Planetary nebulae in M 32 and the bulge of M 31: Line intensities and oxygen abundances

Astronomy and Astrophysics Supplement Series, Jul 2018

We present spectroscopy of planetary nebulae in M 32 and in the bulge of M 31 that we obtained with the MOS spectrograph at the Canada-France-Hawaii Telescope. Our sample includes 30 planetary nebulae in M 31 and 9 planetary nebulae in M 32. We also observed one HII region in the disk of M 31. We detected [OIII]λ4363 in 18 of the planetary nebulae, 4 in M 32 and 14 in the bulge of M 31. We use our line intensities to derive electron temperatures and oxygen abundances for the planetary nebulae.

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Planetary nebulae in M 32 and the bulge of M 31: Line intensities and oxygen abundances

Astron. Astrophys. Suppl. Ser. Planetary nebulae in M 32 and the bulge of M 31: Line intensities and oxygen abundances M.G. Richer 2 G. Stasinska 1 M.L. McCall 0 0 Dept. of Physics and Astronomy, York University , 4700 Keele Street, Toronto, Ontario , Canada M3J 1P3 1 DAEC, Observatoire de Meudon , 5 place Jules Janssen, F-92195 Meudon Cedex , France 2 Instituto de Astronom a, UNAM , Apartado Postal 70-264, 04510 Mexico D. F. , Mexico We present spectroscopy of planetary nebulae gas. Outflow probably begins when supernovae have raised in M 32 and in the bulge of M 31 that we obtained the internal energy of the gas enough to allow it to escape with the MOS spectrograph at the Canada-France-Hawaii the potential well (e.g., Brocato et al. 1990). Telescope. Our sample includes 30 planetary nebulae in M 31 and 9 planetary nebulae in M 32. We also observed one H ii region in the disk of M 31. We detected [O iii] 4363 in 18 of the planetary nebulae, 4 in M 32 and 14 in the bulge of M 31. We use our line intensities to derive electron temperatures and oxygen abundances for the planetary nebulae. galaxies; individual; M 31; M 32 galaxies abundances | galaxies; elliptical and lenticular galaxies; evolution | galaxies; ISM 1. Introduction One of the most important clues concerning the early evolution of dynamically hot galaxies (DHGs: ellipticals, dwarf spheroidals, and bulges of spirals) in the fundamental plane of galaxies is the existence of a well-de ned relationship between metallicity and mass (e.g., Bender et al. 1993) . The fundamental lesson taught by this relation is that star formation in DHGs stopped because of gas loss, with less massive systems losing greater fractions of their Send o print requests to: M.G. Richer ? The overwhelming majority of my work on this project was done while I was at the Observatoire de Meudon as a member of the DAEC. ?? Visiting Astronomer, Canada-France-Hawaii Telescope, operated by the National Research Council of Canada, the Centre National de la Recherche Scienti que de France, and the University of Hawaii. Most commonly, the metallicity in DHGs is measured via the Mg2 index. While the Mg2 index is an excellent means of ranking galaxy metallicities, it does not yield an abundance directly, i.e., the number density of a particular element relative to hydrogen, and calibrations of the Mg2 index (model-dependent) are usually in terms of the iron abundance, an element whose production is notoriously di cult to model. Though this may be best for some purposes, e.g., studies of stellar populations, it is not su cient for all purposes. To study the chemical evolution of DHGs requires the abundance of an element whose production is well understood. Were such abundances available, there would be some hope of quantifying the gas fraction at which DHGs of di erent masses begin to lose mass. Knowledge of the abundances would admit studying the yield of heavy elements, and hence the slope of the stellar initial mass function during the star formation epoch. Given the known photometric and dynamical properties of DHGs today, abundances would also allow us to study the global energetics involved during their star formation phase. This paper is one of a sequence investigating the oxygen abundances of DHGs. Here, we present oxygen abundances for samples of planetary nebulae in M 32 and in the bulge of M 31. These two nearby systems are good representatives of typical DHGs. Though M 32's light pro le may be truncated compared to isolated ellipticals, its structural, dynamical, and spectral properties are perfectly typical for an elliptical of its luminosity (Kormendy 1985; Bender et al. 1993) . Similarly, recent work on the DHG fundamental plane has shown that the photometric, dynamical, and stellar population properties a These are minimum spectral ranges. The actual spectral range will depend upon the object's position within the spectrograph's eld of view. of bulges follow those of pure ellipticals (Bender et al. 1992, 1993) . Oxygen is an excellent element with which to study the evolution of galaxies. Oxygen is a primary element whose sole signi cant production site is type II supernovae (Wheeler et al. 1989) , so its abundance is tied directly to the history of massive star formation, and the enrichment time scale is short compared to the gas consumption time scale. Oxygen abundances are also easily observable in planetary nebulae. Planetary nebulae have high electron temperatures, so the temperature-sensitive [O iii] 4363 line is observable, making it possible to determine accurate electron temperatures in high metallicity environments. Further, the dominant ionization stages of oxygen, O+ and O++, have observable lines, while other ionization stages are easily accounted for using ratios of readily detectable helium lines (e.g., Kingsburgh & Barlow 1994) . Planetary nebulae are good sites in which to probe the oxygen abundance, and they are the only sites that are directly accessible in DHGs. Since planetary nebulae are bright in strong emission lines (e.g., [O iii] 5007), they are easily located within their parent galaxies using emission-line and continuum-band imaging (e.g., Ciardullo et al. 1989) . Observational and theoretical evidence indicates that the stellar precursors of most planetary nebulae do not modify their initial oxygen abundance (Iben & Renzini 1983; Henry 1989; Perinotto 1991; Forestini & Charbonnel 1997) . Hence, a planetary nebula's oxygen abundance reflects that in the interstellar medium at the time of its precursor's formation. Finally, most of the stellar populations in DHGs are old, so they will produce planetary nebulae at comparable rates per unit mass. As a result, planetary nebulae sample the oxygen abundances in DHGs according to the mass in each stellar population. The resulting mean oxygen abundance for the planetary nebula population in a DHG should then be a mass-weighted mean of the oxygen abundances in its stellar populations. Apart from their utility for studying the chemical evolution of M 31 and M 32, the spectroscopic data for the planetary nebulae we present are interesting for what they reveal about the evolution of the planetary nebulae themselves. Though there may exist a good qualitative understanding of planetary nebula evolution, it is unclear how well it stands up to quantitative scrutiny. This situation arises primarily because the distances to planetary nebulae are di cult to establish within the Milky Way. Traditionally, this constraint has made it di cult to study such absolute properties as the luminosity and size of planetary nebulae, as well as the temporal evolution of these quantities. Extragalactic planetary nebulae are especially valuable in this regard because their distances are known. The addition of the data sets for M 32 and the bulge of M 31 is particularly helpful since these planetary nebulae arise from old stellar populations. They will thus provide an intriguing contrast with the planetary nebula populations in the Magellanic Clouds, which are the product of recent star formation (Richer 1993) . Whether the evolution of planetary nebulae depends upon the progenitor mass or metallicity are among the questions that we may hope to answer through a comparison of the properties of planetary nebulae in M 31 and M 32 with those elsewhere. A better quantitative understanding of planetary nebula evolution would be a great help in understanding and using the planetary nebula luminosity function as a distance indicator. In this paper, we present our spectroscopic data for our samples of planetary nebulae in M 32 and in the bulge of M 31. The observations and their reductions are described in Sect. 2. The line intensities and reddenings we deduce are presented in Sect. 3. The reddening-corrected line intensities are then used to calculate electron temperatures and oxygen abundances in Sect. 4. Summary comments are given in Sect. 5. In companion papers, we will use the data we present below to study the chemical evolution of DHGs and the evolution of planetary nebulae in di erent environments. 2. Observations and reductions limitation, since no bias pattern was obvious on either of the rst two nights. Finally, pixel-to-pixel variations were Our ultimate purpose for making these observations was removed using spectra of the internal quartz lamp. to study the chemical evolution of M 32 and the bulge of Extracting the spectra proved challenging on account M 31. In M 31, we chose planetary nebulae in the inner of the nature and faintness of the sources, and on acbulge in order to probe the highest levels of enrichment. count of the characteristics of the spectrograph. The planOn account of the bright galaxy background, we also pref- etary nebulae in M 31 and M 32 are su ciently faint erentially chose planetary nebulae that were known to be that we were unable to detect their continuum emisbright in [O iii] 5007. All of the objects we observed in sion. Only the emission lines were visible, appearing as M 31 are found within the inner half e ective radius of a sequence of dots, so it was impossible to trace these M 31's bulge. In M 32, we observed as many objects as we spectra. Furthermore, the spectra spanned the full width could, again emphasizing bright objects on account of the of the detector, so they su ered from geometric distorgalaxy background. In this case, the objects we observed tion (pin-cushion) introduced by the optics of the specextended to many e ective radii. trograph. Fortunately, we had to include star apertures We obtained our observations over three nights in when de ning the spectrograph's focal plane mask to perAugust 1994 at the Canada-France-Hawaii Telescope mit accurate re-alignment on the eld when ready to (CFHT) with the multi-object spectrograph (MOS). The do spectroscopy. We used these stars (6 for M 31, 3 for MOS is an imaging, multi-slit spectrograph that employs a M 32) to map the geometric distortion imposed by the grism as the dispersing element (see Le Fevre et al. 1994 optics, and corrected this distortion using the tasks in the for details) . Objects are selected for spectroscopy using noao.twodspec.longslit package (Anderson 1987) . At this focal plane masks that are constructed on-line from previ- point, we had images in which the wavelength axis was ously acquired images. The detector was the Loral3 CCD, parallel to the rows of the CCD, and we could use the a thick CCD with 15 m square pixels in a 2048 2048 brightest line in each spectrum to de ne an extraction format, coated to enhance the quantum e ciency in the aperture (e.g., Massey et al. 1992) . Except for the U900 blue. The Loral3's read noise was 8 electrons and its gain spectra, the individual spectra were extracted from each was set to 1.9 electrons/ADU. For the observations of both image and then combined to produce the nal combined M 31 and M 32, we used slits 1500 long by 100 wide. No spectra. To better de ne the extraction apertures for the order-sorting lter was used for any of these observations. U900 spectra, the spectra were combined rst, after ver Table 1 presents a log of our observations. During the ifying that the individual images had the same spatial course of the observations, we used three di erent grism coordinate scales. In all cases, extraction involved local set-ups in order to optimize throughput, wavelength cov- subtraction of the underlying galaxy and sky spectra. erage, and spectral resolution. We used the B600 grism Establishing a consistent sensitivity scale across all only because of the disappointing throughput of the U900 three grism set-ups was a primary consideration of our grism. Although the precise dispersion and wavelength data reduction. We calibrated the instrumental sensitivity coverage depend upon each object's position within the for each set-up using observations of the spectrophotometeld of view, Table 1 lists typical values for all three grisms ric standard stars listed in Table 1. We veri ed that our (minimal ranges for the wavelength coverage). slitlet-to-slitlet sensitivity scale was secure in three ways. We used the standard IRAF routines to reduce the First, the observations of the standard stars were made in data (noao.imred.ccdred), and followed the standard re- pairs through two di erent slitlets. These slitlets were cut duction procedure. First, the overscan bias was removed at the red and blue extremes of the eld of view to enfrom all of the images. Next, for the rst two nights, se- sure that our standard star observations spanned the full quences of zero exposure images were combined and sub- wavelength range of our planetary nebula observations. tracted from the other images to remove any bias pattern. These paired observations of the standard stars had 500 A, This was not done on the third night because the CCD 800 A, and 1900 A of spectrum in common for the U900, dewar began to warm up before we had a chance to ob- B600, and O300 grisms, respectively. In these overlap retain the zero exposure images. This is unlikely to be a gions, the sensitivity functions for each grism (on each night) were in agreement. Second, we obtained a spectroscopic sky flat through the standard star mask with the B600 grism on the last night. This mask contained two slitlets in addition to those used for the standard star observations. Comparing the night sky spectra through these four slitlets indicates that variations in the wavelength sensitivity between di erent slitlets are less than 4.5% (rms). Finally, observations of NGC 6720 were obtained through a di erent mask than the standard stars, and no wavelength-dependent trends are seen in its sensitivity calibration (see Table 3 below). Therefore, though we did not observe the standard stars through the slitlets used for our program objects, we have no reason to believe that our sensitivity calibration is slitlet-dependent. We then chose the O300 observations of the planetary nebulae in M 32 as our reference data set. This choice was motivated by a number of considerations. First, these planetary nebulae were observed with all three grisms. Second, the O300 grism has good sensitivity over the H { H wavelength range (Le Fevre et al. 1994) , which contains the strongest lines in the spectra. Third, our reddening values for these planetary nebulae (see Tables 6, 7, and 8) were reasonable, typically E(B − V ) < 0:2 mag, and invariably positive. These reddenings were consistent with previous observations of PN1 in M 32 (Ford et al. 1978) . The reddening towards M 32 is also expected to be small if it is in front of the disk of M 31 (e.g., Burstein & Heiles 1984) . We ensured that there were no systematic di erences between the B600 and O300 data sets by comparing the intensities of H , H , [O iii] 4959, and He i 5876 measured relative to [O iii] 5007 for the planetary nebulae in M 32. In making these comparisons, we considered only those objects for which we had the best detections of these lines. For these objects, we computed the ratio of the line intensity in the B600 spectrum to that in the O300 spectrum. Table 2 lists the mean value of this ratio, the standard error in the mean, and the objects we considered for each line. Clearly, the main wavelength-dependent trend in Table 2 is a systematic decrease in the B600 sensitivity relative to the O300 sensitivity as one goes to longer wavelengths. Simply tting a line to the values in Table 2 as a function of wavelength, however, yields a rather poor correction at H . As a result, for wavelengths between any two lines found in Table 2, we corrected for the difference in sensitivity calibrations by interpolating linearly between the corrections in Table 2. For lines to the blue of H or to the red of H , we adopted the H or H corrections, respectively. We wondered if the upturn at H in Table 2 could be due to second order contamination, but this seems unlikely. Both the O300 and B600 grisms have very low e ciency at 3250 A, and a second order contamination would a ect the sensitivity calibration for both grisms similarly. Consequently, the upturn at H appears to be real. The corrections in Table 2 were applied Line H Hγ H H H8 H9 H10 H11 H12 H Hγ H H H8 H9 H10 H11 H12 H Hγ H H H8 H9 H10 H11 H12 100:00 42:07 24:03 12:99 7:14 5:79 3:97 2:85 2:32 100:00 42:29 24:06 13:13 7:24 5:59 3:89 2:84 2:19 100:00 42:51 24:72 13:11 7:91 5:58 3:77 3:15 2:60 a The derivation of the uncertainties in the line intensities and reddening is described in Sect. 3. to the spectra of the planetary nebulae in both M 32 and the bulge of M 31. The U900 data required no correction to put them on the O300 sensitivity scale. We deduced this from direct comparison with the B600 and O300 data (Tables 6, 7, and 8), and independently using a spectrum we obtained of the Galactic planetary nebula NGC 6720. Table 3 lists the intensities and reddening values for hydrogen lines in three regions of NGC 6720. The reddening values we derive from Hγ, H , H9, H10, H11, and H12 are in very good agreement in all three apertures, indicating that our U900 sensitivity calibration is good to 3750 A. Our reddening values at H are consistently 0.16 mag lower than calculated from Hγ, so our U900 sensitivities may be underestimated by 15% near 4100 A. Our H8 reddening values are consistently high, but H8 was blended with He i 3889. We corrected the blend for the He i 3889 contribution using the He i 4471 intensity assuming no radiative transfer correction, thereby removing the maximum possible He i 3889 contribution (e.g., Aller 1987) . Thus, it is perhaps not surprising that our H8 reddenings are too high. Overall, our Balmer line intensities for NGC 6720 indicate that our U900 sensitivity calibration is secure from 3750 A to H . Similarly, for the planetary nebulae in M 32 (Tables 6, 7, and 8), the U900 line intensities for [O ii] 3727, [Ne iii] 3869, and He ii 4686 are in excellent agreement with their B600 and O300 counterparts. Figures 1 through 6 display the O300, B600, and U900 spectra of the planetary nebulae in M 32, while Figs. 7 through 12 display the B600 spectra of the planetary nebulae in the bulge of M 31. The object designations 1100 (Ciardullo et al. 1989) are shown next to the spectra. 1000 Normally, the spectra are scaled such that H occupies the 1841000 full intensity scale, so stronger lines from adjacent spec- 1100 tra overlap, but some of the U900 and B600 spectra are scaled such that Hγ and H , respectively, occupy the full a Kingsburgh & Barlow 1994 ICF. intensity scale. This scaling allowed the best compromise in demonstrating the signal-to-noise for various lines and an assessment of the background sky and galaxy subtraction. The full wavelength range is shown for the B600 and U900 spectra, but only the wavelength range below 7350 A is shown for the O300 spectra. Cosmic rays were not removed unless they interfered with the measurement of line intensities, and many remain in the spectra displayed in Figs. 1 through 12. 3. Line intensities and reddening Tables 6 to 9 (at end) list the adopted reddening-corrected line intensity ratios and reddening values for the planetary For the planetary nebulae in M 32, Tables 6 to 8 list nebulae in M 32 and the bulge of M 31. We use the object the reddening-corrected O300, B600, and U900 line indesignations from Ciardullo et al. (1989) . The line intensi- tensities, in addition to our adopted line intensities. The ties were measured using the software described by McCall adopted intensities are those listed under the object name. et al. (1985). The uncertainties quoted for the line ratios Generally, we adopted the U900 line intensities in the blue are 1 uncertainties that incorporate the uncertainties in and the O300 line intensities in the red, with the dividboth the line and H fluxes. The uncertainties in the line ing line being He ii 4686. He ii 4686 is the only common fluxes include contributions from the t to the line itself exception to this rule. For He ii 4686, we normally chose and from the noise in the continuum. In those instances the line intensity from the spectrum in which the line was where there is no line intensity value, but there is a line measured with the lowest relative error. intensity uncertainty, e.g., He ii 4686 in PN5 in M 32, the \uncertainty" is a 2 upper limit to the strength of undetected lines, and is based upon the noise observed in the continuum. Note that PN4 and PN17 in the M 32 eld have radial velocities indicating that they belong to the background disk of M 31 (Ford & Jenner 1975) . The H ii region in the background disk of M 31 that we observed in the M 32 eld is that denoted H ii 1 by Ford & Jenner (1975) . The reddening-corrected line intensities in Tables 6 through 9 are related to those we observed via log I ( ) I (H ) F ( ) F (H ) where F ( ) and I( ) are the observed and reddeningcorrected line intensities, respectively, E(B−V ) is the reddening, and A( ) is the extinction for E(B −V ) = 1:0 mag from the reddening law of Schild (1977) . All of the line intensities for the planetary nebulae in M 32 in Tables 6, 7, and 8 have been corrected for reddening using E(B − V ) determined from the O300 H =H ratio. For the U900 spectra that did not extend to H , we corrected intensities relative to Hγ using the O300 reddening, then adopted I(Hγ)=I(H ) = 0:47. For the planetary nebulae in the bulge of M 31, we determined the reddening from the Wavelength a \3727" denotes the sum of [O ii] 3726, 3729. H =H ratio in the two cases when it was available, Note that the line intensities for PN408 in M 31 are but used the reddening calculated from the Hγ=H ra- not corrected for reddening. For this faint object, we did tio otherwise. In all cases, we assumed intrinsic ratios of not detect Hγ, and H fell outside our spectral window. I(H )=I(H ) = 2:85 and I(Hγ)=I(H ) = 0:47, which are Since our reddenings are based upon di erent line appropriate for an electron temperature of 104 K and an intensity ratios for di erent objects, we consider them electron density of 104 cm−3 (Osterbrock 1989) . The red- in greater detail before proceeding. All of our H dening uncertainties reflect the 1 uncertainties in the H based reddenings in Tables 6 through 9 are positive. or Hγ line intensities. Wavelength The overwhelming majority of our Hγ-based reddenings in (1994), which employs the line intensities of He ii 4686 Table 9 are also either positive or consistent with no red- and He i 5876 to correct for unseen ionization stages of dening, but our 1 Hγ line intensity uncertainties do allow oxygen. Further details may be found in Stasinska et al. negative reddenings in four cases (PN3, PN43, PN48, and (1998). Tables 4 and 5 present two oxygen abundance calPN53). We considered not using Hγ to determine the red- culations. The abundances in Col. 3 are simply the sum dening, but rejected this option for four reasons. First, for of the O+ and O++ ionic abundances. The abundances in the four planetary nebulae in M 32 for which we measured Col. 4 are those from Col. 3 corrected for the ICF. The an Hγ intensity from the B600 spectrum, the reddening- ICF is normally small because He ii 4686 is weak. The corrected Hγ intensity has the expected value of approxi- oxygen abundances in Col. 4 will be adopted in future mately 47% that of H after correcting for reddening us- work. ing the O300 H intensity. In these four cases, then, H In calculating the oxygen abundances, we assumed an and Hγ would yield similar reddenings. Second, our ulti- electron density of 4000 cm−3 in all cases. With electron mate aim is to calculate electron temperatures and oxygen densities of 1 cm−3 and 2 104 cm−3, the oxygen abundance abundances from these line intensities. If we measured the changes by a maximum of −0:02 dex and +0:07 dex, reintensity of [O iii] 4363 relative to Hγ and [O iii] 4959, spectively, for the planetary nebulae in M 31, and by a 5007 relative to H , and assumed I(Hγ)=I(H ) = 0:47, we maximum of −0:03 dex and +0:11 dex, respectively, for would obtain nal intensities for the [O iii] lines that would the planetary nebulae in M 32. be statistically indistinguishable from those obtained by In instances where only upper limits to intensities were correcting for reddening using the Hγ intensity. Applying available, we adopted the following approach. When we a negative reddening correction does a ect the oxygen had upper limits for the intensities of the helium lines abundance we derive by reducing the [O ii] 3727 inten- these limits were used to calculate the ICF. If we did not sity, but this e ect has less impact on the oxygen abun- observe He i 5876 (because it was outside our spectral dance than the uncertainty in the electron temperature window), we made no correction for unseen stages of oxysince there is so little oxygen in the form of O+. Third, gen regardless of the intensity of He ii 4686. (Only in two forcing I(Hγ)=I(H ) = 0:47 via a reddening correction, cases, PN29 and PN30 in M 31, did we detect He ii 4686 even if negative, accounts for any errors in the sensitiv- when He i 5876 was outside our spectral window.) When ity calibration that might otherwise systematically a ect we only had an upper limit to [O iii] 4363, we used this to the [O iii] lines and the subsequent oxygen abundances. derive an upper limit to the electron temperature, and this Fourth, on average, our H - and Hγ-based reddenings temperature limit was then used to derive a lower limit to agree. The mean H -based reddening for all objects (both the oxygen abundance. In these instances, we did not comM 31 and M 32) is E(B − V ) = 0:18 0:04 mag, while the pute an error for either the electron temperature or the mean Hγ-based reddening for all of the planetary nebu- oxygen abundance, and have indicated the results listed in lae in the bulge of M 31 is E(B − V ) = 0:18 0:08 mag, Tables 4 and 5 as limits. When we had an upper limit for if negative reddening values are included, or E(B − V ) = [O ii] 3727, we adopted this limiting intensity for the line. 0:25 0:06 mag, if negative reddening values are set to zero In this case, the O+ ionic abundance is over-estimated, (the uncertainties are the standard errors in the means). but its contribution to the total oxygen abundance was Thus, the reddenings computed from H and Hγ are sim- normally small. ilar. For comparison, the foreground reddening to M 31 Our uncertainties for the electron temperatures and is E(B − V ) = 0:093 0:009 mag (mean of McClure & oxygen abundances reflect the uncertainties in the [O iii] Racine 1969; van den Bergh 1969; and Burstein & Heiles line intensities alone. As noted earlier, reddening in1984) . It is not surprising that the mean reddening for the troduces a further uncertainty through its e ect upon planetary nebulae is 0.10 mag greater than the foreground [O ii] 3727, but this has less influence upon the oxygen value, for planetary nebulae su er additional reddening abundance than the uncertainty in the electron temperadue to internal dust and dust within M 31 and M 32. ture. The electron temperature uncertainty that we quote Consequently, we have chosen to correct for \reddening" is simply the temperature range permitted by the (1 ) even when E(B − V ) is negative. limiting values of the [O iii] line intensities. Similarly, our oxygen abundance uncertainties are derived from the 4. Oxygen abundances abundances calculated using the extreme values of the electron temperature. Tables 4 and 5 present the electron temperatures and the oxygen abundances for the planetary nebulae in M 32 and 5. Discussion in the bulge of M 31, respectively. We only observed two ionization stages of oxygen, O+ and O++. We accounted In Tables 4 and 5, we derive oxygen abundances for apfor unseen stages in our oxygen abundance calculations proximately half of the planetary nebulae we observed. For using the ionization correction factors (ICF) computed the rest, we derive lower limits. Many of the oxygen abunaccording to the prescription of Kingsburgh & Barlow dance limits, however, are very useful. Six of the fourteen temperature limits in Table 5 are below 104 K, and one is even below 9000 K. If we separate the planetary nebulae in Table 5 on the basis of whether they have temperatures or temperature limits, the mean oxygen abundances of the two sets di er at the 92% con dence level, with the set of objects with temperature limits having a higher mean oxygen abundance by at least 0.11 dex. Table 5 shows clearly that we are able to measure oxygen abundances up to approximately the solar value (12 + log(O=H) = 8:93 dex; Anders & Grevesse 1989). Since the Loral3 CCD has only modest sensitivity at [O iii] 4363, the quantum e ciency being about 22%, these results are by no means the limit of what is possible with 4m-class telescopes. In several companion papers, we shall exploit the spectroscopic observations of planetary nebulae in M 32 and in the bulge of M 31 in several ways. First, we intend to study the evolution of the planetary nebulae in these galaxies relative to those in the Milky Way and the Magellanic Clouds. 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M. G. Richer, G. Stasińska, M. L. McCall. Planetary nebulae in M 32 and the bulge of M 31: Line intensities and oxygen abundances, Astronomy and Astrophysics Supplement Series, 203-219, DOI: 10.1051/aas:1999172