1 Introduction

With a plethora of emission lines, from the far ultraviolet all the way to the far infrared, H II regions are a very rich source of spectral information. The huge amount of information contained in the emission lines makes these regions ideally suited for investigating the properties of the local gas-phase interstellar medium (ISM). In principle, it is possible to deduce many of the important characteristics of an H II region from its spectrum.

Among such characteristics are the temperature and density of the electron gas, which can be derived from taking ratios of line fluxes. Well known examples of these are the often employed optical [O III] 5007 Å/4363 Å and [S II] 6717 Å/6730 Å line ratios for the electron temperature and density, respectively. Ratios from spectral lines emitted by different ionic species of the same element constrain the ionization structure inside the H II region. The ionization structure throughout the region reflects the hardness and density of the local radiation field which is a signature of the number and the mass of the major ionizing sources in the H II region. Given a set of electron temperatures and densities, the chemical composition of the gas can be derived from the spectral lines. The abundances thus determined provide clues for the local nucleosynthetic and star-formation history.

The analysis of an H II region in terms of the physical structure, stellar content and chemical composition is, however, not without difficulties. The derivation of the physical structure is often limited to the electron temperature and density as found with the "classical'' ratios of lines from O⁺⁺ and S⁺, but by limiting the analysis to these lines, one is actually probing only a part of the total emitting volume of the H II region. This is the result of the stratified nature of the ion distribution within the region. The electron properties found are taken to be representative of the entire nebula, but the existence of density variations and temperature gradients throughout the nebula is not ruled out (the t² problem, Peimbert 1967). Gradients and fluctuations in the temperature structure are immediately reflected in the abundances one derives. The abundances are underestimated because of the tendency of temperature fluctuations to bias the temperatures found towards higher values.

The accurate determination of elemental abundances is also hampered by the limited number of ionization stages one can observe for a particular element. For many astrophysically important elements, one often has to make corrections for unseen ionization stages using Ionization Correction Factors (ICFs) in order to derive the total elemental abundance. Many prescriptions for these ICFs have appeared over the last thirty years, either based on coincidences in ionization potentials of various ions (Peimbert & Costero 1969) or on photoionization models (e.g. Grandi & Hawley 1978; Mathis & Rosa 1991). The most difficult element in this respect proved to be sulfur (Natta et al. 1980; Garnett 1989).

Many different diagnostics have been proposed for constraining the effective temperature of the ionizing source of an H II region. Helium recombination line diagnostics like He I 5875/H$\beta$ are robust but insensitive for T_* higher than 39 000 K (Kennicutt et al. 2000). Single-element forbidden-line diagnostics, like [O III] 5007/H$\beta$ and [Ne III] 3869/H$\beta$, as well as composite diagnostics, like the $\eta{\arcmin}$ parameter (Vilchez & Pagel 1988), are also often used. These diagnostics, however, can be sensitive to abundance effects and variations in ionization parameter, especially in spatially resolved nebulae.

 

 

Table 1: The objects contained in the sample. Given are the coordinates of the ISO SWS and LWS pointings. The designations, except for 30 Doradus, are those of Henize (1956).

Object	SWS Coordinates (J2000.0)			LWS Coordinates (J2000.0)
	RA (h  m  s)	$\delta$ ( $^o\ ~\arcmin\ ~\arcsec$)	Date	RA (h  m  s)	$\delta$ ( $^o\ ~\arcmin\ ~\arcsec$)	Date
N66	0  59  03.7	-72  10  39.9	01-Apr.-1996	0  59  08.0	-72  10  26	01-Apr.-1996
N81	1  09  13.6	-73  11  41.1	14-May-1996	1  09  09.1	-73  12  07	22-Apr.-1996
N88A	1  24  08.1	-73  09  02.5	14-May-1996	1  24  09.4	-73  09  16	22-Apr.-1996
N79A	4  51  47.4	-69  23  47.6	10-Sep.-1996	4  51  49.3	-69  24  36	29-Apr.-1996
N4A	4  52  08.4	-66  55  23.4	29-Apr.-1996	4  52  06.4	-66  55  27	29-Apr.-1996
N83B	4  54  25.2	-69  10  59.8	14-May-1996	4  54  25.2	-69  10  59	01-Nov.-1997
N11A	4  57  16.2	-66  23  18.3	22-Apr.-1996	4  56  48.3	-66  24  11	17-Jul.-1997
N159-5	5  40  02.4	-69  44  33.4	10-Jul.-1996	5  40  04.2	-69  44  42	06-May-1996
N157B	5  37  51.6	-69  10  23.1	14-May-1996	5  37  51.8	-69  10  22	29-Apr.-1996
30 Dor#1	5  38  33.5	-69  06  27.1	05-Dec.-1997	5  38  33.5	-69  06  27	29-Apr.-1996
30 Dor#2	5  38  35.5	-69  05  41.2	13-Apr.-1996	5  38  43.8	-69  05  17	29-Apr.-1996
30 Dor#3	5  38  46.0	-69  05  07.9	06-May-1996	5  38  45.8	-69  05  11	29-Apr.-1996
30 Dor#4	5  38  54.2	-69  05  15.3	05-Aug.-1997	5  38  54.2	-69  05  11	29-Apr.-1996
N160A1	5  39  43.3	-69  38  51.4	10-Jul.-1996	5  39  44.2	-69  38  42	06-May-1996
N160A2	5  39  46.1	-69  38  36.6	10-Jul.-1996	5  39  44.2	-69  38  42	06-May-1996

Figure 1: The Magellanic Clouds in H$\alpha$. The LMC is left and the SMC is right. Indicated are the H II regions in the sample. North is up and East is left. The images were obtained with the Parking Lot Camera (Kennicutt et al. 1995; Bothun & Thompson 1988).

To address the problems described above, one needs to use as many spectral lines as possible. By extending the data base from the optical into other spectral regions, such as the infrared, many of the problems can at least be alleviated. The inclusion of more spectral lines makes it possible to derive electron temperatures and densities for more different ions, giving a better coverage of the temperature profile of the nebula. A larger spectral coverage gives us access to more different, otherwise unobservable, ionic species. This makes a direct calculation of the elemental abundance possible without having to resort to ICFs. It also puts more constraints on the ionization structure of the nebula and, therefore, on the nature of the ionizing sources.

For this purpose a large set of new optical and infrared spectroscopic data has been obtained of several bright H II regions in the Magellanic Clouds. The infrared spectra were taken with the Short Wavelength Spectrometer (SWS) and Long Wavelength Spectrometer (LWS) on board the Infrared Space Observatory (ISO, Kessler et al. 1996), as part of a comprehensive ISO Guaranteed Time Program on H II regions in Local Group Galaxies.

The Magellanic Clouds form excellent targets for such studies. With their small distance of 55 kpc (LMC, e.g. Feast 1999) and 65 kpc (SMC, e.g. Kovács 2000), star formation and star-forming regions can be studied in considerable detail, down to the scale of individual stars. Given their high elevation above the Galactic plane, the problem of extinction towards the Magellanic Clouds is less acute than for Galactic H II regions. The study of Galactic H II regions is often limited to the local neighbourhood because of the extinction in the Galactic plane, whereas in the Magellanic Clouds H II regions can be studied on a galaxy wide scale.

In this first of a series of papers, the data are presented. The data set comprises optical driftscan spectra, as well as infrared spectra obtained with the Infrared Space Observatory (ISO). The layout of the paper is as follows. Section 2 gives an overview of the objects in the sample and their environment, Sect. 3 gives a detailed account of the reduction of the optical and infrared spectra, and Sect. 4 discusses the quality and reliability of the spectral data presented. The analysis will be deferred to subsequent papers.