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2 Observations and reductions

Our ultimate purpose for making these observations was to study the chemical evolution of M 32 and the bulge of M 31. In M 31, we chose planetary nebulae in the inner bulge in order to probe the highest levels of enrichment. On account of the bright galaxy background, we also preferentially chose planetary nebulae that were known to be bright in [O III]$\lambda$5007. All of the objects we observed in M 31 are found within the inner half effective radius of M 31's bulge. In M 32, we observed as many objects as we could, again emphasizing bright objects on account of the galaxy background. In this case, the objects we observed extended to many effective radii.

We obtained our observations over three nights in August 1994 at the Canada-France-Hawaii Telescope (CFHT) with the multi-object spectrograph (MOS). The MOS is an imaging, multi-slit spectrograph that employs a grism as the dispersing element (see Le Fèvre et al. 1994 for details). Objects are selected for spectroscopy using focal plane masks that are constructed on-line from previously acquired images. The detector was the Loral3 CCD, a thick CCD with $15\,\mu$m square pixels in a $2048\times 2048$ format, coated to enhance the quantum efficiency in the blue. The Loral3's read noise was 8 electrons and its gain was set to 1.9electrons/ADU. For the observations of both M 31 and M 32, we used slits 15$^{\prime\prime}$ long by 1$^{\prime\prime}$ wide. No order-sorting filter was used for any of these observations.

Table 1 presents a log of our observations. During the course of the observations, we used three different grism set-ups in order to optimize throughput, wavelength coverage, and spectral resolution. We used the B600 grism only because of the disappointing throughput of the U900 grism. Although the precise dispersion and wavelength coverage depend upon each object's position within the field of view, Table 1 lists typical values for all three grisms (minimal ranges for the wavelength coverage).


  
Table 1: Observing log

\begin{tabular}
{lllllll}
\hline
\noalign{\smallskip}
Date & Object & Grism & Di...
 ...$~s & \object{Feige 15}\vspace*{2mm} \\ \noalign{\smallskip}
\hline\end{tabular}
$^{\mathrm{a}}$ These are minimum spectral ranges. The actual spectral range will depend upon the object's position within the spectrograph's field of view.

We used the standard IRAF routines to reduce the data (noao.imred.ccdred), and followed the standard reduction procedure. First, the overscan bias was removed from all of the images. Next, for the first two nights, sequences of zero exposure images were combined and subtracted from the other images to remove any bias pattern. This was not done on the third night because the CCD dewar began to warm up before we had a chance to obtain the zero exposure images. This is unlikely to be a limitation, since no bias pattern was obvious on either of the first two nights. Finally, pixel-to-pixel variations were removed using spectra of the internal quartz lamp.

Extracting the spectra proved challenging on account of the nature and faintness of the sources, and on account of the characteristics of the spectrograph. The planetary nebulae in M 31 and M 32 are sufficiently faint that we were unable to detect their continuum emission. Only the emission lines were visible, appearing as a sequence of dots, so it was impossible to trace these spectra. Furthermore, the spectra spanned the full width of the detector, so they suffered from geometric distortion (pin-cushion) introduced by the optics of the spectrograph. Fortunately, we had to include star apertures when defining the spectrograph's focal plane mask to permit accurate re-alignment on the field when ready to do spectroscopy. We used these stars (6 for M 31, 3 for M 32) to map the geometric distortion imposed by the optics, and corrected this distortion using the tasks in the noao.twodspec.longslit package (Anderson 1987). At this point, we had images in which the wavelength axis was parallel to the rows of the CCD, and we could use the brightest line in each spectrum to define an extraction aperture (e.g., Massey et al. 1992). Except for the U900 spectra, the individual spectra were extracted from each image and then combined to produce the final combined spectra. To better define the extraction apertures for the U900 spectra, the spectra were combined first, after verifying that the individual images had the same spatial coordinate scales. In all cases, extraction involved local subtraction of the underlying galaxy and sky spectra.

Establishing a consistent sensitivity scale across all three grism set-ups was a primary consideration of our data reduction. We calibrated the instrumental sensitivity for each set-up using observations of the spectrophotometric standard stars listed in Table 1. We verified that our slitlet-to-slitlet sensitivity scale was secure in three ways. First, the observations of the standard stars were made in pairs through two different slitlets. These slitlets were cut at the red and blue extremes of the field of view to ensure that our standard star observations spanned the full wavelength range of our planetary nebula observations. These paired observations of the standard stars had 500 Å, 800 Å, and 1900 Å of spectrum in common for the U900, B600, and O300 grisms, respectively. In these overlap regions, 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 different slitlets are less than 4.5% (rms). Finally, observations of NGC 6720 were obtained through a different 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$\beta$ - H$\alpha$ wavelength range (Le Fèvre 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)<0.2mag, 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 differences between the B600 and O300 data sets by comparing the intensities of H$\beta$, H$\alpha$, [O III]$\lambda$4959, and He I$\lambda$5876 measured relative to [O III]$\lambda$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.

  
Table 2: B600/O300 sensitivity correction

\begin{tabular}
{lcl}
\hline \noalign{\smallskip} Wavelength &
$I$(B600)/$I$(O30...
 ... PN8, PN11, PN17, PN24 \vspace*{3mm}\\  \noalign{\smallskip}
\hline\end{tabular}
$^{\mathrm{a}}$ PN4 and PN17 are background objects in the disk of M 31 (Ford & Jenner 1975).


  
Table 3: Hydrogen line intensities for NGC 6720

\begin{tabular}
{llrl}
\hline
\noalign{\smallskip}
 & Line & Intensity${}^{\math...
 ....37$\space & $0.184 \pm 0.151$\space \\ \noalign{\smallskip}
\hline\end{tabular}
$^{\mathrm{a}}$ The derivation of the uncertainties in the line intensities and reddening is described in Sect. 3.


  
Table 4: Oxygen abundances in M 32

\begin{tabular}
{llll}
\hline
\noalign{\smallskip}
Object & $T_{\rm e}$\space & ...
 ...$\gt 7.84$\space & $\gt 7.92 $\space \\ \noalign{\smallskip}
\hline\end{tabular}
$^{\mathrm{a}}$Kingsburgh & Barlow 1994 ICF.
${}^{\mathrm{b}}$ PN4 and PN17 are background objects in the disk of M 31 (Ford & Jenner 1975).


  
Table 5: Oxygen abundances in M 31

\begin{tabular}
{llll}
\hline \noalign{\smallskip}
Object & $T_{\rm e}$\space & ...
 ....38 \pm 0.09$\space & $8.39 \pm 0.09$\\ \noalign{\smallskip}
\hline\end{tabular}
$^{\mathrm{a}}$Kingsburgh & Barlow 1994 ICF.

Simply fitting a line to the values in Table 2 as a function of wavelength, however, yields a rather poor correction at H$\alpha$. 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$\beta$ or to the red of H$\alpha$, we adopted the H$\beta$ or H$\alpha$corrections, respectively. We wondered if the upturn at H$\alpha$in Table 2 could be due to second order contamination, but this seems unlikely. Both the O300 and B600 grisms have very low efficiency at 3250 Å, and a second order contamination would affect the sensitivity calibration for both grisms similarly. Consequently, the upturn at H$\alpha$ appears to be real. The corrections in Table 2 were applied 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 H$\gamma$, H$\epsilon$,H9, H10, H11, and H12 are in very good agreement in all three apertures, indicating that our U900 sensitivity calibration is good to 3750 Å. Our reddening values at H$\delta$ are consistently 0.16mag lower than calculated from H$\gamma$, so our U900 sensitivities may be under-estimated by 15% near 4100 Å. Our H8 reddening values are consistently high, but H8 was blended with He I$\lambda$3889. We corrected the blend for the He I$\lambda$3889 contribution using the He I$\lambda$4471 intensity assuming no radiative transfer correction, thereby removing the maximum possible He I$\lambda$3889 contribution (e.g., Aller 1987). Thus, it is perhaps not surprising that our H8 reddenings are too high. Overall, our Balmer line intensities for NGC 6720 indicate that our U900 sensitivity calibration is secure from 3750 Å to H$\beta$. Similarly, for the planetary nebulae in M 32 (Tables 6, 7, and 8), the U900 line intensities for [O II]$\lambda$3727, [Ne III]$\lambda$3869, and He II$\lambda$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 (Ciardullo et al. 1989) are shown next to the spectra. Normally, the spectra are scaled such that H$\beta$ occupies the full intensity scale, so stronger lines from adjacent spectra overlap, but some of the U900 and B600 spectra are scaled such that H$\gamma$ and H$\alpha$, respectively, occupy the full 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 Å 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.

  
Table 6: Line intensities for PNe in M 32
\begin{table}
$^{\mathrm a}$\ \lq\lq 3727'' denotes the sum of
[{O} {II}]$\lambda\lambda$3726, 3729.\end{table}


  
Table 7: Line intensities for PNe in M 32 (continued)


  
Table 8: Line intensities for PNe in M 32 (continued)
\begin{table}
\hspace*{1.5cm}$^{\mathrm a}$\ \lq\lq 3727'' denotes the sum of
[{O} {II}]$\lambda\lambda$3726, 3729.\end{table}

  
\begin{figure}
\includegraphics [angle=90,width=13cm,height=9cm]{oap1to6.eps}\end{figure} Figure 1: The O300 spectra for PN8, PN11, PN2, PN7, H II1, and PN25 in the M 32 field. The spectra are displayed such that H$\beta$ spans the entire free intensity scale. Consequently, lines stronger than H$\beta$ overlap in adjacent spectra. We show only the spectral range blueward of 7350 Å. In all of the spectra we present, cosmic rays were not removed unless they interfered with measuring line intensities, so many obviously remain. PN25 is very close to M 32's nucleus, so the sky subtraction is poorer for this object. H II1 is an H II region in the background disk of M 31 (Ford & Jenner 1975)

  
\begin{figure}
\includegraphics [angle=90,width=13cm,height=9cm]{oap7to12.eps}\end{figure} Figure 2: The O300 spectra for PN24, PN6, PN5, PN1, PN4, and PN17 in the M 32 field. The format is identical to Fig. 1. Like PN25, PN24 is also very close to M 32's nucleus and suffers from somewhat poorer background subtraction. Note that PN4 and PN17 are background planetary nebulae in the disk of M 31 (Ford & Jenner 1975)

  
\begin{figure}
\includegraphics [angle=90,width=13cm,height=9cm]{bap1to6.eps}\end{figure} Figure 3: The B600 spectra for PN8, PN11, PN2, PN7, H II 1, and PN25 in the M 32 field. For PN8, PN11, and PN25, the scaling is such that H$\alpha$, not H$\beta$, defines the free intensity range. The full useful wavelength range of the spectra, 3690 Å to 6750 Å, is shown. See Fig. 1 for comments on individual objects

  
\begin{figure}
\includegraphics [angle=90,width=13cm,height=9cm]{bap7to12.eps}\end{figure} Figure 4: The B600 spectra for PN24, PN6, PN5, PN1, PN4, and PN17 in the M 32 field. For PN24 and PN17 (M 31), H$\alpha$, and not H$\beta$, defines the free intensity range. The full useful wavelength range of the spectra, 3690 Å to 6750 Å, is shown. See Fig. 2 for comments on individual objects

  
\begin{figure}
\includegraphics [angle=90,width=13cm,height=9cm]{uap1to6.eps}\end{figure} Figure 5: The U900 spectra for PN8, PN11, PN2, PN7, H II 1, and PN25 in the M 32 field. Only for PN11 and PN25 does H$\beta$define the full intensity scale. For the other objects, the full intensity scale is defined by H$\gamma$. The full useful wavelength range, 3690 Å to 5050 Å, is displayed. See Fig. 1 for comments on individual objects

  
\begin{figure}
\includegraphics [angle=90,width=13cm,height=9cm]{uap7to12.eps}\end{figure} Figure 6: The U900 spectra for PN24, PN6, PN5, PN1, PN4, and PN17 in the M 32 field. H$\beta$ defines the full intensity scale for PN24 and PN17, but H$\gamma$ does so for the other objects. The full useful wavelength range, 3690 Å to 5050 Å, is displayed. See Fig. 2 for comments on individual objects

  
\begin{figure}
\includegraphics [angle=90,width=13cm,height=9cm]{ap1to5.eps}\end{figure} Figure 7: The B600 spectra for PN172, PN31, PN80, PN30, and PN29 in the M 31 bulge field. The intensity scaling is set so that H$\beta$ occupies the full free intensity scale in all cases, and the entire useful wavelength range is shown

  
\begin{figure}
\includegraphics [angle=90,width=13cm,height=9cm]{ap6to10.eps}\end{figure} Figure 8: The B600 spectra for PN28, PN23, PN12, PN10, and PN1 in the M 31 bulge field. The intensity and wavelength scales are as in Fig. 7. Note that the background subtraction is poorer for PN12 than is normally the case

  
\begin{figure}
\includegraphics [angle=90,width=13cm,height=9cm]{ap11to15.eps}\end{figure} Figure 9: The B600 spectra for PN3, PN38, PN36, PN53, and PN52 in the M 31 bulge field. The intensity and wavelength scales are as in Fig. 7

  
\begin{figure}
\includegraphics [angle=90,width=13cm,height=9cm]{ap16to20.eps}\end{figure} Figure 10: The B600 spectra for PN42, PN45, PN43, PN48, and PN95 in the M 31 bulge field. The intensity and wavelength scales are as in Fig. 7

  
\begin{figure}
\includegraphics [angle=90,width=13cm,height=9cm]{ap21to25.eps}\end{figure} Figure 11: The B600 spectra for PN47, PN408, PN93, PN92, and PN91 in the M 31 bulge field. The intensity and wavelength scales are as in Fig. 7. Note that the signal-to-noise is poor for the very faint object PN408

  
\begin{figure}
\includegraphics [angle=90,width=13cm,height=9cm]{ap26to28.eps}\end{figure} Figure 12: The B600 spectra for PN97, PN387, and PN380 in the M 31 bulge field. The intensity scaling is set so that H$\beta$occupies the full free intensity scale in all cases, and the entire useful wavelength range is shown. For PN97 and PN380, though, the wavelength range extends to 6600 Å

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