cross correlation. Before aligning and adding

FIG. 2.ÈCentral12A.8]12A.8 of the stellar center HD 97950 in NGC 3603. The image is on a logarithmic scale. A total integration time of about 8400 s went into this K-band raw image. The FWHM of the point source is 180 mas. The faintest stars traced by the contour lines are about 15.5th magnitude. 280 EISENHAUER ET AL. Vol. 498 the images with an equivalent pixel size of 100 mas, we Ðrst 3. The point source must be located at least0A.5 away regridded them with a linear interpolation by a factor of 2. from the edge of the detector. Otherwise artifacts from alia- The photometric calibration was carried out during the sing may a†ect objects close to the opposite edge of bright night of 1996 April 28, switching several times quickly stars. between the central region of NGC 3603 and the reference 4. The star must be detected simultaneously in the J, H, stars SA 100-280 and SA 106-1024. and K bands, and its position must coincide in all three The next step was to detect the stars and to measure their wavelength bands within less than 150 mas. This criterion Ñuxes. Since the observed Ðeld is extremely crowded, we prevents including deconvolution artifacts in the immediate could not apply simple aperture photometry but had to vicinity of bright point sources in the source list. Such arti- assign the Ñuxes to the stars by deconvolution. From the facts arise mainly from the imperfect knowledge of the true variety of available algorithms we chose the Lucy- PSF. We have already discussed the difficulties in the deter- Richardson algorithm(Lucy 1974). This algorithm provides mination of the local PSF that arise from anisoplanatism a maximum likelihood solution for the inverse convolution and crowding. In addition, ADONIS introduces some tri- problem with known point-spread function (PSF). Unfor- angular coma. This coma results in artifacts in the decon- tunately, the PSF in AO images varies with distance from volution whose orientation and strength vary within the the reference star because the wave fronts from di†erent di†erent images and wavelength bands. The above criterion stars pass through slightly di†erent patches of the atmo- is well suited for rejecting these artifacts. sphere. With increasing distance, this anisoplanatism 5. Last, but not least, the star should be identiÐable by becomes more and more serious and Ðnally limits the Ðeld eye in the raw, not deconvolved, K-band images, which are of view in AO observations. We thus could not use the the observations of highest spatial resolution and limiting image of the wave-front reference star for deconvolving the magnitude. other Ðelds. Especially toward the central star cluster HD Having identiÐed the stars and measured their Ñux in 97950, no bright star is isolated enough to extract the local each Ðeld of the mosaic, we Ðnally combined the source lists PSF. Even the median image from the 70 brightest stars still accordingly. We Ðrst corrected for optical distortion, which shows signiÐcant structure in its outer regions. We thus amounts to about 2% in the corner of a single image. We extended the central part(0A.7 in diameter) of this median then assigned two or more detections within the overlap PSF with the wings of the wave-front reference star. We region of the di†erent Ðelds to a single star if the detections scaled the images in a way that the radial intensity proÐle in were less than 150 mas apart. The standard deviationp in the resulting PSF is continuous with a continuous deriv- m the magnitude m of multiply detected stars served as a ative. We also cleaned the wave-front reference star from measure for the typical photometric error and its depen- faint stars in its vicinity before combining the images. In the dence on the brightness of the star. Besides the statistical same way we generated the PSF in the other Ðelds from at component, this error also includes part of the systematic least 30 stars per image. Then we iterated the Lucy- errors from sky subtraction and Ñat-Ðelding. Not included, Richardson algorithm 10,000 times for every Ðeld and each however, are systematic errors from the deconvolution and of the three wavelength bands J, H, and K. All statements the photometric calibration. We Ðtted a model of the form below will refer to these ““ Lucy maps,ÏÏ if not stated other- p a10bm c to the measured distribution of errors. This wise. m \ ] model was then used to estimate the photometric error in a Within these Lucy maps we subsequently searched for single detection.Figure 3 shows both the measured errors point sources. To avoid mistaking artifacts from the decon- and the associated error model. Similar to the combination volution for stars, every accepted point source had to fulÐll of the di†erent Ðelds in the mosaic, we included the sources several criteria: from the detailed images in our source list. The only di†er- ence was that we introduced a relative magniÐcation and 1. The star must be a local maximum in the Lucy map. rotation as additional parameters. To avoid apparent spatial superresolution, we convolved the Lucy map with a Gaussian of 200 mas FWHM before searching the local maxima. 2. The Ñux of the point source must signiÐcantly exceed the median Ñux in its vicinity. Within a radius of 100 mas the Ñux was assigned to the star. A box of 1A ] 1A deÐned its vicinity. The detection was declared to be signiÐcant if the Ñux in the point source exceeded the median Ñux in its vicinity by the limiting Ñux. We deÐned this limiting Ñux in a very conservative manner. In the raw (not deconvolved) image, the Ñux in the brightest pixel of a point source must be at least 5 times higher than the root-mean-square back- ground noise. This criterion leads to limiting magnitudes of 18.4 in K, 18.9 in H, and 18.4 in J for the detailed obser- vations with the 50 mas pixel~1 scale, and up to 19.3 in K, 19.3 in H, and 19.2 in J for the mosaic. If we deÐned the limiting magnitude from the cross-correlation with the PSF, we would get up to 1.8 mag fainter limits. The latter deÐni- FIG. 3.ÈStatistical photometric errors. The individual points indicate tion, however, can only be applied for areas with well- the standard deviation in the brightness of the stars detected in at least two separated point sources, a criterion which is not at all di†erent Ðeld positions. The best-Ðt model (solid line) provided the estimate fulÐlled in the observations of NGC 3603. for the photometric errors of single detections. No. 1, 1998 STELLAR CONTENT OF NGC 3603 281

In the common Ðeld of our detailed images and the pre- 1. There is an area surrounding each object where one vious HST observations(Mo†at et al. 1994), there are 166 cannot detect stars *m mag fainter than the object. The size stars in total, 40 of them seen by the HST . Only four out of of this area depends on the di†erence *m in brightness and these 40 stars were left unrecognized by our source detec- is described by the critical distance d(*m). tion. Just by eye, however, one Ðnds three of them in our 2. At any given position in the Ðeld, there is a single raw K-band images, which is a strong hint that our detec- dominant neighboring object that inhibits the detection of tion limits are indeed conservative. In the outer region of stars. The inÑuence from several nearby objects does not NGC 3603, the superiority of the new data over HST is sum up. even more prominent. We have detected in total 881 stars We assume implicitly a binary probability for the detection simultaneously in all three wavelength bands J, H, and K. of a star. The model simply predicts if one will detect the In the following discussion we will restrict ourselves to the star or will not. Knowing the critical distance d(*m), one 820 stars located within the three Ðelds of the mosaic closest can easily decide from the spatial distribution and bright- to the stellar center of NGC 3603. This region is marked by ness of the objects in the observed region whether a star of the dotted line inFigure 1. We will ignore the stars in the given brightness can be detected at any given position. In Ðve Ðelds north of Sher 25 because the reduced integration addition, we assume that within the region where a bright time complicates accurate statistics for the very faint stars. object may have prevented the detection of a neighbor in COMPLETENESS CORRECTION AND SYSTEMATIC our observations, the fainter stars are evenly distributed. 4. More explicitly, the probability density for the presence of PHOTOMETRIC ERRORS fainter stars is constant within this region. For the small More than 11 mag lie between the brightest and the faint- areas of interest, with radii of a few arcseconds, we expect est detected stars. With this high contrast, it is no surprise this assumption to be fulÐlled within most of the Ðeld. that the detection of a star is not simply limited by the To derive the critical distance d(*m), we Ðrst calculated background noise but also by blending from neighboring the distribution f (*m, d), which describes how often two bright objects. In particular, toward HD 97950 confusion is stars with a di†erence *m in brightness and a relative dis- the limiting factor. When constructing statistics or fre- tance d were found within the observations. Dividing this quency distributions, we therefore have to account for the frequency by the area of the corresponding annulus, the new possibility of missing stars and apply an appropriate com- distribution F(*m, d) P f (*m, d)/d should converge toward pleteness correction. In addition, the strong contrast within large radii because the stars do not blend with each other the images may corrupt the photometry of faint stars anymore. In contrast, one should Ðnd no pairs of stars with located close to bright objects. distances smaller than the critical distance d(*m). The dis- tribution F(*m, d) will vanish for small values of d. Finally 4.1. Completeness Correction we had to select the appropriate wavelength band. Since we required simultaneous detection in the J, H, and K bands, We approached the completeness correction in two dif- the observations of poorest quality primarily limit the ferent ways: First we used the classical approach using detection. These are the J-band images. They show less Monte Carlo simulations and artiÐcial probe stars. Second, detail and fewer stars than the images in H and K. And we interpreted the measured spatial distribution of the stars indeed, in the J band, the measured distribution F(*m, d) themselves to predict the probability of missing stars. follows the expected behavior.Figure 4 shows this distribu- The Ðrst technique, based on Monte Carlo simulations, is tion F(*m, d). From its decay at small separations we then rather simple. One spreads artiÐcial stars of known inten- sity randomly across the Ðeld of interest, repeats the com- plete data reduction, and analyzes the fraction of probe stars recovered. However, this technique has two disadvan- tages. First, one should use the true PSF to convolve the artiÐcial stars and then deconvolve the images with the measured PSF. Since we only know the measured PSF, we end up with an unsolvable problem; it is exactly the incom- plete knowledge of the PSF that led to the most serious artifacts in the deconvolution. The second disadvantage is the excessive amount of computing time necessary to reduce the simulated data in the same way as the real observations. Owing to the large brightness contrast between the inner part of NGC 3603 and the outer regions in our map, the completeness limit varies strongly across the Ðeld. For a complete Monte Carlo simulation that samples the Ðeld of view on a grid corresponding to 1A and the brightness in steps of 1 mag, we would need about 106 iterations of the Lucy-Richardson algorithm. We avoided this expenditure of computing time by ana- FIG. 4.ÈCritical distance d(*m). The logarithmic gray-scale representa- lyzing the measured spatial distribution of the stars them- tion indicates the frequency F(*m, d) in Ðnding two stars with a di†erence selves and by introducing a simple model for the *m in magnitude and separated by the distance d. The frequency for every distance is normalized to the area of the corresponding annulus. The solid completeness correction in highly crowded Ðelds. The line traces the decay in the distribution toward smaller distances. For model was conÐrmed with a reduced set of Monte Carlo distances smaller than this critical distance d, no stars fainter by *m can be simulations. The basic assumptions of this method follow: found in the vicinity of a brighter object. 282 EISENHAUER ET AL. Vol. 498 determined the function for d(*m), the critical distance. For importance of a single detection close to the stellar center. every di†erence *m in brightness, the critical distance corre- The completeness correction applied in our analysis was sponds to the step function that best Ðts the observed dis- always calculated from a surrounding region 10A in diam- tribution. Assuming a smooth variation of the critical eter. distance with di†erence in brightness, we reduced the sta- To verify whether this new method of completeness cor- tistical Ñuctuations by averaging over 2 mag. The relation rection is valid for our observations of NGC 3603, we com- for the critical distance is shown as a solid line in Figure 4. pared its predictions with the Ðndings of a Monte Carlo Finally we could draw circles around every detected star in simulation. We randomly chose 128 positions within the the Ðeld, indicating how this star limits the detection of central25A.6]25A.6 of NGC 3603 and tried to recover the nearby objects.Figure 5 shows the resulting map of limiting artiÐcial stars implemented there. Just as in the semianalytic magnitudes. model, the step width in brightness was 1 mag. Figure 6 This information was then used to weight individually compares the results of this Monte Carlo simulation and every detection of a star. The weight factor is given by the the above model. The mean-square di†erence is 1.3 mag; reciprocal proportion of the area in its neighborhood where the average of the prediction di†ers by less than 0.8 mag. the detection of a similar star would not have been prevent- This reasonable agreement encouraged us to use the ed by confusion with nearby objects. Since this correction is modeled completeness correction for the subsequent data valid only in regions of evenly distributed stars, and our reduction. Still, the derived completeness correction should observed Ðeld covers both the very dense central core of be regarded more as a lower limit. First, the predicted limit- NGC 3603 and less crowded outer regions, we had to deÐne ing magnitudes exceed systematically the Ðndings from the the completeness correction locally. A single weight factor Monte Carlo simulation for stars fainter than 15th magni- deÐned for the whole Ðeld cannot represent the statistical tude, and we therefore expect the correction to underesti-

FIG. 5.ÈLimiting magnitude in the J band derived from the critical distance and the position and brightness of every star in the Ðeld of view No. 1, 1998 STELLAR CONTENT OF NGC 3603 283

FIG. 6.ÈDi†erence between the limiting magnitudes as predicted by the model and derived from Monte Carlo simulations. Each cross rep- resents the standard deviation of artiÐcial stars within an interval of 1 mag in brightness. The numbers attached to each cross indicate the number of Ðeld positions of corresponding limiting magnitude.

FIG. 7.ÈSystematic photometric errors. The errors were estimated from the di†erence of measured and true brightness of artiÐcial stars. Each mate the number of faint hidden objects. Second, a realistic cross represents the median deviation of artiÐcial stars within an interval of cluster proÐle implies a huge gradient of the projected star 1 mag in brightness. density toward the center, and the assumption of a uniform 18th magnitude. The observed population of faint red stars star distribution within the averaged area no longer holds inFigure 9 is thus not mimicked by systematics in the at the place where most stars are expected. photometry. In the mass function shown inFigure 16, the histogram of the uncorrected number counts is plotted with dotted lines, 5. FIELD STARS while the solid lines refer to the corrected number counts. A priori we cannot rule out a signiÐcant number of Ðeld stars in our images that could systematically a†ect our sta- 4.2. Systematic Errors in the Photometry tistics. Since NGC 3603 is located very close to the Galactic In° 3 we have already discussed the statistical errors in plane (b \[0.53¡), and since we detect stars down to the Ñux measurements. The systematic errors possibly rather faint magnitudes(mJ \ 19.2), we have to address this induced by the aperture photometry on deconvolved issue carefully. On the other hand 92% (704) of the 765 images, however, cannot be estimated from statistical argu- detected stars brighter thanmJ \ 18.6 are located within ments. Although the Lucy-Richardson deconvolution (Lucy the 0.35 arcmin2 considered in our statistics (see Fig. 1). 1974) is known to provide mostly linear photometry in This is a rather small area, and our work should be much crowded regions in the case of the aberrated HST Wide less a†ected by foreground and background contamination Field Camera 1 images(Busko 1994), the possibility of sys- than seeing-limited surveys in nearby star-forming regions, tematic photometric errors in our investigation could still for example the o Ophiuchi cloud and the Taurus molecular exist. We expect the most serious systematic errors in our cloud. photometry in the wings of bright stars. The additional Since no reference Ðeld with comparable sensitivity is photon noise a†ects the measurement, and an incomplete available, we could only obtain a rough estimate of the knowledge of the PSF falsiÐes the deconvolution. Also we expected number of Ðeld stars. The most straightforward might have chosen too small an aperture radius (100 mas) upper limit can be derived from the number of stars in our to collect all the Ñux from a point source. images separated by more than 50A from the cluster center, We investigated these e†ects with the same Monte Carlo assuming that they are all Ðeld stars. Within this area of 244 simulations that were used for the completeness correction. arcsec2 there are 16 stars brighter than the limiting magni- From the comparison of the detected Ñux with the input tude, which ismJ \ 18.6 in this region. If we extrapolate this Ñux we got the systematic errors of the photometry as a value to the limiting magnitude ofmJ \ 19.2 in the Ðeld function of the brightness of the stars.Figure 7 shows these considered in our statistics by assuming a typical slope of systematic errors against the measured brightness for the dNlog (m)/dm B 0.3 in the luminosity function of faint artiÐcial stars. Each cross represents the median deviation Ðeld stars10 (Allen 1976, p. 243), the upper limit results in 357 of artiÐcial stars within a measured brightness interval of 1 Ðeld stars arcmin~2, i.e., 126 stars in the considered Ðeld of mag, centered on integer values. With the exception of stars 0.35 arcmin2. The true Ðeld star density should be much with brightness close to the limiting magnitude, no signiÐ- smaller, because the measured star density is far from con- cant systematic error occurs. In all three wavelength bands vergence toward larger radii in our Ðeld (seeFig. 8). Many the di†erence between measured and true brightness is less of the stars near the edges of our Ðeld are still cluster than 0.07 mag for stars brighter than 15th magnitude. The members. size of the aperture used for the photometry in the decon- Not just an upper limit, but a real estimate for the density volved images is large enough to collect all the Ñux from a of Ðeld stars can be deduced from number counts (Allen star. Closer to the detection limit (m º 16), the artiÐcial 1976, p. 243) in the visible wavelength region. We start from stars appear brighter than they actually are, but the system- the fact that for stars brighter than 10th magnitude the atic errors do not exceed 0.27 mag. In addition, the system- number of Ðeld stars in the blue spectral band follows the atics behave similarly in the three bands, so that almost no visible distribution if one simply shifts the limiting Ñux by 1 systematic errors are induced in the colors. The shift in the mag. Accordingly, we assume the same distribution in the color index J[K is less than 0.05 mag for stars down to near-infrared, with a di†erent shift in limiting magnitude. 284 EISENHAUER ET AL. Vol. 498

We interpret our observations of the intermediate-mass population in NGC 3603 with an approach similar to the one used by the above authors for the high-mass popu- lation. From the measured luminosities and colors, we derive the evolutionary state of the intermediate-mass population.Figure 9 shows the (J[K, J) color-magnitude diagram of the 820 stars detected in the central region of NGC 3603. The colors and magnitudes have been corrected for a global extinction ofAV \ 4.6 and a distance modulus m [ M \ 14.3. The diagram therefore shows the absolute magnitudes and intrinsic colors of the stars if there is no additional individual extinction. The value for the distance modulus was adopted fromMelnick et al. (1989). Compared with his best extinction estimate ofAV \ 4.44 within a radius of 50A from the stellar center, our observations indi- cate a slightly higher extinction. Since our estimate is based mainly on the color index J[K, a small photometric inac- curacy of 0.027 mag in J[K could explain this di†erence in extinction. The conversion from the visible extinction to the near-infrared followsRieke & Lebofsky (1985). At high luminosities the diagram is dominated by a main-sequence FIG. 8.ÈProjected stellar density in NGC 3603. The solid histogram population of high-mass stars. This main sequence seems to includes the completeness correction, while the dotted histogram shows be truncated below 4M_. Instead, a red lower luminosity the actual detections. The best-Ðt isothermal model (solid curve) corre- population is found clearly separated from the main sponds to a core radius of 13A or 0.45 pc. Owing to uncorrectable confusion sequence atM ¹ 0. Bearing in mind that the PMS evolu- in the central region, this core radius must be regarded as a mere upper J limit. tion of a star with ¹2M_ takes º10 Myr, the age of 2È3 Myr of the high-mass stars immediately suggests the PMS nature of this red population if all stars in NGC 3603 are coeval. However, the color-magnitude diagram is not the only We derive this shift from the color of main-sequence stars, indicator of the evolutionary state of a stellar population. B V which also reproduces the shift [ \ 1.0 between the Its state of evolution also shows up in the statistical dis- J blue and visible bands. In the band one therefore must tribution of stellar luminosities.Figure 10 shows the J-band add 2 mag to the limiting Ñux to read the expected number luminosity function (LF) observed in NGC 3603. If the of Ðeld stars from the tables for the visible wavelength range inAllen (1976, p. 243). This results in an estimate of 64 Ðeld stars arcmin~2. Normalized to the area of the Ðeld con- sidered in the star number counts, we expect 23 Ðeld stars. This corresponds to a fraction of less than 3%. By compari- son, the standard deviation of the total number of stars in an ensemble of similar clusters would be 8201@2 stars, or 3.5%. We therefore must consider the Ðeld star population when analyzing the properties of individual stars in the Ðeld but not when investigating the statistical properties of the whole cluster.

6. HISTORY OF STAR FORMATION IN NGC 3603 All previous investigations of the star formation history in NGC 3603 are based on the observations of stars with high luminosities and masses (M º 15M_). The existence of strong W-R features in the stellar center indicates an age of about 2 Myr for the highest mass stars(Santos & Bica 1993). This result agrees with the general distribution of the high-mass stars in the color-magnitude diagram, which follows the isochrone of 2È3 Myr reasonably well (Melnick et al.1989; K.-H. Hofmann et al. 1995). The scatter in this distribution, however, indicates that the higher mass stars did not all form at the same time but spread over a period of about 2 Myr. From the simultaneous existence of the blue supergiant Sher 25, two other blue supergiants in the vicin- FIG. 9.ÈColor-magnitude diagram (J[K, J) of the 820 stars detected ity of the cluster core(Mo†at 1983), and several O3 V stars in the central region of NGC 3603. The diagram has been corrected for global extinction ofAV \ 4.6 and distance modulus m [ M \ 14.3. The (Drissenet al. 1995), Brandner et al. (1997b) infer at least error bars indicate the typical photometric errors. In addition to the two distinct episodes of star formation separated by about empirical main sequence, the diagram also includes the theoretical PMS 10 Myr. isochrones for 0.3 Myr and 3 Myr. No. 1, 1998 STELLAR CONTENT OF NGC 3603 285

FIG. 11.ÈTheoretical color-magnitude diagram of PMS stars. The FIG. 10.ÈJ-band LF of NGC 3603. No completeness correction was dotted curves trace the evolution for given stellar mass, and the solid lines applied in order to avoid artifacts in the shape of the LF. The dip at represent the isochrones for an ensemble of stars. MJ B 0.75 (shaded region) is an indication of the age (0.3È1 Myr) of the underlying stellar population. The upper abscissa shows the corresponding mass of a ZAMS star. shows the isochrones and evolutionary tracks in the near- infrared, calculated on a grid of nine masses from 0.6 to 6 M and eight epochs from the birthline up to 100 Myr. stellar population in NGC 3603 were a simple main- _ In analogy to the age determination in globular clusters, sequence distribution with an underlying power-law IMF, we can derive the age of a very young population from the the LF itself would follow a power law because the mass- truncation of the main sequence. In contrast to the main- luminosity relation in this mass range of main-sequence sequence turno† in globular clusters, however, a truncation stars can be described by a power law as well. The J-band toward low masses is an indicator of the age of NGC 3603. LF of NGC 3603, however, exhibits a distinct local The isochrones for ages of 0.3 Myr and 3 Myr are also minimum for stars of M 0.75. J B shown in the color-magnitude diagram of NGC 3603 (Fig. A quantitative model of the star formation in NGC 3603 9). The distribution of stars in this diagram clearly excludes should reproduce both the dip in the LF and the distribu- an age of º3 Myr for the intermediate-mass population. tion of stars in the color-magnitude diagram. SpeciÐcally the measured truncation in the main sequence occurs at too high masses. In contrast, the isochrones for an 6.1. Color-Magnitude Diagram age of 0.3È1 Myr reproduce quite well the main-sequence While evolution away from the main sequence has to be truncation and the distribution along the PMS branch. considered for the highest mass stars in NGC 3603, we will However, there are two potential problems with this inter- apply PMS models to derive the evolutionary state of the pretation. intermediate-mass population. First, the above models of PMS evolution include only We use the numerical calculations on the PMS evolution the stellar photosphere. PMS stars, however, often display of intermediate-mass stars (0.6È6M_) published by Palla individual reddening and excess emission in the near- & Stahler(1993). Since the cited calculations follow the infrared. The individual reddening results from dust stellar evolution in the space of e†ective temperature and envelopes not yet completely dispersed by the star. In low- bolometric luminosity, we Ðrst have to transform their mass star-forming regions like the Taurus-Auriga complex, results into near-infrared magnitudes and colors. We prefer the extinction of the T Tauri stars is in the range AV B 0È10 the bolometric correction of zero-age main-sequence (Lada& Adams 1992). The near-infrared excess results (ZAMS) stars to the one of giants because PMS stars seem mainly from the accretion disks surrounding the youngest to have surface gravities closer to dwarfs than to giants stars(Lada & Adams 1992). The typical excess emission of (Shiavon,Batalha, & Barbuy 1995). The bolometric correc- the classical T Tauri stars in the Taurus-Auriga complex tion for ZAMS stars was calculated from the bolometric accounts for about *J \ 0.0, *H \ 0.2, and *K \ 0.5 luminosity, visual brightness, and e†ective temperature (Meyer 1996). Both extinction and excess emission result in tabulated inSchmidt-Kaler (1982) and the colors according a horizontal shift in the (J[K, J) color-magnitude diagram. toKoornneef (1983). Stars of less than 6M_ live for at least This shift could lead to a misinterpretation of the properties 10 Myr on the main sequence. Since this is much longer of the stellar population. A set of 3 Myr old stars with than the expected age of NGC 3603, we can neglect their circumstellar disks and/or local extinction could mimic a postÈmain-sequence evolution in our discussion. Figure 11 much younger population. 286 EISENHAUER ET AL. Vol. 498

The standard procedure for estimating the individual 97950, a bubble of about 1 pc (28A) in radius expands with extinction and excess emission is based on the H[K and velocities of about 50 km s~1 into the ambient medium, J[H colors of the stars.Figure 12 shows the color-color probably driven by stellar winds from the central three W-R diagram for the stars in NGC 3603. As in the color- stars(Balick, Boeshaar, & Gull 1980; Clayton 1986, 1990). magnitude diagram, the colors have been corrected for the Such strong activity will have a big impact on the dust global extinction ofAV \ 4.6. We will follow the scheme of envelopes of surrounding stars as well. Winds from the Strom,Strom, & Merrill (1993), dividing the stellar popu- central W-R and O stars may have destroyed these lation into three groups: envelopes. In addition, the central density of more than 105 M pc in HD 97950(K.-H. Hofmann et al. 1995) even 1. Stars with colors of giant or dwarf stars and additional _ ~3 surpasses the one in R136 by a factor of 3(Mo†at et al. reddening from extinction (region I). This group includes 1994; K.-H.Hofmann et al. 1995). Gravitational inter- PMS stars with low excess emission. The weak-line T Tauri action, therefore, may have led to fast disruption of the stars are typical representatives of this subclass of PMS circumstellar disks(Heller 1995), so that we cannot Ðnd the stars. distinct excess emission as known from single classical T 2. Stars that show excess from circumstellar disks in Tauri stars. In summary, we conclude from the near- addition to their photospheric emission. Beside an addi- infrared colors that the distribution in the color-magnitude tional extinction, their location (II) in the (H K, J H) [ [ diagram indeed reÑects the photospheric properties of the diagram is dominated by the excess in the K band. These stars. Ongoing spectroscopic observations hopefully will stars lie to the red of region I. Toward the red, their location provide deeper insight into the processes associated with is limited by the point (H K 1, J H 1), which rep- [ \ [ \ star formation and dispersion of circumstellar material in resents the maximum excess that can be explained by cir- NGC 3603. cumstellar disks(Lada & Adams 1992). This group includes The second criticism refers to the model of star formation the classical T Tauri stars and the Herbig AeBe stars. itself. The results from the PMS model calculations depend 3. Stars with thermal excess emission from surrounding heavily on the accretion rate assumed during the star for- dust envelopes (region III). Their excess is too large to be mation process(Palla & Stahler 1992). This accretion rate is explained by circumstellar disks any longer. They are expected to be on the order of 10 yr , as deter- located even further to the red than stars in region II. ~5 M_ ~1 mined from the model itself and the empirical birthline In the color-color diagram of NGC 3603(Fig. 12) nearly all observed in nearby star-forming regions(Palla & Stahler stars are in region I within the photometric errors. This 1993). There is no evidence yet whether the estimated accre- means that there are no signiÐcant indications of individual tion rate is valid for starburst regions as well. This accretion reddening or distinct excess emission in the K band. Com- rate determines when and at what mass the central burning pared with the Taurus-Auriga complex, this may appear to of deuterium, the shell burning of deuterium, and the be a rather surprising result. However, one should keep in central burning of hydrogen set in(Palla & Stahler 1992). mind that NGC 3603 is not a molecular cloud but an H II For higher accretion rates, the birthline will be shifted region with a very dense star cluster that generates strong toward higher luminosities and lower temperatures. As a winds and ionizing radiation. From the stellar center HD result, one would underestimate the age of the intermediate- mass stars. In addition,Palla & Stahler (1993) used a sim- pliÐed description of the accretion process by assuming photospheric outer boundary conditions to construct the starting models for their PMS calculations. These photo- spheric boundary conditions mimic the loss of the speciÐc entropy of the accreting material through heat radiation from the disk faces. A spherically symmetric collapse, on the other hand, is described by shock-type boundary condi- tions. Luckily, even these two extreme cases lead to only slightly di†erent protostellar evolution(Palla & Stahler 1992). This kind of simpliÐcation can thus be neglected compared with the previous inÑuence of the accretion Ñow. While individual extinction, excess emission, and the outer boundary conditions of PMS evolutionary models seem not to alter our Ðndings essentially, the inaccurate knowledge of the protostellar accretion rate remains a potential problem in the interpretation. 6.2. Age Distribution in NGC 3603 In the previous section we determined the population age by a Ðt of PMS models to the observed distribution in the (J[K, J) color-magnitude diagram. More quantitatively, one can assign to every single star its counterpart in a stellar library. Our library includes the PMS stars of 0.6, 1, 1.5, 2, 2.5, 3, 3.5, 4, 5, and 6M_ with ages of 0.1, 0.3, 1, 3, 10, 30, and 100 Myr, respectively, as adopted fromPalla & Stahler FIG. 12.ÈNear-infrared color-color diagram (H[K, J[H) of the 820 stars detected in the central region of NGC 3603. The diagram has been (1993), and the empirical properties of dwarfs, giants, and corrected for global extinction of AV \ 4.6. supergiants as published byKoornneef (1983) and Schmidt- No. 1, 1998 STELLAR CONTENT OF NGC 3603 287

Kaler (1982). The counterpart in the stellar library has been by the corresponding mass (M)-absolute magnitude (M) chosen such that the weighted distance s2\; relation M 7 M; in the special case of a bijective relation i/J,H,K (Mi [ Milibrary)2/pi2 from the detected star in the luminosity one gets the following: space is minimized.Figure 13 shows the resulting age dis- tribution for stars with M \ 4M . No postÈmain- number of stars with M ½ (M, M ] dM) _ m(M) \ sequence evolution was considered in our stellar library. logarithmic mass interval d log M Stars matching the birthline are displayed in the 0.03 Myr 10 bin. The age distribution reproduces well the best eye Ðt dN(M) \ (1) from the color-magnitude diagram. Most of the stars with d log M 10 M \ 4M_ are 0.3È1 Myr old. The logarithmic average age of these stars is 0.5 Myr. In order to estimate the e†ect of the number of stars with M ½ (M, M ] dM) t(M) \ photometric errors on the derived age distribution, we absolute magnitude interval dM simulated our observations of NGC 3603 numerically. We started from the theoretical near-infrared luminosities for 1 dM P m(M) . (2) 800 stars, whose masses are distributed according to a M dM power-law IMF m P M! of index ! \[0.73, and recalcu- lated the age distribution with additional photometric We apply this relation to constrain the age of the errors. The photometric errors were taken from a Gaussian intermediate-mass stars in NGC 3603. Assuming a coeval probability distribution, with a standard deviation accord- population and a power law for the IMF, we try to repro- ing to the empirical error model discussed in° 3 (Fig. 3). We duce the distinct minimum in the J-band LF near MJ B found both the age spread and the average age consistent 0.75. Stochastic Ñuctuations from a power-law LF that with the real observations. Only stars assigned to ages might have mimicked the observed dip between MJ \ 0 greater than 10 Myr are signiÐcantly over abundant in the andMJ \ 1.5 can be excluded by a s2 test at a signiÐcance observations. Most likely this discrepancy results from the level of 3%. We assumed a simple power-law IMF under- non-Gaussian contribution, for example, from blending lying the stellar population because no observations indi- from neighboring stars to the real photometric errors. cate strong variation in the IMF on the scale of our LF feature. Moreover, we were not interested in reproducing 6.3. Determination of Age through the L uminosity Function meticulously the shape of this feature but in reproducing its In addition to the distribution of stars in the color- position. According toequation (2), a dip in the LF arises magnitude diagram, the LF also(Fig. 10) can be used to from small o dM/dM o; the starsÏ magnitudes (luminosities) constrain the cluster age. We will restrict our analysis to the have to vary rapidly with increasing mass. A steep mass- J-band luminosity, because for classical T Tauri stars the absolute magnitude relation with large o dM/dM o Ñanking light in this wavelength range mainly arises from the stellar the shallow segment will further pronounce the minimum in photospheres(Kenyon & Hartmann 1990) and exhibits no the LF.Figure 14 shows the mass-absolute magnitude rela- signiÐcant excess emission from its surrounding dust disk tion for the birthline, the ZAMS, and for a population built (Meyer 1996). The J-band Ñux thus traces best the bolo- metric luminosity of these PMS stars. For a coeval popu- lation, the relation between IMF m and LF t is determined

FIG. 14.ÈTheoretical mass-absolute magnitude relation of PMS stars with masses ¹6M_. The solid line traces the averaged relation for an age FIG. 13.ÈDistribution of stellar ages in NGC 3603 for stars with of 0.3 Myr and 1 Myr of the PMS models fromPalla & Stahler (1993). The M \ 4M_. No postÈmain-sequence evolution was considered. Stars shaded area highlights the segment of an almost horizontal mass- matching the birthline are displayed in the 0.03 Myr bin. luminosity relation, which is responsible for the dip in the measured LF. 288 EISENHAUER ET AL. Vol. 498 up equally from 0.3 Myr and 1 Myr old stars. We have aroundMJ B 0.75. Even though the LF reÑects a convolu- marked the range ofMJ corresponding to the observed dip tion of the age distribution and a (possibly time-dependent) in the LF. In contrast to birthline or ZAMS stars, 0.3 Myr IMF, and therefore the reconstructed age is certainly a non- and 1 Myr old stars exhibit the desired mass-absolute mag- unique solution, the age estimate from the observed feature nitude relation. The physical reason for this nonlinear shape in the LF independently supports the Ðnding from the in the mass-absolute magnitude relation is the out-of-equi- color-magnitude diagram. librium CNO burning of PMS stars shortly before joining We wish to emphasize the principal di†erence in deter- the ZAMS (A. Belikov 1997, private communication). For mining the age from the LF and from the colors and lumi- masses exceeding 1.25M_, their evolutionary tracks nosities of individual stars. In the latter case one determines describe double luminosity maxima and minima in the Ðrst the age of every single star and then computes the age Hertzsprung-Russel diagram near the main sequence (Iben of the population by averaging the ensemble. This method 1965). ignores the frequency of stars along the corresponding iso- Perhaps even more convincing are the theoretical J-band chrone. Since we estimated the age primarily from the J[K LFs(Fig. 15) constructed from the mass-absolute magni- color index, the result could have been falsiÐed by excess tude relations for various population ages. In all graphs we emission in the K band. The analysis of the LF, however, assumed a power-law IMF m P M! of index ! \[0.73 and ignores the age information for the individual stars and typical photometric errors of 0.3 mag. The comparison with determines the population age solely from features like the measured LF of NGC 3603(Fig. 10) clearly excludes an peaks and dips in the distribution of stellar luminosities. age of 3 Myr for the intermediate-mass stellar population. Since the J-band luminosity does not su†er from excess The simulated LFs for 0.3 Myr to 1 Myr, however, repro- emission of surrounding accretion disks, this method pro- duce quite well the observed frequencies, speciÐcally the dip vides an independent and robust estimate of the population age. Not yet mentioned, but equally important as the intrinsic uncertainties in the models of PMS evolution, is the issue that the concept of coeval star formation becomes question- able for stellar ages of 0.3È1 Myr. Before assigning ages to individual stars in this regime, the zero point has to be clariÐed. Conveniently, stellar ages are related to well-deÐned evo- lutionary states during the starsÏ lifetime: High-mass stars evolve from the zero-age main-sequence, PMS stars are born on the birthline, and protostars start their life at the time the molecular cloud becomes gravitationally unstable to collapse and accretion. In this article we referred to the birthline as the zero point of age because the evolutionary PMS tracks byPalla & Stahler (1993) started from there. For a coeval stellar population, however, the proper zero point should be deÐned by the time a superior action initi- ated the simultaneous formation of stars. As the onset of accretion on the protostar is related closely to the physical origin of star formation, we should add the accretion time to our PMS-evolutionary time to obtain a common measure. In the simple model of a constant accretion rate of 10~5 M_ yr~1, we therefore must add 0.1 Myr for a 1 M_ star and 0.4 Myr for a 4M_ star. Such an accretion rate corresponds to an isothermal sphere with a temperature of 30 K. More detailed calculations byBernasconi & Maeder (1996), which also include turbulent pressure supporting the molecular cloud, predict an increase in accretion rate as protostellar evolution proceeds. The total accretion times, however, are still comparable to the previous estimates. Although accretion times of 0.1È0.4 Myr are not negligible in our discussion of stellar ages in NGC 3603, the basic result that the intermediate-mass stars are younger than about 1 Myr will not be dramatically altered. On the other hand, if cloud collapse happened over timescales of 106 yr and some W-R/O stars with heavy winds and strong radi- ation Ðelds abruptly stopped the accretion within the whole FIG. 15.ÈTheoretical J-band LFs for various ages of a stellar popu- region, the concept of the birthline as the zero point still lation. The distributions have been constructed from a power-law IMF holds as the best common physical event in NGC 3603. m P M! of index ! \[0.73 and for typical photometric errors of 0.3 mag. Even more complicated is the comparison between the The comparison with the measured LF of NGC 3603(Fig. 10) clearly excludes an age of 3 Myr for the intermediate-mass stellar population. The age of low- and intermediate-mass stars determined from simulated LFs for 0.3È1 Myr reproduce quite well the observed fre- PMS tracks and the age assigned to high-mass stars from quencies, speciÐcally the dip around MJ B 0.75. their postÈmain-sequence evolution. Ages of high-mass No. 1, 1998 STELLAR CONTENT OF NGC 3603 289 stars refer conveniently to the ZAMS phase, and therefore agree with each other:K.-H. Hofmann et al. (1995) state the accretion phase is ignored altogether. In addition, the ! \[1.59 ^ 0.22 for the stars with masses M [ 15 M_ concept of a proper ZAMS seems to fail for very massive within their Ðeld of6A.4]6A.4; Mo†at et al. (1994) measured stars(Bernasconi & Maeder 1996) because these stars have ! \[1.4 ^ 0.6 for the stars exceeding 30M_ in their Ðeld already burned a substantial fraction of their central hydro- of 66A ] 66A. Both groups assumed the same age of 3.2 Myr gen content at the time they emerge from their parental to convert the stellar luminosities into masses. For compari- clouds. Ages of high-mass stars should thus be treated with son,Salpeter (1955) derived !B[1.35 for stars with the same care as the ones derived from PMS evolution of masses of 0.4È10M_ in the vicinity of our Sun. More recent intermediate-mass stars. investigations(Scalo 1986) indicate a slightly steeper local Nevertheless, the discrepancy between the age of the IMF with !B[1.7 for stars with masses of 2È10 M_. somewhat less than 1 Myr old intermediate-mass stars and Based on the new photometric data, we can now extend the presumably 2È3 Myr old high-mass stellar content in the IMF in NGC 3603 down to less than 1M_. Since we NGC 3603 is puzzling. According to standard ideas of star know the age of the stellar population, we can interpret the formation, high-mass stars should form last (Herbig 1962; observed J-band luminosities in terms of masses and thus Larson1982, 1985), and they are thus not expected to derive the IMF from the LF. We describe the intermediate- exceed the age of the associated intermediate-mass popu- mass population (M ¹ 6M_) by interpolating the mass- lation. Especially the observations of strong W-R features absolute magnitude relations for 0.3 Myr and 1 Myr (Fig. in NGC 3603(Santos & Bica 1993) and the spectral classi- 14)from Palla & Stahler (1993). The higher mass stars do Ðcation of three stars as WN6]abs(Drissen et al. 1995) not pass through a PMS phase, instead they are born on the seemed to provide a Ðducial age estimate of 2 Myr for the ZAMS and develop very quickly away from it. F. Bertoldi high-mass content. Recent modeling of similar W-R star (1996, private communication) calculated their tracks in the candidates in R136(De Koter, Heap, & Hubeny 1997), near-infrared for solar metallicity Z \ 0.020 by combining however, uncovered the possibility that these stars are the evolutionary models fromShaerer et al. (1993) with O-type main-sequence stars still in the core hydrogen KuruczÏs (1992) models of stellar atmospheres. We use the burning phase of their stellar evolution. They therefore tracks corresponding to 3 Myr to interpret absolute J-band should be at most 2 Myr old, probably optically visible for magnitudes of stars with masses M [ 6 M_. only about 1 Myr. This surprising result is mainly based on The resulting IMF is plotted inFigure 16. The solid his- the presence of intrinsic absorption lines of hydrogen and togram includes the completeness correction; the dotted helium in the spectra of the previous W-R candidates. Since histogram indicates the raw number counts. No systematics such absorption lines are present in the spectra of the three are included in the 1 p error bars; they simply represent the WN6]abs stars in NGC 3603(Drissen et al. 1995) as well, Poisson statistics.Figure 16 also shows the best-Ðt power these stars are not likely as evolved as true W-R stars and law (solid line) with ! \[0.73. This power law reproduces may be signiÐcantly younger than the expected 2 Myr. the observed mass distribution in the mass range from 1 In our comparison of intermediate- and high-mass stars M_ up to 30M_ within the statistical uncertainties. No we assumed the onset of accretion to be the proper age zero cuto† or turnover in the IMF is evident down to the detec- point. Prior to the inside-out collapse, however, the pro- tion limit of less than 1M_. The IMF steepens toward high genitors of low- and intermediate-mass stars need to over- stellar masses. Our best estimate ! \[1.7 of its slope come the magnetic support of the molecular cloud by (dashed line) for masses M º 15M_ agrees well with pre- ambipolar di†usion(Shu et al. 1993). This ambipolar di†u- vious results fromK.-H. Hofmann et al. (1995) and Mo†at sion outÑow phase is expected to last for 1È10 Myr (Lizano et al.(1994). The uncertainty in the age of the intermediate- & Shu1989). In contrast, massive stars were part of mag- mass population of NGC 3603 and the possibly insufficient netically supercritical clumps, which collapse on a free-fall completeness correction dominate the error in the slope of timescale of 0.1È1 Myr(Scott & Black 1980). Even though the IMF for masses M ¹ 6M_. Both the assumption of a the following accretion takes longer for stars of higher Ðnal somewhat higher age of the stellar population and a more mass, the delay from the ambipolar di†usion phase may complete detection of fainter stars in the central crowded mimic a younger age of the low- and intermediate-mass region would steepen the IMF. The value ! \[0.73 stars. should thus be regarded as an upper limit. Since the sta- Taking into account the uncertainties in the age determi- tistical photometric errors result primarily in a smoothed nation of both high- and intermediate-mass stars and the luminosity and mass function, their contribution to the lack of an appropriate common age zero point, we therefore error in the overall slope of the IMF is negligible compared would not exclude the possibility of coeval star formation in with the systematic uncertainties. NGC 3603 starting 1È2 Myr ago. 7.2. Comparison with Other Regions of High-Mass Star Formation 7. INITIAL MASS FUNCTION IN NGC 3603 Among the most interesting open problems in the theory 7.1. Extension to Intermediate-Mass Stars of star formation is the question whether regions of high- Prior to this work, the IMF in NGC 3603 had only been mass star formation also form large numbers of low-mass measured for high-mass stars. BothMo†at et al. (1994), stars(Zinnecker et al. 1993; Zinnecker 1996). There has with HST imaging, andK.-H. Hofmann et al. (1995), with been evidence for a truncated IMF both within the Galaxy speckle interferometry, found the IMF in agreement with a and in extragalactic starbursts(Scalo 1986). The most con- power law m P M!, which was Ðrst proposed by Salpeter clusive hint for such a truncated IMF toward lower masses (1955) for the solar vicinity. These measurements di†er in our Galaxy arises from the observed element abundance slightly in their Ðelds of view and lower mass limits, but the gradients with galactocentric distance. In order to repro- values for the exponent ! derived from the two studies duce especially the large 16O abundance variation with a 290 EISENHAUER ET AL. Vol. 498

cated IMFs has been questioned in all regions of high-mass star formation studied so far. Young Galactic open clusters, which are the smallest clusters including high-mass stars, show star formation down to the observational complete- ness limit (M [ 1.4M_; Phelps & Janes 1993). Also the Orion Nebula cluster with its Trapezium system is certainly not deÐcient in low-mass stars(Zinnecker et al. 1993). The IMF is found to be rising down to 0.2M_ from number counts of spectroscopically classiÐed stars (Hillenbrand 1997). Most recent modeling of the properties of the M82 starburst, including evidence for a more complex spatial- temporal evolution and up-to-date dynamical data, also does not require a lower mass cuto† º1etM_(Satyapal al.1997; FoŽ rster-Schreiber et al. 1998). A Salpeter-like IMF from 0.1È100M_ seems consistent with the observed properties. In addition, high angular resolution imaging has resolved extragalactic giant starburst regions like M82A into individual superstar clusters(OÏConnell et al. 1995) that are supposed to be the progenitors of globular clusters. Ho& Filippenko (1996) have determined dynamical masses for some of these superstar clusters in NGC 1569 and NGC 4214. The similarity of these superstar clusters with globular clusters in terms of total mass and mass-to-luminosity ratio FIG. 16.ÈInitial mass function of NGC 3603. The mass of every indi- vidual star was derived from its absolute J-band magnitude. We calculated probably implies that these starburst regionsÈas globular the mass-absolute magnitude relation of the intermediate-mass stars clustersÈcontain many low-mass stars(Ho & Filippenko (M ¹ 6M_) by interpolating the PMS evolutionary models of Palla & 1996). Stahler(1993) for 0.3 Myr and 1 Myr. High-mass stars were modeled by a Nevertheless, all the knowledge on the low-mass stellar coeval population with an age of 3 Myr(Shaerer et al. 1993; Kurucz 1992; F. Bertoldi 1996, private communication). The solid histogram includes the content in extragalactic starburst regions is based on indi- completeness correction for missing stars in the crowded central region. rect evidence from integrated properties. The Ðnal conclu- The dotted curve shows the actual detections. No cuto† or turnover in the sion as to whether starburst regions form low-mass stars in IMF is evident down to the observational limit of less than 1M_. A large numbers can only be drawn from star counts. The Salpeter power law (solid line) with index ! \[0.73 reproduces the largest H II region resolvable in its intermediate-mass star observed mass distribution in the mass range from 1M_ up to 30M_. On account of possible systematic uncertainties, the slope ! \[0.73 should population is 30 Doradus in the Large Magellanic Cloud be regarded as an upper limit. Our best estimate of the slope (dashed line) with its stellar cluster R136. No cuto† is apparent down to for stars with masses M º 15M_ is ! \[1.7. B3M_ (Hunter et al. 1995; Brandl et al. 1996). Unfor- tunately, present observational techniques do not provide the spatial resolution and sensitivity to extend the know- self-consistent model of star formation and chemical evolu- ledge on the IMF toward lower masses. The cuto† suspect- tion of the Galaxy,GuŽ sten & Mezger (1982) postulated the ed in M82 is still beyond our observational capabilities. For IMF in spiral arms to be truncated at masses below M B 2 the time being, the Galactic starburst template NGC 3603 is M_. In contrast, no low-mass cuto† is needed down to 0.1 thus the best object to constrain the theory of star forma- M_ in the interarm regions. It thus appears to be a charac- tion in starburst regions. To the extent that our result is teristic of the increased star formation rate that it produces representative for all starbursts, high-mass star formation primarily medium- and high-mass stars. Rieke et al. (1980, does not necessarily imply a truncated IMF. Our high 1993) concluded that there is a low-mass cuto† in the IMF spatial resolution imaging and deep photometry proves at masses M B 3M_ in the starburst galaxy M 82. Other- that stars form down to at least one solar mass. wise they could not explain simultaneously the stellar lumi- Our result is fully consistent with recent studies ranging nosity, the dynamical mass, and the high Lyman continuum from the small open clusters up to the highest mass super- photon production rate. Since the early 1980s, this result is star clusters.Table 1 summarizes the properties of typical one of the most cited pieces of evidence for a truncated IMF objects covering the whole mass range of star-forming in starbursts. regions. The luminosities in the visible wavelength band In conjunction with this observational work, the theoreti- and in Ha have been corrected for extinction. Both in terms cal framework(Silk 1977; Larson 1985) for a low-mass of its luminosity as well as regarding its OB star content, cuto† in star formation was developed. The formation of NGC 3603 is closer to the extragalactic starburst R136 and stars will heat the surrounding gas, and thus the Jeans mass the superstar clusters in M82A than to the Orion Nebula in nearby star-forming regions will increase as well. cluster. In the visible, the Orion Nebula cluster (MV \ Depending on the detailed scenario, this process will lead [5.1) is more than 4 mag fainter than NGC 3603 (MV \ either to a Ðxed lower mass in the IMF or a varying lower [9.6). On the other hand, NGC 3603 is almost as luminous mass cuto† with star formation rate(Scalo 1986). In as some of the fainter superstar clusters in M82A (MV \ summary, both observations and theoretical work led to the [9.6 to [13.2;OÏConnell et al. 1995). With its approx- opinion that the IMF is truncated in all regions of high- imately 20 equivalent O5 V stars necessary to ionize the mass star formation and especially in starbursts. associated H II region, NGC 3603 also surpasses the Orion With the development of large and sensitive optical and Nebula cluster by a factor of 100 in its Lyman continuum infrared detector arrays, however, the paradigm of trun- photon output. Only 1 order of magnitude is missing for the No. 1, 1998 STELLAR CONTENT OF NGC 3603 291

TABLE 1 PROPERTIES OF HIGH-MASS STAR-FORMING REGIONS A.

STELLAR CLUSTER

Upper Lower Mass Mass Distance Diameter L Va M M b Age Cuto† Cuto† OBJECT (pc) (pc) M a (L )(MStars)(MOB) (Myr) ! (M )(M) V V,_ _ _ _ _ Young open clusters .... Galaxy 1.1È8.7c ...... 12È630d 7È42c[0.4 to [1.8c 8È26d \1d OB associations ...... Galaxy ...... 100È1600e 0È11f,g [0.7 to [2.0f 15È120f \5f Orion Nebula cluster . . . 470h 5i[5.1j 9.7]103 1850i 155k ¹1i Scaloi 50i 0.2i NGC 3603 ...... 7.2]103l \4.2l[9.6m 6.1]105 ... 2.6]103n \3 [0.7o to [1.7p 60q \1 R136/NGC 2070 ...... 5.0]104r 5.2s,t [11.3t 2.9]106 ... 1.7]104u 3.5v[1.6v 100w \3v M82A ...... 3.6]106x 130y[18z 1]109 1.3]108aa ... 5bbÈ30cc Salpeter 50bbÈ100cc 0.1bb or 3cc M82A superclustersgg ... 3.5ee,z [9.6 to 6.1]105È ...... [13.2z 1.7]107 NGC 1569 A ...... 2.5]106t 4.4s[14.1t 3.9]107 3.3]105ff ... º15t ......

B.

ASSOCIATED H II REGION

Diameter L a M Equivalent M b a ` OBJECT (pc) (ergsH s~1)(MH ) log N a (M )OB Equivalent N(O5 V) _ 10 Lyc _ Orion Nebula cluster ...... 5hh 1]1037hh 50hh 48.85hh 35hh (0.2)hh NGC 3603 ...... 100hh 1.5]1039hh 3.9]104hh 51.05hh 5]103hh 20hh R136/NGC 2070 ...... 1.2]1040ii . . . 51.94jj 4]104ii 170kk M82A ...... 450hh 4]1040hh º2.5]105hh 52.45hh 1.3]105hh 600hh M82A H II complexesgg ...... 52.2bb 7]104ii 300kk

a Corrected for extinction. b Stars with spectral type earlier than B2. c Phelps& Janes 1993. d Adopted from mass functions (Figs. 4, 5, 8, 9, 10, 11, 13) inPhelps & Janes 1993. e Calculated from the number of OB stars(6È93)f and the typical mass of a B0 V star (17.5 M_). f Massey,Johnson, & De Gioia-Eastwood 1995. g Stars with M [ 25 M_. h Genzelet al. 1981. i Hillenbrand 1997. j V \ 5, AV \ 1.7i for stars earlier B2. k CalculatedNGC1976 from Table 5 in Hillenbrand 1997. l Melnicket al. 1989. m V \ 9.1, AV \ 4.44.l n WithinNGC3603 a diameter of 2 pc. o M Z 1 M_. p M Z 15 M_. q Initial mass of hydrogen-rich WNL stars;Langer et al. 1994. r Panagiaet al. 1991. s Half-light diameter. t OÏConnell,Gallagher, & Hunter 1994. u Stars with M [ 9M_ within 40 pc]60 pc in NGC 2070 excluding R136; calculated from Table 2 inParker & Garmany 1993. v Brandlet al. 1996. w DeKoter, Heap, & Hubeny 1997. x Freedmanet al. 1994. y From Fig. 1 (lower panel)inOÏConnell etal. 1995. z OÏConnellet al. 1995. aa Calculated fromL V \ 1.3]109L V _ andM/L V B 0.1 M_/L V _ inOÏConnell et al. 1995. bb FoŽ rster-Schreiberet al. 1998. , , cc Satyapalet al. 1997. dd Riekeet al. 1993. ee FWHM after deconvolution. ff Ho& Filippenko 1996. gg Properties of the stellar clusters and the H II regions in M82 A refer to di†erent objects. hh Kennicutt 1984. ii Calculated fromlog N analogous to Kennicutt 1984. jj Within a diameter of10 140Lyc pc; Walborn 1991. kk Panagia 1973. starburst cluster R136 in 30 Doradus (170 equivalent O5 V With NGC 3603 being the last step on the ladder of stars) and the H II complexes in M82A (300 equivalent O5 V high-mass star-forming regions in our Galaxy that show no stars,FoŽ rster-Schreiber et al. 1998). NGC 3603 thus sets the star count evidence for a low-mass cuto† in the IMF down Galactic cornerstone toward extragalactic starburst regions. to the completeness limit, our Ðndings suggest a non- 292 EISENHAUER ET AL. truncated IMF in all starburst systems. Although the jury is On the other hand, previous investigations(Melnick et al. still out on extragalactic starbursts, more and more evi- 1989; Santos& Bica 1993; K.-H. Hofmann et al. 1995) dence is consistent with a universal IMF similar to the one derived an age of about 2È3 Myr for the high-mass stars in of the Ðeld star population. NGC 3603. Nevertheless, a scenario of coeval star forma- SUMMARY tion may hold in the face of the systematic uncertainties 8. both in PMS evolutionary tracks and in the age determi- Based on AO observations we have performed near- nation of high-mass stars. infrared photometry in the J, H, and K bands of about 900 Starting from the evolutionary state as outlined above stars in the central region of NGC 3603. Since the observed and the appropriate mass-luminosity relation obtained Ðeld is extremely crowded and the faint stars cannot be from the PMS models, we have transformed the measured detected in the bright core, we have corrected the number of luminosities into stellar masses. The IMF shows no turn- detected stars by a semianalytical model. The model itself over or truncation down to at least 1 M_. has been veriÐed by Monte Carlo simulations. A Salpeter power law with index ! \[0.73 reproduces The color-magnitude diagram clearly reveals a popu- the observed mass distribution in the mass range from 1 lation of red, low-luminosity stars, well separated from the M_ up to 30M_. On account of possible systematic uncer- main sequence of the high-mass stars. According to the tainties, the slope ! \[0.73 should be regarded as an present models of PMS evolution, the age of this red, upper limit. Our best estimate of the slope for stars with intermediate-mass population is between 0.3 Myr and 1 masses M º 15M_ is ! \[1.7. Myr. This age is also supported by a pronounced dip in the LF.

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