H e a v y E l e m e n t A b u n d a n c e s in I o n ized N e b u l a e

loannis Tsamis

A thesis submitted in partial fulfilment of the requirements for the degree of Doctor of Philosophy at the University of London

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Optical recombination-line (ORL) abundances for carbon, nitrogen, and oxygen, relative to hydrogen, are presented for 12 Galactic planetary nebulae (PN) and 3 Magellanic Cloud PN (LMC N141, LMC N66 & SMC N87) and an extensive comparison with the corresponding abundances derived from UV {lUE), optical and IR {IRAS, ISO) collisionally-excited lines (CELs) is performed. Two extreme nebulae, NGC2022 and LMC N66, are exposed that exhibit very discordant - ORL versus CEL - abundances; by a factor of > 10 for the O^"*" ion. The Type I PN NGC2440 is revealed to have discrepancy factors similar to those of the pre­ viously studied NGC7009 and M 2-36. Along with previous results it is shown that ORL/CEL ionic abundance enhancements differ among nebulae, spanning a range from ~ 1.5 to > 20. There are indications that the C /0, N /0 elemental ra­ tios derived from ORLs may be larger that the corresponding CEL ratios, hinting at an origin for the two types of line from nebular material of dissimilar chemical history. The abundance enhancements of doubly ionized C, N and O are positively correlated with the difference (AT) between the Baimer recombination continuum and [O III] forbidden-line temperatures, suggesting that temperature variations, real or induced, are partly to blame for the discrepancies. The relative unifor­ mity, however, of the overabundance patterns and lack of consistent correlation with CEL excitation energies point away from Peimbert-type, simple temperature fluctuations as the cause of the problem. Recent results by Garnett & Dinerstein (2000) are confirmed, showing that the abundance discrepancies are: i) anticor­ related with the intrinsic nebular surface brightness, and ii) positively correlated with PN absolute radii, i.e. young, bright nebulae display less abundance discrep­ ancies than older, more extended ones. It is further shown however, that very sim­ ilar correlations exist between AT and the nebular radii and surface brightnesses, suggesting that fainter, more extended objects are more likely to have temperature discrepancies than brighter, smaller ones. These findings strongly indicate that the long-standing ORL/CEL abundance discrepancy problem is associated with the evolution of PN. Optical spectra of five H II regions, two within the Galaxy (M 17 and NGC 3576) and three in the Magellanic Clouds (30 Doradus, LMC N llB and SMC N66) were analyzed. The ORL/CEL discrepancy factors for the 0^+ ion are found for the first time to be in the range of 2-5, thus placing these objects in the abundance discrepancy regime of PN. For the first time also ORL/CEL discrepancies are doc­ umented for extragalactic Hll regions. Accurate, temperature-insensitive ORL C2+ /o 2+ ratios are derived for all five nebulae, showing remarkable agreement within each galactic system. A cknowledgements

I’d like to thank Mike Barlow for taking me on this project, ensuring that it was not left unfinished by sustaining me through my fourth year and then declaring that my thesis is a joy to read (I leave that to the reader); also for making possible the trip to ESO’s mountaintop on La Silla, Chile, where I checked for myself the round Earth hypothesis by observing the southern sky.. .only I forgot to put that in my thesis. He is also thanked for advice on the analysis of H II regions spectra and for other bits of astronomical information. Mike in short is a European student’s ideal supervisor, mainly because he won’t let you give up, once your studentship evaporates right on time when you thought you’ve got a thesis at last. I would therefore strongly recommend Mike in case he wanted to become Departmental chairman, provost or European Commissioner (since he’s not a euro-sceptic). Next I’d like to thank Xiaowei Liu and Pete Storey; the former mainly for the loan of his customized MIDAS line-fitting routine that made my life a lot simpler and the latter for letting me use his nebular optimization code that made Chapter 3 possible. Paul Crowther is also thanked for advice on the spectral characteristics of OB and WR stars. I hope all the people in Physics & Astronomy who did their PhDs at roughly the same time when I did mine, go on to do something even more rewarding after it’s all over; and that the Perren Fund continues as a financial support for non-UK students. Silvia wants to be acknowledged too, so let me say it here that I have no pity for her would be thesis advisor in Guildford whose project she dumped (after having pocketed his 1st semester grants in a typical Italian way) and so she came to cool atoms in UCL and to meet me. These last 2.5 years in London wouldn’t have been so cool - and hot! - otherwise. Finally, and most importantly, I thank my family in Thessaloniki, Greece for their love and support through all this experience. C o n t e n t s

List of Figures 8

List of Tables 10

1 Introduction 13 1.1 Physical processes in photoionized nebulae ...... 14 1.1.1 Photoionization ...... 14 1.1.2 Radiative recombination and line emission ...... 17 1.1.3 Dielectronic recombination ...... 19 1.1.4 Collisionally excited line emission ...... 19 1.2 Determination of electron densities and temperatures in nebulae . 22 1.2.1 Ne, Te determinations from CELs ...... 22 1.2.2 Ne, Te determinations from HI Baimer lines and optical con­ tinuum measurements...... 26 1.3 Determination of elemental abundances in nebulae ...... 28 1.3.1 Derivation from observations of C E L s ...... 28 1.3.2 Derivation from observations of O R L s ...... 29 1.3.2.1 H e liu m ...... 29 1.3.2.2 Heavier elements: Carbon, Nitrogen, Oxygen, and N e o n ...... 30 CONTENTS

1.3.3 Temperature and density variations in nebulae: implications on abundance determinations ...... 31 1.4 Interstellar extinction ...... 35 1.4.1 The derivation of c (H /3 ) ...... 38 1.5 Aims of this thesis and outline ...... 38

2 Elemental Abundances in Galactic and Magellanic Cloud Plane­ tary Nebulae 40 2.1 Optical Observations ...... 40 2.1.1 D ata re d u c tio n ...... 43 2.2 The lUE observations ...... 47 2.3 D ata a n a ly sis ...... 48 2.4 Nebular A n aly sis ...... 52 2.4.1 Reddening ...... 52 2.4.2 Nebular Diagnostics ...... 54 2.4.2.1 Electron D en sities ...... 55 2.4.2.2 Electron Temperatures ...... 57

2.4.2.3 Recombination excitation of the N II and O II au­ roral lines ...... 58 2.4.2.4 Baimer discontinuity electron temperatures .... 63 2.5 Elemental abundances from CELs ...... 64 2.6 Ionic Abundances from O R L s ...... 71 2.6.1 Helium ...... 72 2.6.2 C2+/H+, C3+/H+ and C^+/H+ ...... 72 2.6.3 N2+/H+, N^+/H+ and /E + ...... 76 2.6.4 0 2 + /H + ...... 78 2.6.5 Total C, N and O Abundances from ORLs ...... 93 2.6.6 Total elemental abundances ...... 103 2.6.7 Comparison of ORL and CEL abundances ...... 104 2.6.7.1 Abundance grids and ionic discrepancy ratios . . . 104 CONTENTS 6

2.6.7.2 C /0 and N/G elemental r a tio s ...... 117 2.7 D iscussion ...... 121 2.7.1 Correlations ...... 121 2.7.1.1 Discrepancy factors versus temperature gradients 121 2.7.1.2 Discrepancy factors versus PN intrinsic surface bright­ ness ...... 122 2.7.1.3 Discrepancy factors versus PN absolute radii . . . 123 2.7.1.4 Temperature gradients versus PN radii and surface brightness ...... 124 2.7.2 Temperature fluctuations ...... 131 2.7.3 Density inhomogeneities ...... 134

3 Empirical Composite Models of the Planetary Nebula NGC 5882138 3.1 Introduction ...... 138 3.2 NGC 5882: overview of nebular p roperties ...... 139 3.3 Results and d iscussion ...... 143

4 Elemental Abundances in Galactic and Magellanic Cloud H ll Regions 153 4.1 Observations ...... 153 4.2 D ata re d u c tio n ...... 156 4.3 Nebular analysis ...... 157 4.3.1 Reddening correction ...... 157 4.3.2 Electron temperatures and densities ...... 157 4.4 Ionic and total elemental abundances from C E L s ...... 160 4.5 Ionic abundances from O R L s ...... 160 4.5.1 C2+/H+ andN2+/H+ ...... 160 4.5.2 02+/H+ ...... 162 4.5.2.1 Relative intensities of fine-structure V 1 lines . . . 165 4.5.2.2 Continuum observations and scattered light .... 167 CO NTENTS 7

4.5.2.3 Uncertainty ...... 179 4.6 D iscussion ...... 180 4.6.1 Abundance discrepancies and elemental ratios ...... 180 4.6.2 Causes of the discrepancies ...... 183

5 Conclusions 190

A Emission line fluxes: the catalogue 194 A.l Galactic Planetary Nebulae ...... 195 A.2 Magellanic Cloud Planetary Nebulae ...... 250

A.3 Galactic H II Regions ...... 256 A.4 Magellanic Cloud H ll Regions ...... 263

Bibliography 274 L ist of F ig u r es

1.1 Energy-level diagram for O'*" and ...... 24 1.2 Energy-level diagram for N"*" and ...... 25 1.3 The interstellar extinction curve of the Small Magellanic Cloud. . . 37

2.1 The [C lll]-|-[S ll]-|-0 ll A4070 blend in the spectrum of NGC 5882. 51 2.2 NGC 3242 &; NGC 5882; abundance comparison grids ...... 108 2.2 NGC 5315 & NGC 3918: abundance comparison .grids ...... 109 2.2 NGC 2022 & NGC 6818: abundance comparison grids ...... 110 2.2 N G C2440 & NG C 6302: abundance comparison .grids ...... I l l 2.2 IC4406 & My Cn 18: abundance comparison grids ...... 112 2.2 NGC 3132 & IC 4191: abundance comparison grids ...... 113 2.2 LMC N66 & LMC N141: abundance comparison grids ...... 114 2.2 SMC N87: abundance comparison g r id ...... 115 2.3 k, N^+/0^'^ ORL ratios vs. the corresponding CEL ratios. 118

2.4 A(C2+/H+) & A(Q2+/H+) v s . a t ...... 126

2.5 A(N2+/H+) v s . a t ...... 127

2.6 A(Q2+/H+) & A(C2+/H+) v s . S(H (3) ...... 128

2.7 A(Q2+/H+) & A(C2+/H+) v s . T ...... 129 2.8 A T vs. 5(H/?) k A T vs. R ...... 130

3.1 The lUE, large-aperture SWP spectrum of NGC 5882 ...... 140 LIST OF FIGURES 9

3.2 The spectrum of the planetary nebula NGC 5882 from the blue atmospheric cutoff to 4050 Â ...... 142 3.3 The spectrum of NGC 5882 from 4000 to 4960 Â featuring the promi­ nent recombination lines from C, N, O and Ne ions ...... 151 3.4 The high-order Baimer lines of NGC 5882 as a density diagnostic. 152

4.1 Blue-violet spectrogram of the 06.5 Ib(f) -f WNL supergiant R 139 in 30 Dor ...... 187 4.2 Blue-violet spectrogram of the WN6(h) component of the multiple star R 140 in 30 Dor...... 188 4.3 Spectrogram of the blue supergiant Parker 3157 in LMC NllB,

showing the O II multiplet V 1 absorption lines ...... 188 4.4 The high-order Baimer lines of NGC 3576 as a density diagnostic. 189 L ist of T a bles

2.1 Journal of ESO 1.52m observations ...... 44 2.2 Journal of NTT 3.5m observations ...... 45 2.3 References for atomic data ...... 56 2.4 Electron densities of Planetary Nebulae ...... 58 2.5 Electron temperatures of Planetary Nebulae ...... 59 2.6 The [N II] and [O II] electron temperatures after correcting for the effects of recombination excitation ...... 63 2.7 He & CEL heavy-element abundances in galactic P N ...... 65 2.7 He & CEL heavy-element abundances in galactic PN - cont...... 66 2.8 He & CEL heavy-element abundances in Magellanic Cloud PN . . 67

2.9 High-excitation C II recombination lines ...... 75 2.10 (a.) Recombination line N^'^'/H"'" ...abundances in PN ...... 77 2.10 (b.) Recombination line N^'^/H'^ ...abundances in PN ...... 78 2.11 (a.) Recombination line abundances of ..PN ...... 79 2.11 (b.) Recombination line 0^+/H+ abundances of PN ...... 81 2.11 (c.) Recombination line abundances of PN ...... 85 2.12 Comparison of the observed and predicted relative intensities of

planetary nebula G II lines...... 86 2.13 Carbon abundances from optical recombination lines ...... 95 2.14 Nitrogen abundances from recombination lines in Galactic PN. . . 98

10 LIST OF TABLES 11

2.15 Oxygen abundances from recombination lines in Galactic PN. . . . 102 2.16 Recombination line Carbon and Oxygen in Magellanic Cloud PN. . 103 2.17 Elemental abundances in Planetary Nebulae, in units where log N(H.) = 12.0...... 104 2.18 Ionic ORL/CEL abundance discrepancies in PN ...... 116 2.19 Carbon to Oxygen PN abundance ratios from CELs and ORLs. . . 119 2.20 Nitrogen to Oxygen PN abundance ratios from CELs and ORLs. . 120 2.21 Total H/3 intensities, intrinsic surface brightnesses, angular and ab­ solute radii of PN ...... 125

3.1 Plasma diagnostics of NCC 5882 ...... 141 3.2 Ionic abundances of NCC 5882 derived from ORLs and CELs. . . . 141 3.3 NCC 5882: parameters of empirical models ...... 144 3.4 Comparison of observed line intensities with those from empirical models ...... 145

4.1 Journal of H II regions observations ...... 154

4.2 Plasma diagnostics of H II regions ...... 159 4.3 He and CEL heavy-element abundances in galactic and Magellanic

Cloud H II regions ...... 161

4.4 Ionic carbon abundances from ORLs in H II regions ...... 162

4.5 Recombination line N^"^/H+ abundances in H II regions ...... 163

4.6 Comparison of the observed and predicted relative intensities of H II

region O II lines...... 163

4.7 Observed continuum emission and dust-scattered light in H II regions. 169 4.8 Equivalent widths of nebular emission lines and of stellar absorption

lines in the observed continuum of H II regions ...... 170 4.9 Recombination line 0^"*'/H"^ abundances in Hll regions ...... 172

4.10 Corrected O II intensities and ORL 0^'*'/H‘*' abundances for 30 Dor and LMC N llB ...... 179 LIST OF TABLES 12

4.11 Elemental abundance ratios in H II regions ...... 182

A .l N G C 2022 lin e list...... 195 A.2 NGC 2440 linelist...... 198 A.3 NGC 3132 linelist...... 201 A.4 NGC 3242 linelist...... 204 A.5 NGC 3918 linelist ...... 210 A.6 NGC 5315 linelist...... 215 A.7 NGC 5882 linelist ...... 219 A.8 NGC 6302 linelist ...... 224 A.9 NGC 6818 linelist ...... 230 A. 10 My Cn 18 linelist ...... 235 A .l l i e 4406 linelist...... 237 A.12 IC4191 linelist (fixed-slit observations) ...... 240 A.13 IC4191 linelist (scanning-slit observations) ...... 245 A.14 SMC N87 linelist ...... 250 A.15 LMC N66 linelist...... 252 A.16 LMC N141 linelist...... 254 A.17M17 linelist...... 256 A.18 NGC 3576 AAT linelist ...... 258 A.19 NGC 3576 ESO linelist ...... 261 A.20 SMC N66 linelist...... 263 A.21 LMC NllB linelist...... 266 A.22 30 Doradus linelist ...... 270 C h a p t e r 1

Introduction

The great preponderance of luminous matter in the universe is locked within stars. Only some ten per cent or less of this matter can be found free in the space between stars, in gaseous form. If found in the proximity of a strong energy source the gas becomes ionized and glows at characteristic wavelengths (from x-ray to radio!), observable from ground and space telescopes. The study of this light reveals much about the portion of the universe that is in this fleeting, transitory state. Photoionized nebulae^ are conspicuous, both in our own Galaxy and in others, mostly as diffuse clouds marking the birthplaces of massive stars. The stars’ intense light ionizes the neutral gas from which they spawned, which then forms an H u region^ i.e. an emission nebula, defined by the glowing light of recombining hydrogen. Hll regions then are young objects, associated with that part of the history of matter, when it changes from free-roaming and diffuse to bound and concentrated. Planetary nebulae on the other hand are related to old, medium-heavy (~ 1<

M q < ~ 8Mq) stars that undergo a ‘rapid’ transformation, expelling parts of their

^‘Nebula’ from the Latin for ‘cloud’. The Greek synonym - a bit older! - is the female-gender word ^NecpeXr]’] e.g. see Hesiod, Theogony, 745: ‘There stands the terrible home of murky Night wrapped in dark violet clouds. In front of it the son of lapetus stands ...’

13 1.1. Physical processes in photoionized nebulae 14

mass in a shell around them. When such a star evolves, it builds a carbon-oxygen core which becomes degenerate by virtue of its high density, while further core evolution is limited by extensive mass loss from the surface. At a point in the star’s evolution when its surface temperature reaches about 30 000 K, the expelled matter in the vicinity of the ageing star becomes ionized and glows according to the same basic laws that govern the emission in Hll regions.

Both H II regions and planetary^ nebulae are short-lived objects compared to the stars that they are associated with. Studying the emission from Hll regions, we gain insight into the composition of matter that is incorporated into stars, while observing planetary nebulae gives us information about the altered composition of star-stuff that is returned back into interstellar space and that one day will become part of the next stellar generation. I suppose it is this dynamic nature of nebulae, associated with change, that makes their study so appealing and so fruitful in deciphering the chemical evolution of stars and galaxies.

1.1 Physical processes in photoionized nebulae

A brief outline is presented here of important aspects in the physics of photoionized nebulae.

1.1.1 Photoionization

Let us consider the realistic case of a single, hot evolved star [e.g the nucleus of a Planetary Nebula (PN)] or a compact cluster of hot young stars (e.g. an

OB stellar association in the case of H II regions), surrounded by a gas cloud composed of ~ 90per cent hydrogen and ~ 10per cent helium (by number), along with a sprinkling of trace, heavier elements, such as C, N, O, Ne, etc. The stars emit copious amounts of UV photons; those photons with wavelengths below the

^They do not really have anything to do with planets; the term is purely historical. In the

early days of telescopic observations (200 years ago) the first sightings of these objects revealed

them to be visually similar to the greenish disks of the outer solar system planets. 1.1. Physical processes in photoionized nebulae 15

Lyman limit (912 A) have energies greater than 13.6 eV thus being able to ionize hydrogen in its ground state. If there is enough hydrogen to absorb all the UV photons, the now photoionized nebula is said to be ‘ionization bounded’. In this case the UV radiation does not penetrate beyond an ionization front, exterior to which is the remaining neutral gas cloud. Hll regions usually come under this definition, unlike PN which can be ionization-bounded in some directions, but not in others. In the opposite case we refer to ‘matter bounded’ (or ‘density bounded’) nebulae, in which the radiation runs out of gas to ionize. Relating to such situations Baker & Menzel (1938) defined the so called Cases A and B: under Case A the optical depths at the frequencies of all line and con­ tinuum radiation are small and all photons escape (completely matter-bounded or optically thin nebula) ; under Case B all Lyman photons—both line and continuum radiation arising from recombinations to the n = 2 hydrogen level—are absorbed ‘on the spot’, that is, the optical depths at those frequencies are large (completely ionization-bounded or optically thick nebula). On the contrary, radiation belong­ ing to the Baimer and higher series has small optical depth and escapes the nebula. The optical depth, represents the number of mean free paths of radiation of frequency n travelling radially outwards before it escapes from the nebula. It is a dimensionless quantity defined as

Tu{t) = [ Njio{T)audr, (1.1) Jo where Ajjo is the number density of hydrogen atoms (in cm“ ^), is the ionization cross section of for photons with frequency z/ and r is the radial coordinate.

Optically thin nebulae are characterized by 1, while those for which Tj/ ^ 1 are optically thick. The ionization cross section drops off quite rapidly with energy, varying approximately as It follows that the most energetic photons travel further outwards in the gas than photons near the series limit, so that the radiation field at larger distances from the ionizing source is hardened. In reality. Case B is almost always a better approximation than Case A and under this assumption the hydrogen atom can be treated as though its ground 1.1. Physical processes in photoionized nebulae 16 state does not exist, since every emission of a Lyman photon is cancelled by an absorption from the ground state (apart from Lyman-a, n = 2^1, which is multiply scattered, diffusing slowly outwards through the nebula). In this way, stellar radiation of the Lyman continuum is degraded to photons of the higher hydrogen series, i.e. it is converted to Baimer, Paschen, Brackett etc. photons plus Lyman-o; photons. When ionization equilibrium holds, the emission rate of stellar ionizing photons per unit volume is equal to the number of recombinations; for a pure hydrogen nebula under Case B, this can be written as

/ (L„lhv)dv= fN(R+)NeaBdV, (1.2) J uo JV where Li, is the stellar luminosity and fV(H"'"), Nq are respectively the ion and elec­ tron number densities; a g is then the total recombination coefficient (in cm^ sec“ ^), representing the sum of rates of direct recombination to the second and higher lev­ els of the ion. The recombination coefficients have a rather weak dependence on temperature varying roughly as oc T~^. In the case of an Hll region the left part of Eq. 1.2 is proportional to the total number of stars embedded in it; Eq. 1.2 can be solved to yield the size of a spherically symmetric ionization-bounded neb­ ula, known as the ‘Stromgren sphere’. Typical Stromgren radii are about 0.2 pc for PN, while HII regions have a larger spread in size: the radius of the Orion neb­ ula is a few pc, but that of the giant nebula 30 Doradus in the Large Magellanic Cloud (LMC) is about 100 pc. The excess energy {hu — 13.6 eV) of the ionizing photons is transferred to the expelled photoelectrons in the form of kinetic energy; due to the rather large cross section for electron-electron collisions (~ 10~^^cm~^) this energy is distributed amongst the free electrons, which acquire a Maxwell-Boltzmann energy distribu­ tion. This process is the main source of heating in nebulae. The recombination cross sections on the other hand are much smaller, of the order of 10~^°cm“^, while the photoionization cross sections are larger than that, at ~ 10“ ^^ cm~^ (e.g. Osterbrock 1989). Thus any deviations from ionization equilibrium take time to wear off; this timescale is called the recombination time, r^, which is equal to 1.1. Physical processes in photoionized nebulae 17

(7Veas)“ ^. For typical values of og = 3 x 10“ ^^ cm^ sec~^ and TVg = 5000 cm~^, the recombination time is % 20 years. The heating rate via photoionization is proportional to the number densities of the ions involved. Since hydrogen is the most abundant element it provides the dominant contribution to the heating rate; and as mentioned above due to the hardening of the radiation field further out in the nebula, the mean energy (or equivalently the temperature) of liberated photoelectrons there is greater than those produced nearer to the star. A nebula without any radiation losses would heat up; then its temperature at a given time would be just higher than the temperature corresponding to the mean energy of freshly created photoelectrons (since the slowest electrons are pref­ erentially captured by ions). But, there are radiation losses in nebulae and hence cooling of the gas due to the loss of kinetic energy by the electrons following recom­ binations. Much more importantly, however, heat is being removed from nebulae chiefly by the emission of collisionally-excited lines (CELs; see Section 1.1.4) from

ions such as - via [O III] 52-, 88-/im, AA4959, 5007, O'*" - via [O ll] AA3726, 3729, N+ - via [N ll] AA6548, 6584 lines etc.^ The established thermal balance results at an electron temperature of about 10 000 K for a typical nebula.

1.1.2 Radiative recombination and line emission

This process is the free-bound recapture of an electron by an ion, followed by downward cascades leading ultimately to the ground level of the atom or the next lower stage of ionization. These transitions are accompanied by the emission of recombination lines right across the spectrum from UV to radio wavelengths; the strongest recombination lines in nebulae are those primarily of H I, e.g. the Baimer optical recombination lines (ORLs) like Ha A6563, H/3 A4861, H7 A4340 etc, but

also of He I (A4471, A5876, A6678) and He li (AA1640, 4686). The radiation field within H II regions is not strong enough to produce significant amounts of doubly

^The brackets indicate lines that are forbidden for electric dipole transitions; the Latin num­ bers following the element’s symbol indicate, in the standard astrophysical notation, the degree of ionization, i.e. T for neutral, ‘IF for singly ionized and so on; also A4959 stands for 4959 Â. 1.1. Physical processes in photoionized nebulae 18

ionized helium, so lines of He II are only present in PN spectra. The typical timescale between photoionizations is of the order of 10^ sec, while the transition probabilities following recombinations are about 10^-10^ sec~^. Thus all ions can be considered to be effectively in their ground states and that photoionizations occur from that level. The intrinsic intensity (i.e. corrected for interstellar extinction, in units of ergcm“^s“^) of a recombination line due to an z —^ j transition, received at the Earth, is given by

47T where the nebula is at a distance D, and is the number density of the recombining ion. The emissivity of the line can be written as

hr tx = lV(X+'+‘) % «eff(A, T ) y , (1.4)

and has units of ergcm~^s“ ^. Here CKgff(A, T) is the effective, line recombination coefficient, itself given by

aeff(A) = B(A)c,«g(%+:). (1.5)

It represents the probability for spontaneous radiative decay via line A (i.e. the branching ratio), B{X) = Aik, k

1.1.3 Dielectronic recombination

This process involves the radiationless capture of an electron by an ion which then jumps to the next lower stage of ionization having two electrons in excited states, followed by the emission of radiation as one of the electrons de-excites, resulting in a singly excited bound level; the latter can then decay to the ground level of the ion. Hence dielectronic recombination provides an additional mechanism for populating high-lying ionic levels. The effective dielectronic recombination coeffi­ cient is relatively insensitive to density, but has a strong temperature dependence which can be characterized roughly as oc e ~ ^ (e.g. Nussbaumer & Storey 1984). At sufficiently high electron temperatures dielectronic recombination to states of high principal quantum number becomes the dominant mechanism for recombination. The effect of dielectronic recombination has to be quantified since it can have an impact on abundance determinations in nebulae; for example, if unaccounted for, it can lead to deceptively high ionic abundances being determined by observers who register photons, but assume they originate from the radiative recombination process only. Thus in its general form the effective recombination coefficient for a given transition will be the sum of radiative and dielectronic parts. Effective radiative and/or dielectronic recombination coefficients for a variety of ions have been calculated by e.g. Nussbaumer & Storey (1984), Péquignot, Petit jean & Boisson (1991), Smits (1991), Storey (1994) and Davey et al. (2000).

1.1.4 Collisionally excited line emission

There are emission lines in nebulae created following electron-impact excitation of low-lying (a few eV from the ground state) ionic or atomic levels and subsequent spontaneous radiative decay. These are in most cases the brightest lines observed from both PN and Hll regions, championed by the green [O III] AA4959, 5007 lines of and others of O'*", N""", N^'*’, Ne^’*', S'*", Ar^"*" etc., extending from the UV, through the optical to the far-IR domain. As mentioned previously, such 1.1. Physical processes in photoionized nebulae 20 lines are the major sources of cooling in nebulae, in spite of the low abundance of heavy element ions, compared to that of H and He ions; the latter elements have atomic levels of much higher excitation energies, not easily populated by electron impacts and hence their recombination line emission is of minor importance to the nebular thermal balance. Collisionally excited lines in the optical regime arise from transitions whose up­ per terms are metastable. Such transitions violate the Laporte parity rule, accord­ ing to which a transition between two odd or two even atomic levels is forbidden - whereas the other type that obey this rule are called permitted or electric dipole transitions. These lines due to their very low transition probabilities (~ 10“^ to 10"^ sec"^) are not seen in the laboratory, where they are suppressed via collisional de-excitation. However, the prevailing electron densities in photoionized nebulae are so low (~ 10- 10^ cm“^) that most of the upward excitations are followed by radiative decay after a time of the order of seconds. CELs are also excited usually by transitions for which Al = 0, i.e. they do not involve an angular momentum change. In the UV spectral domain, we also find CELs arising from transitions that violate the A S = 0 selection rule, i.e. they involve a change of spin (or mul­ tiplicity 25 + 1); these are referred to as semi-forbidden or intercombination or intersystem lines and sometimes they do not violate the parity change rule. In nebular plasmas, the level populations of a given ion are determined by the equilibrium between radiative decay and collisions with electrons to and from a given excited level i and other levels j. Due to the rarefied state of nebulae the plasma is far from local thermodynamic equilibrium (LTE) so that the relative level populations in atoms can not be obtained from the Boltzmann formula (i.e. the nebula is not in LTE, such as e.g. a stellar interior); instead these are calculated on the basis of an assumed statistically steady state, whereby the principle of detailed balancing holds. In this, the individual rates that populate and de-populate a level are assumed to balance. For a multi-level ion the statistical equilibrium equations 1.1. Physical processes in photoionized nebulae 21

for each level are (Osterbrock 1989)

^ ^ A/j N q qji + ^ ] A/j Aji = ^ ^ Ni A/g Qij + ^ ^ A/^j , (1-6) jÿéi j>i jÿéi j

YiNj = N, j can be solved to yield the population of each level relative to the ground state. The left part of Eq. 1.6 gives the transitions that populate level z, while the right part gives transitions that depopulate level j. Aij are the radiative transition probabilities for i j (to all lower levels) and qij are the collisional excitation/de­ excitation rates for level i. For i > j, the collisional de-excitation rate is

8.63 X 10-® Q{j, i)

and the collisional excitation rate from j i is given by

Qji = — Qij , ( 1 - 8 ) 9j where Q is the relevant collision strength (a dimensionless constant which is aver­ aged over the Maxwellian velocity distribution and is a slowly varying function of Te), the g's are the respective statistical weights of the levels^, Eij is the energy of level i above j and Tg is the electron temperature. H’s for e.g. O III are tabu­ lated by Aggarwal (1983), while A^’s for the same ion are given by Nussbaumer & Storey (1981). In this way the emissivity of z —^ j transitions can be written as

eij = N(X+”')mAijEij ( 1.9)

(in units of ergcm-®sec-^), where N{X'^'^) is the ionic abundance and ni is the fractional level population. In the low-density limit (collisional de-excitation negligible), a photon is emitted for every collisional excitation, so Eq. 1.9 becomes

eij = N{X+’^)N,qjiEij. (1.10)

under LS'-coupling, the statistical weight of a level g = 2J + 1, where J is the total angular momentum L -t- 5 of the term. 1.2. Determination of electron densities and temperatures in nebulae 22

At higher densities the fractional level population in Eq. 1.9 ceases to be pro­ portional to the collisional excitation rate, but is rather given (in a two-level approximation) by

' 9 jl + A,j/(qijN,y ( ^ ^ Thus, above a certain critical density (for the two-level case),

Ncr = Aij f Qij.) (1.12) the fractional population of the excited state ceases to increase with Ne and the CEL is collisionally suppressed. The critical density of a level is simply the density at which the rate of excitations that populate the level equals the rate of collisional de-excitations that depopulate it. Collisional suppression has to be allowed for when the nebular electron density approaches within an order of magnitude of

N ’cr-

1.2 Determination of electron densities and tempera­ tures in nebulae

1.2.1 Ne, Te determinations from CELs

From the discussion in the previous section it follows that the emissivity ratio, and hence the observed intensity ratio, of two collisionally excited ionic levels physically depends on how the level populations are influenced by collisional processes that populate and depopulate the two levels. We can see that when collisional de­ excitation rates for the two levels are very different the line ratio is sensitive to the nebular electron density; conversely, when the collisional excitation rates for the two levels are very different, the line ratio is sensitive to the electron temperature. Let us consider the intensity ratio of two CELs whose upper levels have very similar excitation energies, but quite different lifetimes (or equivalently transition probabilities), and which decay to a common lower level (see Fig. 1.1). Then from Eqs. 1.9 and 1.11, it follows that the emissivity/ intensity ratio of these two lines is 1.2. Determination of electron densities and temperatures in nebulae 23

a function of since the temperature dependencies largely cancel out; therefore the line intensity ratio can be used to determine the nebular electron density. The useful range of emissivity ratios across which N q determination is possible lies between the low (A^e*CA^cr) and high (A^g^-^cr) limits. At the low density limit, Eqs. 1.7, 1.8 and 1.10 yield line emissivity ratios that are equivalent to the ratio of collision strengths. At the other density extreme, Eqs. 1.9, 1.11 and 1.12 yield an intensity ratio fixed by the ratio of statistical weights and transition probabilities. At optical wavelengths, the doublet transition in the ground term of np^ ions gives rise to lines such as [O li] AA3726, 3729, [S ll] AA6716, 6730, [Ar iv] AA4711, 4740 and [Cl ill] AA5517, 5537 that are extensively used as density indicators in nebulae (see Fig. 1.1). The [O II] and [S II] lines become collisionally suppressed at relatively low densities (A^cr of upper levels ~ 1-4 x 10^ cm~^), while the [Ar iv] and [Cl III] lines continue to be useful diagnostics up to ~ 10^ and

~ 3 X 10'^ cm“^ respectively. Medium to high spectral resolution (FWHM ~ 1 Â) is needed to resolve the [O ll] doublet and to deblend the [Ar iv] A4711 line from a He I A4713 line.

In the UV domain intercombination lines such as e.g. C III] AA1907, 1909 and

N III] AA1748.7, 1752.2 etc. exist that are useful diagnostics, particularly as probes of high density plasma, since they involve transitions with high Agr upper levels 10^-10^^ cm“^; see Rubin 1989 and review by Dufour 1995). They require however the use of high spectral resolution observations for their analysis, e.g. high dispersion lUE or HST spectrograms. In the far-IR the [O III] 52/zm/88//m line ratio has also been used (see review by Dinerstein 1995, Liu et al. 2001); since these lines arise from low critical density upper levels, they are good indicators for low-density regions, but saturate at relatively low densities (~ 3500 cm“ ^ for the [O III] lines, but quite higher for e.g. the [S III] 18.7, 33.5 p,m lines). Regarding the electron temperature this is usually determined in nebular stud­ ies from intensity ratios of CELs arising from widely separated excitation energy levels. In Fig. 1.2 the energy level diagrams of N+ and ions are shown. 1.2. Determination of electron densities and temperatures in nebulae 24

o n

P3/2.1/2

SII X7330 X7319 P i / 23/2

D s/2,3/2 M076 X4068

D 3/2.5/2 À3729

X3726 X6731 X6716

S 3/2 S3/2

Figure 1.1: Energy-level diagram for 0'*' and S'*" showing the [O ll] and [S ll] lines used for Ne and Te determinations; see text for details.

In the simplest situation, which occurs in the low density limit the ion can be considered to have three levels (the ground term treated as one level); for e.g. we assume that for the level the important rates are collisional excitation from the and radiative decay downward in the nebular AA4959, 5007 lines; it further holds that the radiative transition rate from the highest level to the via the auroral A4363 line is equal to the upward excitation rate from the ground term.^ Then from Eqs. 1.7, 1.8 and 1.10 it follows that

/(A4959 + A5007) AE/kTe (1.13) /(A4363) where AE is the energy difference between the and levels. Prom this proportionality it is easy to see that the temperature of the emitting region may

be obtained from a measurement of the intensity ratios of nebular to auroral [O III] or [N 11] transitions. The nebular to auroral line ratios from both these ions are

^The term ‘nebular’ refers to the lines on account of their original identification in

PN spectra; the term ‘auroral’ refers to ^ S - ^ D transitions since the corresponding [O l] A5577

line is present in the Earth’s aurorae (and is also ubiquitous in night sky spectra). 1.2. Determination of electron densities and temperatures in nebulae 25

OIII

S o

N i l

S o X4363

X5754 D 2

D2

A.4959 X6548 X6584 X5007

1 52-|xm ^ P o , l ,2 Tiiijim P o , U

Figure 1.2: Energy-level diagram for N'*' and showing the [N ii] and [O ill] auroral and nebular lines used for Tg determinations; the fine-structure splitting of the ground terms has been exaggerated for clarity; the far-IR [O ill] 52- and 88-//m lines are useful density and abundance diagnostics.

valid temperature diagnostics for Nq< 10^cm“^; above that density the nebular lines become collisionally suppressed (but not the auroral lines which saturate at > 10^cm“^), something which can lead deceptively high Te’s to be derived in the presence of density variations within the nebular volumes. This important matter will be dealt with later in this study. Apart from the aforementioned diagnostics, intensity ratios of [O ll] and [S ll] lines arising from transitions between the and terms (see Fig. 1.1) are potential temperature diagnostics and have been used in this work to derive Te’s for a number of PN and Hll regions; e.g. the [O ll] (A3726 + A3729)/(A7319

+ A7330), [S i i ] (A4068 + A4076)/(A6716 + A6730) ratios. However, these di­ agnostics may be of limited usefulness since the lines involved are often affected by processes other than simple collisional excitation/de-excitation; contributions from recombination excitation have to be accounted for in some cases. In Chapter 2 the related problem will be discussed. 1.2. Determination of electron densities and temperatures in nebulae 26

The forbidden-line ratios determine the nebular temperature in the region where the lines themselves are emitted. It should thus be noted that due to the exponential sensitivity of the emissivities of CELs to the electron temperature (cf.

Eqs. 1.9 and 1.11), the mean derived Tq tends to be weighted more towards re­ gions of higher-than-average electron temperature, in the case of a non-isothermal nebula.

1.2.2 Ne, Te determinations from H I Balmer lines and optical con­ tinuum measurements

It is true that in the limit of low density, the intensity ratio of any two HI Balmer lines is insensitive to Ne (cf. Eq. 1.3), since the only processes that need to be considered in the formation of line emission are electron captures and subsequent downward-radiative transitions. In the case of increasing Ne however, electron and/or proton impacts tend to alter the population of hydrogen atomic states and this effect becomes also increasingly important for levels of high principal quantum number n. The larger collisional cross sections are those for population-altering transitions for which n is fixed and L changes by ±1, followed by those that both n and L change by ±1 (e.g. Osterbrock 1989). With collisions taken into account, the emissivity ratios of high-order Hi Balmer lines (n —^ 2, n > 10) to that of H/3 (n = 4 —> 2) become sensitive to Ne and thus the observed intensities of these lines relative to H/3 provide a valuable density diagnostic. In particular the strength of high-order Balmer lines relative to H/? tends to rise with increasing N q. This effect has been used in this work to investigate the possible existence of high-density ionized material in the case of a few nebulae for which our observations were of sufficient spectral resolu­ tion to allow the accurate recording of Hi Balmer lines with upper level n up to 24. In previous sections the determination of nebular electron temperatures was discussed which under the standard practice makes use of intensity ratios of CELs. 1.2. Determination of electron densities and temperatures in nebulae 27

It was also mentioned that the emissivities of recombination lines have a weak,

similar dependence to Tq\ from this it follows that the intensity ratio of two recom­ bination lines can not be used as a temperature diagnostic due to the very similar Te sensitivity of the various recombination coefficients. On the other hand, the strength of the Hi Balmer continuum at the series limit (~3646Â) decreases ap­ proximately as DC and therefore its intensity ratio to Balmer recombination lines, whose temperature dependencies are roughly oc is a useful measure of

Te. The nebular continuum near the Balmer limit consists of the HI recombination

and two-photon continua, as well as the He I (and He II for PN) recombination continua. The two-photon (27) continuum results from the decay of the highly metastable ^5*1/2 level of the hydrogen atom to the ground state via a virtual p state lying between the n — 1 and n = 2 levels; in the process two photons are emitted with a total frequency equal to that of Lyman-a (n = 2 —>• 1), but spread in a continuum energy distribution longward of Lyo (e.g. Lang 1974). In order to determine Tg one can therefore use the ratio,

4(A3646-)-/c(A3646+) ^ ^ I(H I) (in units of Â~^), which is a function of the nebular electron temperature. The continuum intensity difference corresponds to the Hi Balmer discontinuity, the so called ‘Balmer jump’ (BJ), which results from electron captures directly to the n = 2 level of Hi; the recombination line intensity, 7(Hi), commonly used is that of H/3, although better accuracy can be achieved by using a Balmer line closer in wavelength to A3646, such as e.g. H ll A3770, thus minimizing errors due to potential uncertainties in the correction for interstellar extinction or the flux calibration of the observed spectra. A notable fact is that in non-isothermal nebulae, the Balmer jump temperature tends to be weighted more towards nebular emitting regions of lower-than-average temperature, since the temperature dependence of Balmer recombination line and continuum emission varies roughly as oc with a in the vicinity of 1. 1.3. Determination of elemental abundances in nebulae 28

In this study Balmer jump temperatures obtained thus have been measured

for a number of PN and H I I regions.

1.3 Determination of elemental abundances in nebulae

1.3.1 Derivation from observations of CELs

Traditionally the abundances of heavy elements in nebulae have been determined relative to the most abundant element hydrogen with the use of CELs. From Eqs. 1.3, 1.4, 1.9, the measured intensity ratio of a CEL to that of a Hi Balmer line (the most commonly used line is H/3) can yield the ionic abundance of an element relative to that of H'*' in the following manner:

iV(X'"+) ^ /(A) AT(H+) AijmEij /(H/3)’ ^ ’ with the various terms defined as previously. An important prerequisite however for using CELs in order to derive ionic abundances, is the accurate knowledge of the electron temperature in the emitting region where ion is expected to exist, since by virtue of the population factor Ui in the above equation, the deduced abundance has an exponential dependence on Te (cf. Eq. 1.11). Hence in the standard practice, the ionic abundances of CEL emitting ions (e.g.

O"^, N+, C^"*" etc.), relative to H+, can be determined once the (Te, N q) of the nebula has been measured with one of the methods described in Section 1.2. In order to subsequently derive the total abundance of a given element relative to hydrogen, one adds up the relative abundances of observed ionization stages; the more ionization stages are observed, the better is the achieved accuracy. Whenever important ionization stages are not covered by the spectral observations, then the empirical approach makes use of ionization correction factors (ICFs) and total abundances may be found as: 1.3. Determination of elemental abundances in nebulae 29

A standard ICF scheme is that of Barker (1980), which stems from the similar­ ity of ionization potentials amongst various ions. In this work the ICF scheme of Kingsburgh & Barlow (1994; KB94 hereafter) was used; it is based on results from detailed photoionization modelling of PN, which apart from similarities be­ tween ionization potentials, takes also into account the effects of charge exchange, dielectronic recombination etc.

1.3.2 Derivation from observations of ORLs

It also possible to determine ionic abundances with the use of optical recombina­ tion lines. One can derive the abundance of a given ion relative to H+ from the ratio of observed (corrected for interstellar extinction) line intensities, by means of Eq. 1.3: jV(X+m+i) _ A aeff(H/5) 7(X+’", A) iV(H+) 4861 aeft(X+”>) /(H/3) ’ where /(A) is the intensity of an emission line from the recombined ion

1.3.2.1 Helium

In this manner, the abundance of helium is determined from observations of strong, optical He I recombination lines, e.g. the A4471, A5876, A6678 lines. The pop­ ulations of the upper levels of these lines however are affected also by collisional excitation from the He° 2s metastable level. This results to a non-negligible increase (of the order of a few percent) in the strength of the aforementioned

He I lines that has to be removed, before deriving abundances via Eq. 1.15. The appropriate corrections have been worked out by Clegg (1987) and Kingdon & Ferland (1995a); in this work, the formulae stated in the latter paper were used to estimate the necessary corrections. In order for the total He abundance to be obtained, one has also to add in the amount of He^'^. The He^"^/H"^ ionic fraction is usually derived from observations of the optical He II A4686 line. 1.3. Determination of elemental abundances in nebulae 30

1.3.2.2 Heavier elements: Carbon, Nitrogen, Oxygen, and Neon

Apart from the strong hydrogen and helium recombination lines and the colli­ sionally excited lines of heavier elements, deep-exposure spectra of nebulae reveal the presence of numerous, weak emission features. As it turns out, most of these are permitted lines of heavy ions, such as O"*", N"*", Ne"*", C"*", etc, excited by the same physical mechanism as the hydrogen lines, i.e. recombination. These lines are weaker than the hydrogen lines by roughly the ratio of heavy element to hydrogen abundances (~ 10~^); their accurate recording therefore, has been made possible only rather recently with the arrival on the scene of high-quantum efficiency, large format charge-coupled devices (CCDs). The majority of heavy element recombination lines can be found in the optical part of the spectrum between 4000 and 5000 A, corresponding to 3s-3p, 3p-3d and 3d-4f transition arrays of C I I , N I I , O II and Ne I I . The C II A4267 is usually the strongest heavy element recombination line detected from both PN and H II regions. Since these lines are excited by recombination they are valuable abundance indicators; furthermore, abundances derived from them (via Eq. 1.15), relative to hydrogen, have the advantage that they are almost completely insensi­ tive to the thermal and density structure of the nebulae under study, since the line emissivities have very similar temperature and density dependencies (Eq. 1.4). Early studies of heavy element ORLs from PN showed that they yielded abun­ dances systematically higher than those derived from CELs for the same ion. Early comparisons involved the C^"*"/H"*" ratio derived from the C II A4267 ORL and the

C III] A1908 CEL; the latter had just become accessible to astronomers with the launch of the lU E (e.g. Harrington et al. 1980, Barker 1982, Kaler 1986, Barker 1991). It was found that in PN the ORL C^‘*"/H"*" ratio was in many cases several times higher than the CEL one (reaching up to factors of 10). The validity of recombination as the means of excitation for the A4267 line was questioned and alternative mechanisms were proposed, such as resonance fluorescence by the hard radiation field close to the central stars, even though that was shown earlier not 1.3. Determination of elemental abundances in nebulae 31

to be the case from model analysis of the Orion nebula spectrum (Grandi 1976). Enhanced, unlooked for, dielectronic recombination contribution to the transition was investigated and shown to be insignificant by Storey (1981) and Nussbaumer & Storey (1984). The effective recombination coefficient was calculated with good accuracy by Pequignot et al. (1991) and did not seem to warrant any major up­ ward revision—as it had been suggested by Kaler (1986). Line-blending was also suspected by some authors—but rejected by e.g. Clegg et al. (1987) in their case study of the planetary nebula NGC3918. The discrepancy remained as detection methods improved and CCDs took over from photographic plates and image photon counting systems. Effective recombi­ nation coefficients were also calculated accurately for O II and N I I . ORL 0^+, abundances were then also found to be significantly higher (by factors of 5) than those obtained respectively from [0 III] AA4959, 5007 and N III] A1751 CELs for the planetary nebula NCC7009 by Liu et al. (1995; LSBC hereafter). In their study of NCC6153, Liu et al. (2000) showed that in that nebula ORL C^'*', O^"*" and Ne^"^ abundances were all higher than those derived from CELs for the same ions, by uniform factors of 10. It thus became clear that careful investigation of the nature of the discrepancies had to be pursued. The problem was tied to the intricate thermal and density structure of nebulae, which defies simplistic charac­ terization and has serious implications on the correct interpretation of abundance determinations from either CELs or ORLs. The reader who wishes to see examples of UV and optical nebular spectra can go now to Chapter 3.

1.3.3 Temperature and density variations in nebulae: implica­ tions on abundance determinations

It was mentioned previously that for abundances to be derived accurately from collisionally excited lines, a necessary requirement is the adoption of a reliable

Tq, since the emissivities of CELs depend exponentially on the temperature of 1.3. Determination of elemental abundances in nebulae 32

the emitting medium. Hence the single most important factor that limits the accuracy of the calculations is the error in our knowledge of the nebular electron temperature. In particular, by virtue of the level population factor in Eq. 1.14, any overestimation of the adopted Te will lead to an underestimation of the derived CEL abundance and vice-versa. It is obvious that the effect is more pronounced for CELs arising from high-lying energy levels, since the populations of those levels

or similarly their excitation energies, Eex, depend more strongly on Tq\ when we have the choice, it is better to derive an ionic abundance using a CEL of low Tex (e.g. an optical line) rather that one of high Tex (such as a UV line). Tied to this is the fact that CEL emission, in general, is preferentially weighted more towards higher-than-average Te nebular regions, assuming that nebulae are not thermally homogeneous objects. In that case it also happens that CEL tem­ perature sensitive ratios, such as the [0 III] (A4959 4- A5007)/A4363 ratio, tend to yield the temperature of the hotter-than-average medium, (e.g. that emitting more strongly in A4363 whose Tex is higher that that of the nebular lines). For these reasons, there has been a long-standing debate on the accuracy of the derived nebular abundances of heavy elements (C, N, O, Ne etc.) measured by means of their CELs, especially ever since it w e ls realized that Tg’s measured by means other than CEL ratios yield different results. In particular, Peimbert (1967, 1971) first discovered that the Balmer jump temperatures in a number of PN and several regions of the Orion nebula are systematically lower than those derived from the [O III] forbidden-line ratio. Recall also, that the Balmer jump temperatures will tend to reflect more the Tg in the cooler-than-average emitting regions, in the presence of temperature variations within the nebular volume (Section 1.2.2). Further studies by Liu & Danziger (1993b; and incorporating results from Barker 1978) confirmed that, on average,

Balmer jump temperatures in PN tend to be lower than the [O III] temperatures, thus strengthening the case for the existence of temperature variations in nebulae. In this work BJ temperatures have also been derived for a number of PN and l.S. Determination of elemental abundances in nebulae 33

t w o H ll regions and have been found to be lower, on average, than Tg([0 I I I ] )

tem peratures.

In the light of these results, the accuracy of the CEL derived abundances has been questioned, on the basis that in the presence of temperature fluctuations the traditional temperature diagnostic—the [O III] nebular to auroral line ratio— would be biased towards high Tg regions and hence the derived heavy element abundances would be substantially underestimated. The validity of the derived heavy element abundances by means of CELs has of course been put under even closer scrutiny, ever since it was discovered that ORL abundances in PN are often quite higher than that (cf. Section 1.3.2.2). In order to characterize temperature fluctuations and evaluate their effects on abundance determinations Peimbert (1967) developed the formalism a brief overview of which is given here. For a given ion he defined a mean tempera­ ture To and a temperature fluctuation parameter given by

and mx",= He then expanded the emissivity function, e(Te) (Eq. 1.9), of a CEL in a Taylor series about the mean temperature T q , keeping terms only up to second order (for relatively small temperature variations), so that the intensity of the CEL would be /(X‘+,A) = J [e(To) + ^ ê ] Ne AT(X*+) dV. (1.18)

Eq. 1.18 can be used to calculate the average temperature, T(A), of the emitting region in a given line A, given that To and are known.

Using our knowledge of Te([0 I I I ] ) and Tg(BJ) from the observations, Tq and can be obtained from the following formulae;

Te([0 III]„a) = ro(02+) {l + i [5^ ^ - 3] i"(02+)} (1.19) 1.3. Determination of elemental abundances in nebulae 34

and Te(BJ) = 3o(H+) [1 - 1.67i2(H+)], (1.20) where it is assumed that To(0^'*') = To(H'*‘) and 1^(0^’^} = In this way the true ionic abundance, in the presence of temperature fluctuations, can be derived from Eq. 1.14, if T(A) is used to calculate the emissivity of the CEL and of H/3, instead of Te([0 I I I ] ) which is the standard thermometer in nebular studies. The concept of temperature fluctuations and their possible effects on abun­ dance determinations has been further explored in the works of Kingdon & Ferland (1995b), Peimbert, Torres-Peimbert & Luridiana (1995), Mathis, Torres-Peimbert & Peimbert (1998) and Liu et al. (2001). Another issue which complicates empirical abundance determinations, is the the fact that nebulae are usually far from being homogeneous objects of uni­ form structure. Imaging studies with the Hubble Space Telescope (HST)., ground- based imaging and kinematical studies and models, reveal that nebulae exhibit widely varying appearances, while their volumes are replete with features de­ scribed as knots, filaments, arcs, bullets and other microstructures (cf. PN studies by Meaburn et al. 1992, O ’Dell & Handron 1996, Lopez, Meaburn & Holloway 1998, Meaburn et al. 1998a, Meaburn et al. 1998b, Meaburn, Lopez &: Noriega- Crespo 2000, Corradi et al. 2000). Long-slit spectroscopy samples a narrow slice of the nebula (be it a PN or an Hll region), while in the best of cases, a scanning long-slit can be used so that spectra averaged over the surface of a nebula are obtained—as is the case for some of the data analyzed in this study. Still the 3D structure of the objects is smoothed out and the derived properties, such els tem­ peratures, densities and abundances, correspond to an average over the nebular volume. Consider also the fact that different emission processes (recombination versus collisionally excited emission) are affected differently under varying temperature and density conditions (Section 1.2), so that the derived abundances may be bicised accordingly. Viegas & Clegg (1994) showed through empirical modelling of PN 1.4- Interstellar extinction 35

that the existence of condensations with Nq ~ 10® cm~® can have a significant effect

on the temperature sensitive [O III] (A4959 + A5007)/A4363 ratio via collisional suppression of the nebular AA4959, 5007 lines, which have Ncr = 6.9 xlO®cm“ ® (a hint about this effect was given in Section 1.2); the auroral A4363 line would not be affected since its critical density is much higher, at 2.5 xl0^cm“®. This would lead to an observationally derived higher Tq and subsequently underesti­ mated CEL abundances. The process therefore leads the observer to deduce lower than true abundances, through an effect of pseudo-temperature fluctuations, when in reality density variations are to blame. As mentioned above, imaging and kine­ matical studies point towards the existence of density inhomogeneities in nebulae. Spectroscopically however, the case that sufficiently dense condensations of ionized (or even partly ionized) nebular matter cause the perceived forbidden-line versus BJ temperature disparity, thus bringing about the ORL versus CEL abundance discrepancies, is more difficult to make as will be shown in the course of this work. The theme of density variations and their effects on the temperature structure and abundance determinations in nebulae has been explored, apart from Viegas & Clegg (1994), by Mathis, Torres-Peimbert & Peimbert (1998) and Liu et al.

(2001).

1.4 Interstellar extinction

Before proceeding with any kind of formal analysis, the observed nebular spec­ tra need to be corrected for the effect of interstellar extinction. Extinction is primarily caused by the preferential absorption and scattering of light with de­ creasing wavelength, by interstellar dust particles. The simplest interstellar dust model consists of a population of tiny, submicron sized carbonaceous and sili­ cate grains (e.g. Weingartner & Draine 2001). The extinction increases towards shorter wavelengths, thus reddening the light from an emitting source; light which 1-4- Interstellar extinction 36

is propagating through interstellar dust is reduced according to the equation

h = h o e -^ \ (1.21) where I\o is the intensity that would be received at the earth in the absence of interstellar extinction along the line of sight, I\ is the intensity actually observed and T\ is the optical depth at the wavelength of observation. By observing pairs of stars and measuring the ratio of their brightnesses, the amount of ‘reddening’ along the path to the more reddened star can be deduced, taking measurements at sufficiently long wavelengths, since for A —^ oc, ta 0. Measurements over the years have shown that in the optical domain the wave­ length dependence, /(A), of extinction is roughly the same along any direction of interstellar space (even though deviations exist), while the total amount, C, of extinction varies, so that TA = C /(A ). (1.22)

In this work, adopted extinction laws /(A), are from: Seaton (1979) for the Galactic UV observations, Howarth (1983) for the Galaxy and the LMC at optical wavelengths, and Prevot et al. (1984) for the SMC at UV and optical wavelengths. In Fig. 1.3 a parametrized form of the extinction law of the Small Magellanic Cloud (Prevot et al. 1984) is shown. The reddening functions in these works are given in terms of analytic fits to

" E (B -V ) ’ where x = 1/A in /im~^, A{x) is the total extinction in magnitudes at x and E{B — V) is the colour excess [{B — V) — {B — V)o] in magnitudes. The ratio of total to selective extinction, Ry = Ay/E{B — V), and a value of 3.1 was adopted for the purpose of this work. It is common in studies of nebulae the above reddening functions to be normalized to H/3 such that 1 .4. Interstellar extinction 37

12

10

8 s Ï 6 I 6 4

2

0 2 3 4 5 6 7 8

1/1 (|im )

Figure 1.3: The extinction curve of the Small Magellanic Cloud (adapted from Prevot et al. 1984). Note the conspicuous absence of the 4.6/xm“^ (2175 Â) feature possibly due to the carbon-poor nature of the SMC dust. The 2175 Â bump is prominent in the Galactic extinction curve.

In nebular studies also, Ay is most usually described by c(H/3), which is the logarithmic difference between observed and intrinsic (or else dereddened) H/3 fluxes. c(H/3) is related to the total extinction and colour excess via e.g.

(1.25) and

c(H^) = lA52B(B-V), (1.26) for the galactic reddening law of Howarth (1983). Once c(H/3) is derived by one of the methods mentioned below, then in con­ junction with Eq. 1.24, the observed nebular emission line fluxes can be corrected for interstellar extinction by multiplying them by 10“^^^^^ 1.5. Aims of this thesis and outline 38

1.4.1 The derivation of c(H/?)

The various methods for the practical determination of c(H/5) are discussed in Chapter 2 (section 2.4.1), where the Balmer decrement and the Hi maps (Burstein & Heiles, galactic studies; Rohlfs et al. 1984, LMC ) methods are explained. The radio-H/3 method will be expanded upon at this point. The radio flux emitted from nebulae does not suffer any extinction so its com­ parison with the reddened flux of a hydrogen Balmer line is a useful measure of interstellar reddening. Milne & Aller (1975) provide a formula for the derivation of c(H/?), which involves the integrated, observed H/3 flux, F(H/3), of a planetary nebula (in erg cm"^ sec“^), compared to the H/3 flux predicted from the free-free radio flux density, (in Jansky) at 5 GHz, the electron temperature t (in units of 10 kK) and the He/H, He^'*'/He fractions. According to that,

where x ' = He^^/He and ^ = He/H. Accurate 5GHz fluxes are available for most PN from VLA surveys by Zijlstra, Pottasch & Bignell (1989) and Aaquist & Kwok (1990). Integrated H/3 fluxes and Si^{5 GHz) fluxes can be found in the compilation by Cahn, Kaler & Stanghellini (1992; CKS92 hereafter). In this work c(H/3) extinction coefflcients were derived from the radio-H/3 method for a number of PN (see Chapter 2, section 2.4.1 and Appendix) and used to deredden UV emission line fluxes from lU E observations and integrated F(H/3) fluxes from CKS92.

1.5 Aims of this thesis and outline

In this thesis a thorough analysis of C, N, O recombination line spectra from both

galactic and Magellanic Cloud planetary nebulae (Chapter 2) and H II regions (Chapter 4) is performed. Elemental abundances, relative to hydrogen, are derived from both CELs and ORLs and an extensive comparison of the different results is made. The main goal of the study is to investigate possible patterns in the 1.5. Aims of this thesis and outline 39 discrepancies, involving ionic lines from the UV, through the optical, to the far- IR. Possible connections of the problem with the evolutionary stage of PN are pointed out. An empirical composite model of the planetary nebula NGC 5882 (Chapter 3) is constructed, lending insight on the behaviour of different emission lines under varying thermal and density conditions in the nebula and indicating physical realities that may fit the observations. The conclusions are presented in Chapter 5. A catalogue of all measured emission-line fluxes can be found in the Appendix. C h a p t e r 2

Elemental Abundances in Galactic and Magellanic Cloud Planetary N ebulae

In this Chapter I present the emission line analysis of a sample of fifteen Galac­ tic and three Magellanic Cloud planetary nebulae. A multi-wavelength study of their thermal and density structure and of the nebular elemental abundances is performed. Abundances from both collissionaly excited (CELs) and optical re­ combination lines (ORLs) are derived. The extensive comparison involves also nebulae analyzed previously with similar methods. The study presents us with important new insight on the issue of ORL/CEL abundance discrepancies.

2.1 Optical Observations

The Planetary Nebulae (PN) studied in this work were observed at the European Southern Observatory (ESO), using the 1.52-m telescope for the Galactic objects and the 3.5-m New Technology Telescope (NTT) for the Magellanic Cloud objects. The long-slit observations were taken during July and December 1995, July 1996 and February 1997. An observational journal is presented in Tables 2.1 and 2.2.

40 2.1. Optical Observations 41

The Boiler & Chivens (B&C) spectrograph was used on the 1.52-m telescope in July 1995, equipped with a Ford 2048 x 2048, 15/rm x 15/rm CCD detector. This detector was superseded in 1996 and 1997 by a UV-enhanced Loral 2048 x 2048,

15yum X 15/im CCD of improved quantum efficiency. A 2 arcsec wide, 3.5 arcmin long-slit was employed. During all runs the CCDs were binned by a factor of 2 along the slit direction, in order to reduce the read-out noise, yielding a spa­ tial sampling of 1.63 arcsec per pixel projected on the sky. A 2400 lines per mm grating was used in first order, along with an order-sorting WG360 filter to cover the AA3995-4978 wavelength range at a spectral resolution of 1.5 Â FWHM. A second grating in first order, along with a WG345 filter, was used to cover the AA3535-7400 range at a resolution of 4.5 Â FWHM (excluding NGC 2440 for which 4.5 Â resolution spectra were not obtained). For three objects (NGC 5882, NGC 6302 and NGC 6818) an extra wavelength range (AA3040-4040) was covered with the holographic 2400 lines per mm grating at a resolution of 1.5 A, so that the high-order Hi Balmer lines—transitions n-^2, from upper levels of principal quantum number up to n = 24—, were accurately recorded. The lower resolution (4.5 A), wider spectral coverage spectra, yielded fluxes of forbidden lines of 0 +, S’*”, N+, CP+, Ar^+, Ne^+ and 0^+ ions, as well as H, He and He+ lines, and the Balmer jump. The high resolution (1.5 A), deep spectra recorded the optical recombination lines (ORLs) of O I I, C II, C I I I , N I I, N II and Ne I I , the majority of which lay between 4000-5000 Â. Typical exposure times for these galactic PN ranged between 20-sec and 35- min. The short exposures were carefully chosen so that strong nebular lines, such as HcK, H/? or the [O III] AA4959, 5007, [N ii] AA6548, 6584 lines, would not be affected by CCD saturation effects. On the other hand, the deep exposures were aimed at capturing the faint ORLs, whose typical intensities can be as low as 10“^ that of H/?, at an appropriate resolving power. For nine out of twelve observed galactic nebulae (NGC entries 2022, 2440, 3132, 3242, 3918, 5315, 5882 and IC 4406, IC 4191) mean spectra were obtained 2.1. Optical Observations 42 by uniformly scanning the long-slit of the B&C spectrograph across the nebu­ lar surfaces (IC 4406 was scanned only at a lower resolution of 4.5 Â). For these nebulae total fluxes of all detected lines can be derived, when total nebular H/3 line fluxes obtained from the literature are taken into account, e.g. large-aperture H/3 fluxes from the compilation of Cahn, Kaler & Stanghellini 1992 (hereafter CKS92). Furthermore, these mean spectra are directly comparable with UV and IR spectra obtained from space-borne facilities (such as the lUE, IR A S and ISO)., whose large apertures usually capture most or all of the nebular emission, de­ pending on the angular size of the observed nebulae and the pointing coordinates of such observations. In this way, several ionization stages of the same element can be traced, allowing a more complete picture for properties such as reddening, excitation class and elemental abundances to be drawn. For the remaining three Galactic PN, plus IC 4406 in the AA3995-4978 range only, flxed-slit spectra were obtained. In those cases, the slit was positioned through the nebular center—using the central star of the PN as a guide, if visible— and passing through the visually brightest parts of the nebula. The Magellanic Cloud PN LMC N66, LMC N141 and SMC N87 were observed at ESO with the 3.5-m NTT, in December 1995. The observing log is presented in Table 2.2. The ESO Multi-Mode Instrument (EMMI) was used in the following modes: red imaging and low dispersion grism spectroscopy (RILD), blue medium dispersion spectroscopy (BLMD) and dichroic medium dispersion spectroscopy (DIMD). The detector was a TEK 1024 x 1024, 24/im x 24^m CCD (no. 31), used while in BLMD observing mode and a TEK 2048 x 2048, 24//m x 24/im CCD (no. 36), while in RILD mode. Both cameras are in use when observing in DIMD mode. In those cases, a dichroic prism is inserted into the beam path so that light is directed to the blue and red grating units in synchronization, allowing simultaneous exposures to be obtained in the blue and red part of the optical spectrum. In all exposures, both CCDs were binned by a factor of 2 in both directions, in order to reduce the read-out noise. The spatial sampling was 2.1. Optical Observations 43

thus 0.74 and 0.54 arcsec per pixel projected on the sky, for CCDs no. 31 and no. 36, respectively. Four wavelength regions were observed with two different gratings #7) at spectral resolutions of approximately 2 A FWHM (AA3635-4145, AA4060-4520, AA4515-4975), and 3.8 Â (AA6507-7828), respectively. An extra wavelength range, AA3800-8400, was covered only for LMC N66 at a resolution of 11 A, using a grism unit (#3). An OG530 filter was used when observing in DIMD mode. The slits used were 5.6 arcmin long and 1, 1.5 arcsec wide, but wide-slit 5 arcsec grism observations were also taken for LMC N66. Exposure times for these Cloud PN ranged from 180-sec to 30-min.

2.1.1 Data reduction

The two-dimensional galactic PN spectra were reduced with the LONG92 package within M ID A S . They were bias-subtracted, fiat-fielded via division by a normal­ ized flat field, cleaned from cosmic-rays, and then wavelength-calibrated using exposures of a He-Ar calibration lamp. The effect of atmospheric extinction was removed by correcting the spectra using extinction coefficients for the La Silla site.

During the 1997 run, twilight sky fiat-fields were also obtained with the purpose to correct the small variations in illumination along the slit, which were ~ l - 3 per cent. For the 1995 and 1996 runs, the spectra were flux-calibrated using wide-slit (8 arcsec) observations of the HST standard stars FeigellO and (the central star of) NGC 7293 (Walsh 1993). The flux distribution of the standards was modeled with high-order spline fits. In 1997, the CTIO standard stars LTT4363, LTT6248

(Hamuy et al. 1994) and the HST standard HD 49798 (Walsh 1993) were used. The former part of the data reduction was done by X.-W. Liu, i.e. the conver­ sion of the raw galactic PN frames to wavelength and flux calibrated 2D spectra. Subsequent reduction and analysis of the galactic PN data, as well as the full reduction and analysis of the Cloud PN data was done by me. The AA3040-4040 spectra of NGC 5882, NGC 6302 and NGC 6818 are affected 2.1. Optical Observations 44

Table 2.1: Journal of ESO 1.52m observations.

PN Date A-range FWHM Mode Exp. Tim e (UT) (A) (A) (sec)

ESC 1.52-m

NGC 3242 7 Feb 1997 3530-7428 4.5 scanning 25, 480 8,10 Feb 1997 3994-4978 1.5 ” 1614, 2 X 1800, 2100

NGC 5882 12 Jul 1996 3040-4050 1.5 fixed slit 2 X 1800

7 Feb 1997 3530-7428 4.5 scanning 20, 220

9,10 Feb 1997 3994-4978 1.5 ” 3 X 1500, 1800

NGC 5315 7 Feb 1997 3530-7428 4.5 scanning 30, 240

9,10 Feb 1997 3994-4978 1.5 ” 1080, 1350, 2 X 1500

NGC 3918 7 Feb 1997 3530-7428 4.5 scanning 20, 22, 40, 200 8,10 Feb 1997 3994-4978 1.5 ” 1020, 1200, 2 X 1620

NGC 2022 7 Feb 1997 3530-7428 4.5 scanning 40, 1155 8,10 Feb 1997 3994-4978 1.5 ” 3 X 1500, 1800, 2100

NGC 6818 11 J u l1996 3040-4050 1.5 fixed slit 2 X 1800

29 Jul 1995 3530-7428 4.5 60, 300 23 Jul 1995 3994-4978 1.5 ” 3 X 1800

NGC 3132 7 Feb 1997 3530-7428 4.5 scanning 70, 1147

8,10 Feb 1997 3994-4978 1.5 2 X 1620, 2 X 1800

NGC 2440 9 Feb 1997 3530-7428 4.5 scanning 3 X 180, 600 9 Feb 1997 3994-4978 1.5 ” 1800

NGC 6302 11 Jul1996 3040-4050 1.5 fixed slit 2X 1800 29 Jul 1995 3530-7428 4.5 ” 30, 60, 300

23 Jul 1995 3994-4978 1.5 ” 3 X 1800

IC 4406 7 Feb 1997 3530-7428 4.5 scanning 50, 1080

29 Jul 1995 3530-7428 4.5 fixed slit 2 X 300 22 Jul 1995 3994-4978 1.5 ” 5 X 1800

IC4191 7 Feb 1997 3530-7428 4.5 scanning 2 X 30, 500

8 Feb 1997 3994-4978 1.5 ” 900, 1200

29 Jul 1995 3530-7428 4.5 fixed slit 5, 10, 300, 600 22 Jul 1995 3994-4978 1.5 ” 3 X 1800

My Cn 18 29 Jul 1995 3530-7428 4.5 fixed slit 2 X 120, 600 23 Jul 1995 3994-4978 1.5 ” 4 X 1800 2.1. Optical Observations 45

Table 2.2: Journal of NTT 3.5m observations.

PN Date A-range FWHM Mode Exp. Time (UT) (A) (A) (sec)

N TT 3.5-m

SMC N87 17 Dec 1995 6507-7828 3.8 fixed slit 1200

3635-4145 2 ” 1200

4060-4520 2 ” 2 X 1800 4515-4975 2 ” 900, 3 X 1800

LMC N66 17 Dec 1995 3800-8400 11 fixed slit 2 x 300

6507-7828 3.8 ” 300

3635-4145 2 ” 300

4060-4520 2 ” 360

4515-4975 2 ” 600

LMC N141 17 Dec 1995 6507-7828 3.8 fixed slit 1200 3635-4145 2 1800 ” 4060-4520 2 2 X 1800 4515-4975 2 ” 180, 4 X 1800 by ozone absorption bands, shortwards of 3400 Â (Schachter 1991). The bands have a typical width of ~15 Â, i.e. much wider than the nebular emission lines (~1.5 Â). In order to remove the ozone absorption we used the ozone-opacity tem plate derived by Liu et al. (2000) for their study of NGC 6153 from the same site (ESO, La Silla) as the observations presented herein. The standard stars Feige 110 and the nucleus of NGC 7293 were observed with a narrow 2-arcsec slit, i.e. the same as the one used for nebular observations. These narrow-slit spectra were used to derive the ozone opacity per unit airmass as a function of wavelength, relative to the mean atmospheric exctinction curve at La Silla. The opacity curve, scaled by the airmasses of the nebular exposures, was then used to divide out the ozone absorption bands. The Cloud PN two-dimensional spectra were reduced with similar procedures as above; the wavelength calibration was done with respect to Th-Ar and He-Ar 2.1. Optical Observations 46

calibration lamps within M ID A S , while f l u x calibration was performed with respect

to wide-slit (8 arcsec) observations of the standard Feige 110, within the I R A F software package. For all flux-calibrated spectra suitable sky windows were selected on each side of the nebular emission, their increments were summed-up, scaled to the total number of CCD pixels along the slit direction and subtracted. In this way we were able to subtract the sky emission for all PN, since their angular diameters were always effectively contained within our long-slit. The resulting spectra were subsequently averaged along the spatial nebular emission—defined by the extent of the B./3 line—to yield the one-dimensional spectra on which our nebular analysis was performed. In a couple of cases, such as PN NGC 6302 and NGC 6818 for which only flxed-slit spectra were obtained, the ORLs detected were rather weak and the spectra slightly noisier than the norm, so we decided to integrate the spectra along the spatial extent of the C II A4267 line which is one of the strongest heavy element recombination lines in PN (of course in this case the H/? line was measured over exactly the same extent as A4267). This resulted in better S/N ratios for the O II and N II lines too, which are expected to originate in similar nebular volumes as the C II A4267 line. In other objects it was found that light from the PN nucleus affects consid­ erably the nebular line fluxes or the shape of the observed continuum spectrum. For instance, NGC 3132 harbours a binary nucleus whose primary component is the A-type star HD 87892 (Kohoutek & Laustsen 1977, Sahu & Desai 1986). The absorption line spectrum of this star affects severely the nebular Balmer lines by diminishing their intensities, so a few central CCD pixels had to be excluded from the integration in order to prevent such contamination and to restore the nebular line profiles. Another example is NGC 5882 which possesses a weak emission line {'weV), hydrogen-deficient nucleus whose continuum energy distribution evidently affects the shape of the nebular continuum in the region of the Balmer disconti­ nuity (~3646 Â). Thus in order to get an accurate measurement of the Balmer 2.2. The lUE observations 47 jump intensity and the high-order Balmer lines from the deep AA3040-4040 spectra of this object, we again excluded several central pixels when averaging the two- dimensional CCD frames. This restored the nebular continuum to its expected shape at those wavelengths.

2.2 The lUE observations

We complemented our optical dataset with International Ultraviolet Explorer (lUE) observations for 13 PN, the Cloud PN included; the remaining two galactic PN IC 4191 and My Cn 18 were never observed by the lUE. Low-resolution, large aper­ ture spectra obtained with the SWP and LWP cameras were accessed via the MAST at the STScI website. The wavelength coverage of the two cameras is 1150-1975, and 1910-3300 Â, respectively. All observations were obtained with the lU E large aperture, a 10.3 x 23 arcsec^ oval. The data were retrieved in final calibrated form. When several spectra for one object were available, they were co-added weighted by the integration time. For NGC 3918 the line fluxes reported by Clegg et al. (1987) were used, combined with those retrieved from a more recent exposure (SWP47514). Similarly, in the analysis of NGC 5315 and NGC 6818 I used the line fluxes recently listed by Feibelman (1998; see his Table 3) and by Hyung, Aller & Feibelman (1999, hereafter HAF99; see their Tables 4 and 5), respectively. The apparent angular sizes of all PN in our sample, apart from NGC 5315 and the Cloud nebulae, are larger than the lU E large aperture (see Table 2.21). Thus, in most cases only a certain fraction of the UV nebular emission was captured by the satellite. This delicate aspect of the observations is discussed in the following Section. 2.3. Data analysis 48

2.3 Data analysis

With respect to the optical data we proceeded by normalizing the individual one­ dimensional spectra to a flux scale such that -F(H/3) = 100, in order to merge them afterwards. When H/3 was found to be saturated in one spectrum, the scaling was

achieved using a suitable unsaturated H I B aimer line from a shorter duration exposure. The AA3040-4040 galactic PN spectra were brought to scale with the rest, using emission lines detected in common with the AA3535-7400 spectra, such as the H 9 A3835, H 8 A3889 or [Ne ill] AA3969, 3967 lines. In a similar fashion, the Cloud PN spectra from different grating settings were brought to scale via their overlapping portions. For two of them, SMCN87 and LMCN141, their AA6507- 7828 spectra had no overlap in wavelength with other grating settings. In this case we normalized the spectra to F(H/0) = 100, using the theoretical intensity ratio /(Ha)//(H/?) = 2.85, a value predicted for representative nebular conditions of Te = 10 000 K and Nq = 5000 cm“ ^ (Storey & Hummer 1995); we also took into account the amount of reddening estimated towards these objects (see Section 2.4.1). Finally, the normalized spectra taken at each grating setting were co-added, weighted by the exposure time, in order to achieve an optimum S/N ratio. Regarding the UV data, a way had to be found to scale them to the normal­ ized optical spectra. For all galactic nebulae (except NGC5315, NGC6818 and

IC 4406), we used the theoretical intensity ratio for He II A1640/A4686, calculated for the appropriate nebular temperatures and densities using the recombination theory predictions of Storey & Hummer (1995). In this way we scaled the UV spectra to T(H/3) = 100 using our measurements of the dereddened /(A1640) and /(A4686) fluxes. In cases where the [O ll] A2470 doublet line was detected, as in NGC 6818 and IG4406, we used instead the predicted intensity ratio [O ll] (A7320 + A7330)/A2470 = 1.33, which holds independently of nebular conditions since all lines involved arise from the same upper levels (vl-values taken from Zeippen 1982). The method was found to be highly reliable; for instance in NGC 5315 whose angular diameter fits completely within the lU E large aperture—and there 2.3. Data analysis 49

is thus no need for any aperture correction to be applied—the [O II] A2470 intensity predicted using the optical [O ll] AA7320, 7330 intensities, is within 2 per cent of the actual UV measurement. For NGC 5315 and the three Cloud PN (SMC N87, LMC N66 and LMC N141) the dereddened UV line fluxes were scaled directly to F(H/?) = 100, with the use of total H/5 fluxes from CKS92 in the former case and from Meatheringham et al. (1988), in the latter. We proceeded with measuring the emission line fluxes from the UV and optical spectra using the M ID A S package. First we scrutinized carefully the low resolution lU E spectra in order to determine the nature of the emission lines, which in some cases can have a stellar and not nebular origin (e.g. N V A1240, C IV A1550, He II A1640). The FWHM of the lines was compared against that of lU E low resolution spectra which is ~6-7 Â. Lines with FWHM of > 8-9 Â were checked for blending or for the possible presence of P Cygni features that would betray their origin in a central star’s wind, as is the case for instance in NGC 5315 and NGC 5882. Most of the line fluxes—and certainly all those of heavy element recombination lines—were derived using Gaussian line profile fitting techniques, apart from the strongest ones, for which their fluxes were measured by simply integrating over the lines. Gaussian line fitting within M ID A S is performed through least-squares fitting using the Newton-Raphson method. Even in our 1.5 Â FWHM high-resolution optical spectra, blending of emission lines is evident throughout the wavelength range covered and it particularly af­ fects weak features like the ORLs of interest. The line broadening is dominated by instrumental broadening and not by nebular dynamics, since the typical FWHM of 1.5 Â translates to ~93 kms“^ at 4861 Â, whereas typical PN expansion veloci­ ties are <25 kms“^ (e.g. Kwok 1994). Therefore, in order to de-convolve features affected by blending, multiple Gaussian fitting was employed. In such cases a successful estimate of the continuum emission level was deemed to be a very im­ portant first step. After subtracting the local continuum, line fluxes were retrieved by fitting multiple Gaussians of appropriate central wavelength and usually equal 2.3. Data analysis 50

FWHM. The FWHM was taken to be the same as that of nearby unblended lines of similar strength to the ones fitted. Their relative wavelength spacing was con­ strained to that from laboratory wavelengths. This procedure assured accurate fiux retrieval and aided line identification in the case of ambiguous features. In rare instances where the fitting process was not adequately yielding unique fluxes for closely spaced lines, we made use of the theoretical intensity ratios of selected lines, assuming F5'-coupling to hold for the ions in question (e.g. Kuhn 1969). As an example, in PN of relatively high excitation class, like NGC 3242 or NGC 5882, the [S ll] A4068 line is blended with the C I II V 16 triplet, whose components at 4067.94, 4068.91, and 4070.26 Â have intensity ratios under pure LF-coupling of 1.00 : 1.31 : 1.71 (Kuhn 1969, Allen 1972). Also present in this blend are the O II V 10 multiplet AA4069.62, .89 and 4072.16 lines. We therefore fitted 6 Gaussians to the 4069 Â complex, fixing the C I II triplet relative intensities to 1.00:1.31:1.71, while allowing the intensities of the other lines to vary independently. After reducing by 2 the number of free parameters in the minimization prob­ lem, convergence was more readily achieved. Of course in this way, the resulting fiux for the C I I I triplet refers to the sum of its components only, while information on the true allocation of fiux among them is lost. Individual fluxes, however, are retrieved for the [S ll] and O II lines and the overall shape of the blend is accurately reproduced (see Fig. 2.1). As another example, the blend at 4650 Â can be fitted with 5 Gaussians, including the O I I AA4649.13, 4650.84 lines from multiplet V 1 and the C I II V 1 triplet AA4647.42, 4650.25 and 4651.47 lines, the latter having relative intensities of 5 : 3 : 1 as predicted under L5-coupling. In general, apart for the aforementioned cases we avoided fixing the relative strengths of ORLs, since we were interested to compare our observations with the predictions of current recombination theory, especially for O II. Typical line fiux errors are estimated to be less than 5 per cent for lines with observed fluxes F{X) > 0.2 [in units of F(H/?) = 100], 10 per cent for those with 2.3. Data analysis 51

0.40

0.35

0.30

0.25

5 0.20

0.10

0.05

0.00

-0.05 4064 4068 4072 4076 4080 4084 4088 4092 Wavelength (A)

Figure 2.1; Medium-high resolution spectrum (1.5 Â FWHM) of the medium excitation PN NGC 5882: the blend of O il quartet V 10, C ill triplet ,V 16 and [S ll] doublet is well htted by multiple Gaussians {thick line)-, also present are O II 4f-3d transitions of which A4089.3 (J = 11/2-9/2) is the strongest predicted line. See text for more details.

0.1 < F{\) < 0.2, 20 per cent for 0.05 < F{\) < 0.1, and 30 per cent or more for lines with F{\) < 0.05. For a given flux range, the errors for lines shortwards of 3400 À are slightly larger, due to the decreasing CCD efficiency towards the end of the spectral coverage and the effects of ozone absorption. 2 .4. Nebular Analysis 52

2.4 Nebular Analysis

2.4.1 Reddening

Before proceeding with any kind of formal analysis, nebular spectra have to be corrected for the effects of interstellar extinction caused by dust. In nebular studies the amount of interstellar reddening is most usually described by c(H/5), which is the logarithmic difference between observed and intrinsic (or else dereddened) }lf3 fluxes. The reddening-corrected nebular fluxes (in units of H/? = 100) are then given by 7(A) = io

are studied. In this way we derived reddening values for our galactic PN using the formulae of Milne & Aller (1975), employing total H/? and 5 GHz continuum fluxes taken from CKS92, along with He'^'/H"*" and He^"*"/H'*' fractions deduced from our optical observations (see Section 2.6). Similarly, one can also derive c{R(3) from a

comparison between the He II A1640 and A4686 or [0 ll] A2470 and AA7320, 7330 integrated line fluxes. In this work we used the value of c(H/?) derived from the radio-H/3 method, denoted c(H/?)^^^, to deredden the UV line fluxes in cases where the apparent PN diameter was larger than the lU E large aperture. All UV fluxes were then scaled to the optical spectra as described previously in Section 2.3. In our sample, apart from NGC 5315 and the Cloud nebulae, the PN angular diameters are larger than the lU E large aperture, so the derived c(H/5)^®^° values are consistently greater than the one given by c(H/3)'^^^, since the derivation involves a comparison between a fraction of the He II A1640 line flux with the total nebular He I I A4686 flux; therefore c(H/3)^^'^° values were not used in our analysis. The Cloud PN spectra were dereddened in a two step process. First the con­ tribution from Galactic foreground reddening was removed; this was estimated from the reddening maps of Burstein & Heiles (1982), using the extinction law of Howarth (1983) in all cases; the remaining extinction due to the Clouds’ interstel­ lar medium was taken from Barlow (1987) in the case of the UV data, and from the observed Baimer decrement in the case of the optical data. The SMC extinc­ tion law of Prevot et al. (1984) was used in the case of SMC N87, while the mean LMC extinction law of Howarth (1983) was used for LMC N66 and LMC N141. Another kind of correction had to be applied for medium- to high-excitation objects with a high He I I A4686 flux [e.g. > 65 when H/3 = 100]; in such nebulae the He II Pickering series (transitions n—j-4) begins to contribute significantly to the H i Baimer series. Relevant cases include PN NCC 2022, N C C 2440, N C C 6302 and LMCN66, where the contribution to H/1 was estimated using the observed

He I I A4686 flux and emissivities from Storey & Hummer (1995). Fluxes of the 2.4- Nebular Analysis 54

contaminating He II A4859 (n = 8-^4) amounted to ~3-5per cent and were sub­ tracted from /(H/5) before any further analysis was performed. Throughout the above analysis we used total nebular F(H/5) and 5 GHz con­

tinuum fluxes as listed by CKS92; from them total nebular He II A4686 fluxes were derived, using the observed /(A4686)//(H/?) ratios, from our scanned neb­ ular spectra. In the case of nebulae that were observed only with a hxed-slit we proceeded as follows: for NGC 6302 we used the total nebular He II flux from CKS92 directly, without resorting to our hxed-slit 7(A4686)/7(H/3) ratio, while for NGC 6818 the above optical ratio was assumed to be representative of the whole PN.

2.4.2 Nebular Diagnostics

The nebular electron temperatures and densities were derived from several CEL diagnostic ratios by solving the equations of statistical equilibrium for multi-level

(> 5) atomic models using the program EQUIB (originally written by I.D. Howarth and S. Adams). The following diagnostic ratios were used:

Te(0 III): 7(A4959 + A5007)/7(A4363)

7;(N II): 7(A6548 -|- A6584)/7(A5754)

7 ;(0 II): 7(A3727)/7(A7320 A7330)

Te(S II): 7(A4068)/7(A6731 + A6716)

Te(BJ): 7c(A3646- - A3646+)/7(H 11)

A^e(Ar IV): 7(A4740)/7(A4711)

ATe(Cl III): 7(A5537/7(A5517)

74(S II): 7(A6731)/7(A6716)

Afe(0 II): 7(A3729)/7(A3726) Nebular Analysis 55

The procedure to derive temperatures and densities from CEL line ratios is as follows: we assumed a representative initial electron temperature of 10000 K

to derive A^e(Cl III) and Ne{Ar iv). We then used the mean density from these

diagnostics to derive T^(0 III) and iterated once to get the final values. In a similar manner, T^(N ll) was derived in conjunction with A^(0 ll) and A^(S ll). The atomic data used for the purpose of this analysis, but also throughout this study, are listed in Table 2.3.

2.4.2.1 Electron Densities

In all spectra the [Ar iv ] A4711 line is blended with the He I A4713 line; individual fiuxes for the two transitions were thus derived by means of Gaussian profile fitting, as described previously in Section 2.3. Both of the [Ar iv] AA4711, 4740 fiuxes were measured on high-resolution spectra (1.5 A for the galactic PN and 2 A for the Cloud PN) in all cases; they were then used to derive the A^e(Ar iv) values listed in Table 2.4. We were able to derive N ^s using the [O ll] A3729/A3726 ratio only for NGC 5882, NGC 6302 and NGC 6818, for which the [O ii] AA3726, 3729 doublet was resolved in the high-resolution AA3040-4040 spectra. For the remaining PN the [O ll] doublet is covered in lower resolution only and is blended with the H 14

A3721.9, [S III] A3721.6 and H 13 A3734.4 lines; the overall [O ll] fiux was derived as follows: the [S III] A3721 fiux was estimated from a comparison with the dered­ dened flux of the [S III] A6312.1 line, which originates from the same upper level.

The latter line is often blended with He II A6310.8 in high excitation objects and that fiux was retrieved via its theoretical ratio relative to He II A4686. Finally, the /(H 13)//(H/?) and /(H 14)//(H/5) intensity ratios were estimated using Hi line emissivities from Storey & Hummer (1995).

Electron densities from the [O III] 52-/im/88-/im ratio were also derived, using the ISO LWS fiuxes of Liu et al. (2001). The derived electron densities for our sample of PN are listed in Table 2.4. It 2.4- Nebular Analysis 56

Table 2.3; References for atomic data.

Ion Collisionally excited lines

Transition probabilities Collision strengths

C III Keenan et al. 1992 Keenan et al. 1992 Flem ing et al. 1996

N II Nussbaumer & Storey 1979 Strafford et al. 1994

N III Fang et al. 1993 Blum & Pradhan 1992

0 II Zeippen 1982 Pradhan 1976

0 III Nussbaumer & Storey 1981 Aggarwal 1983

Ne II Mendoza 1983 Bayes et al. 1985

Ne III Mendoza 1983 Butler & Zeippen 1994

Ne IV Zeippen 1982 Giles 1981

S II Mendoza & Zeippen 1982a Keenan et al. 1996 Keenan et al. 1993

S III Mendoza & Zeippen 1982b Mendoza 1983

S IV Storey (unpublished) Saraph &; Storey 1999

Cl III Mendoza & Zeippen 1982a Butler & Zeippen 1989

Ar III Mendoza & Zeippen 1983 Johnson & Kingston 1990

Ar IV Mendoza & Zeippen 1982a Zeippen et al. 1987

Optical recombination lines

Ion Effective recomb, coeffs. Case

H I Storey & Hummer 1995 B

He I Brocklehurst 1972 B: singlets A: triplets

He II Storey & Hummer 1995 B

C II Davey et al. 2000 B

C III Pequignot et al. 1991 A

Nussbaumer & Storey 1984

N II Escalante & Victor 1990 B: triplets A: singlets

N III Pequignot et al. 1991

Nussbaumer & Storey 1984

O II Storey 1994 B: quartets

Liu et al. 1995a A: doublets

Ne II Kiselius et al. 1998 B: doublets Storey (unpublished) A: quartets Nebular Analysis 57 should be borne in mind that the derived densities correspond to mean values, since the employed line ratios have been taken from scanned nebular spectra in most cases. In addition, all observations sampled lines of sight throughout the nebular volumes as well. It is thus quite probable that localized density variations on a small scale have been effectively smoothed out, or that they are not even seen due to the relatively low spatial resolution of our observations. On the other hand, the Nq{S ll) values for the whole sample are consistently lower than the A^(Ar i v ) values, by about ^30per cent. This behaviour is con­ sistent with the presence of strong density variations in the nebulae, so that the diagnostic line ratios with higher critical densities (e.g. [Cl III] and [Ar iv ] diag­ nostics) yield higher derived nebular electron densities.^ The effect is shown to be more pronounced from a comparison between the [O III] b2-(im/S8-fjLm and [Cl III] and [Ar iv] densities. In eight PN for which values from all these three ratios have been derived, the latter diagnostics yield Ne's that are on average a factor of six higher than those obtained from the far-IR ratio (see Rubin 1989, Liu et al. 2001).

2.4.2.2 Electron Temperatures

We derived Te’s using ionic CEL ratios of nebular to transauroral [S II] transitions and of nebular to auroral transitions for [N ll], [O ll] and [0 III]. The results are presented in Table 2.5. We were able to derive Te(0 lll)’s for all 15 PN and Te(N li)’s for 13 PN, excluding SMCN87 and LMCN141 for which the auroral A5754 line did not fall within our wavelength coverage. The correct choice of electron temperatures is crucial when deriving abundances from forbidden lines, so in what follows we discuss the pattern of TqS for the current sample of 15 PN, as measured from the above ionic ratios.

^[Cl III] AA5517, 5537, N r = 6 4 0 0 , 34000cm~^ respectively; [Ar iv] AA4711, 4 7 4 0 ,

N cr = 14 000, 1 3 0 0 0 0 respectively; [S ll] AA6716, 6730, N cr = 1200, 3300 cm~^ respec­ tively. 2.4- Nebular Analysis 58

Table 2.4: Electron densities of Planetary Nebulae.

PN [O II] [S II] [G1 III] [Ar iv ] [O III]" A7325/A3727 A6731/A6716 A5537/A5517 A4740/A4711 52/im/88/xm Ne (cm “ ^)

NGC 3242 * 1970 1200 3040 775 NGC 5882 4750 4000 2700 5000 1175 NGC 5315 13040 8200 22825 12300 2290 NGC 3918 * 4600 5500 6900 1380 NCC 2022 1970 1050 850 2150 * IG419T 13530 12750 12375 13750 * IG 4 1 9 P 11930 7900 10150 12800 1700 IG 4406 * 950 3500 1250 540 N CC 6818 1800 1700 2400 2350 * NGC 6302 5750 12900 22450 14900 1380 My Gn 18 * 5025 9420 6300 * NGC 2440 3000 3200 3800 4000 * NGC 3132 * 550 720 530 355 SM G N 87 2850 3950 * 9500 * LMGN141 2300 7400 * 9500 * LM GN66 * 1900 * 5700 *

“ Prom ISO LWS line-fluxes by Liu et al. (2001); ^ Values for the hxed-slit spectrum of the nebula;

^ Values for the whole nebula from a scanned spectrum.

2.4.2.3 Recombination excitation of the N II and O II auroral lines

The mean Te(0 III) for the whole sample is 12 300K, while the mean Te(N ll) is 11 BOOK. For 13 PN where both Te’s were measured, the values differ by about

5 per cent, Te(0 III) being only slightly higher than T^(N ll). It could be argued that generally in PN, N+ is present in a lower ionization zone than and that the tem perature is lower too. However, for 8 PN where both Tg(0 III) and

Te(0 li) can be measured, it is found that on average, Te(0 III) = 11 500 K, while

Te(0 ll) = 13 400 K, i.e. higher by 1900K. The disagreement between [O ll] and

[O III] electron temperatures, as derived by the nebular to auroral line ratios 7(A3727)//(A7320 + A7330) and /(A4959 + A5007)//(A4363), respectively, is at its most extreme in the cases of PN NGC 3242, NGC 5882 and My Cn 18. In these 2.J^. Nebular Analysis 59

Table 2.5: Electron temperatures of Planetary Nebulae.

PN [8 II] [ 0 i i ] [N II] [0 III] BJ To T e(K )

NGC 3242 4800 19550 13400 11700 10200 10770 0.032 NGC 5882 6900 15300 10800 9400 7800 8270 0.034 NCC 5315 11400 * 10800 9000 8600 8770 0.0085 NCC 3918 9350 10400 10800 12600 12300 12430 0.0063 NGC 2022 ** 14700 15000 13200 14080 0.037 IC 4191“ 7750 * 11575 10700 10500 ** IC 419P 8550 * 12225 10000 9200 9470 0.017 IC 4406 8650 9000 9900 10000 9350 9570 0.014 NCC 6818 5750 11200 11100 13300 12140 12650 0.024 N CC 6302 10000 20200 14225 18400 16400 17610 0.041 M y Cn 18 12825 12300 10225 7325 ** * N CC 2440 12000 * 13500 16150 14000 15125 0.045 N CC 3132 8120 * 9350 9530 10780 10330 * SMC N87 >20000 ** 12250 ** * LMC N141 10100 ** 11850 * * * LMC N66 * * 12700 18150 * * *

“ Values for the hxed-slit spectrum of the nebula; ^ Values for the whole nebula from a scanned spectrum.

objects the derived Te(0 ll)’s are higher than the Te(0 lll)’s by 7850 K, 5900 K and 4980 K, respectively. A variety of hypotheses can be invoked in order to explain the difference of electron temperatures in these 3 PN. Radiation hardening in the nebular volumes where singly ionized species are expected to exist, could result in higher tempera­ tures than those of the 0^“^ zone, but the effect should be similar for both Te(N ll) and Te(0 ll), since the ionization potentials of N+ and are very comparable. An alternative proposition is as follows; in NGC 3242, NGC 5882 and My Cn 18 most of the N and O atoms are in their doubly ionized stages (0^+/0 = 0.85, 0.96 and 0.63, while N^"""/N = 0.65, 0.72 and 0.63, respectively; cf. Section 2.5). Thus, as discussed by Rubin (1986) and Liu et al. (2000), recombination of N^"^ and can be important in contributing to the excitation of the [N ll] nebular and [O ll] nebular and auroral lines. 2.J^. Nebular Analysis 60

We will attempt to quantify this effect for these three extreme cases, using the expressions of Liu et al. (2000) (see their Section 3.3; Eqs. 1,2). They employed radiative recombination coefficients from Pequignot at al. (1991) and dielectronic recombination coefficients from Nussbaumer & Storey (1984), to show that the intensity of the [N ll] A5754 auroral line due to recombination excitation can be fitted by

where t = 7^/10^ and 0.5

Jr(A7320 + A7330) _ 0,44 ^ ,2 2)

For NGC 3242, which has 7^(0 III) = 11 700 K, the observed A1750 flux yields N2+/H+ = 1.90 X 10-5, for 7^ = 11 700K and = 1200cm'^ (Tables 2.5, 2.4). Inserting this value into Eq. 2.1, we have 7r(A5754)//(H/3) = 0.0064, or 10 per cent of the observed intensity of the A5754 line relative to Hyd, 7(A5754)/7(H/?) = 0.0635. After subtracting 7r(A5754) from the observed intensity, the [N ll] nebular to au­ roral line ratio yields 7^ = 10 850 K, i.e. 2550 K lower than the value derived before the correction (cf. Table 2.5). If one uses instead = 1.28 x 10“^, as de­ rived from N II recombination lines (cf. Table 2.14), we get 7r(A5754)/7(H/3) = 0.0408, which is about 64 per cent of the observed value. After correcting for this 7r(A5754), the temperature deduced from the revised [N ll] line ratio is only 7950 K. Regarding the [O ll] auroral lines, for = 2.80 x 10"^ eis derived from the AA4959, 5007 lines (cf. Table 2.5), Eq. 2.2 predicts a recombination intensity relative to H/? of 0.262 on a scale where H/1 = 100 for the [O ll] AA7320, 7330 lines, or 29 per cent of their observed intensity. According to Liu et ah, the re­ combination contribution to the AA3726, 3729 nebular lines is 7.7 times that of

7r (A7320-^ A7330), at Te = 10 000 K. Thus we have, 7r (A3726-h A3729) = 2.02 on Nebular Analysis 61 a scale where H/3 = 100, or 21 per cent of the observed overall intensity. Again, if we use instead the ORL abundance ratio of = 6.28 x 10“^ derived for this PN (cf. Table 2.15), we get a recombination contribution of 0.588 to the [O ll] auroral lines (65per cent of the observed value), while a contribution of 4.53 relative to H/? (46per cent of the observed intensity), is found for the nebular lines. Therefore, using as given by the CEL lines, the revised [O ll] nebular to auroral line ratio yields Tq = 15 800K, i.e. 3750 K lower than the value derived before the correction. If we instead use the abundance ratio given by the ORLs, the corrected Te(0 ll) is 7500 K lower, at 12 050K, and only 300K higher than the Te(0 III).

Applying the same procedure to NGC 5882, which has Te(0 III) — 9400 K, we obtain from = 1.10 x 10“ ^, as given by the CELs, a corrected nebular to auroral line ratio which yields a 7^(N ll) of 9900 K; if we assume = 1.63 x 10"^, as given by the ORLs, the corrected temperature is 7^(N ll) = 9550 K. Similarly, the corrected [O ll] nebular to auroral line ratio results in a 7^(0 ll) of 14 500 K, us­ ing = 4.48 X 10“ ^ as given by CELs; in the case of = 9.70 x 10"^, as given by ORLs, the corrected temperature is 7^(0 ll) = 13 750K.

For My Cn 18, which has 7^(0 III) = 7325 K, the [N ll] line ratio does not change appreciably for = 2.20 x 10"^ as given by CELs, so there is a negligible change in 7^(N ll); however if N^+/H+ = 20.4 x lO"'^, as given by the ORLs, the corrected temperature becomes Tg(N ll) = 9450 K. Finally, the corrected [O ll] ratio results in a 7^(0 ll) of 12 150K, using 0^+/H^ = 3.54 x 10“'^ as given by CELs; and for 0^'^/H’*' = 6.43 x 10“^, as given by ORLs for this object, the corrected temperature is 7^(0 ll) = 12 100 K. In the above discussion we have neglected the contribution of recombination excitation on the observed [N ll] AA6548, 6584 nebular intensities; this is estimated to be small and amounts to only 7, 8 and 8 per cent, respectively, for NGC 3242, NGC 5882 and My Cnl8, even when adopting ORL abundances. We 2.4- Nebular Analysis 62

base this on the fact that for pure recombination excitation, the [N ii] nebular to

auroral line ratio has a value of 5.9 at Tq = 10 000 K (Liu et al. 2000). On the other hand, the observed ratios for these three PN are respectively, 10, 12 and 13 times larger than that. We summarize our findings for these three objects in Table 2.6. To conclude, for low-density uniform nebular media, as assumed above, ap­ preciable fractions of the [O ll] nebular and auroral and [N ll] auroral lines can be accounted for by recombination processes. However, even after recombina­ tion excitation has been taken into account, the agreement between the various CEL temperature diagnostics remains quite poor, especially for NGC 5882 and My Cn 18 (see Tables 2.6 and 2.5). It is obvious that exact treatment of this prob­ lem requires detailed knowledge of the actual electron temperature and and abundances. To complicate matters more, if the nebulae are not ho­ mogeneous, but also contain a higher-density component, then the observed [O ll] and [N ll] emission pattern can be more difficult to evaluate. In nebulae that contain condensations whose electron density is higher than the critical densities of the low-lying atomic levels—from which the AA6548, 6584 and A3727 nebular lines originate—but lower than that of the auroral lines, the former lines will be preferentially emitted from the lower density medium. This could lead to apparently high temperatures as deduced from the [N ll] and [O ll] nebular to auroral ratios, since their component transitions would be disproportionately affected in regions of high density; there, the nebular lines will be suppressed by collisional de-excitation, but the auroral lines will not be (Viegas & Clegg 1994). As our schematic treatment has pointed out, the [N ll] and [O ll] nebular to auroral temperature diagnostic ratios provide values which are probably poor indi­ cators of Te, in the low-ionization regions of these three objects. Therefore, in the abundance analysis that follows in Section 2.5, we have adopted Te(0 III) when calculating the forbidden-line N+/H+ and abundances for NGC 3242, NGC 5882 and M yC nl8. 2.J^. Nebular Analysis 63

Table 2.6: The [N ll] and [O ll] electron temperatures after correcting for the effects of recombi­ nation excitation.

PN T e([0 III]) Te(BJ) Te"°^([N Il])“ Te"°R[0 Il])“

(K)

NGC 3242 11700 10 300 10850|7950 15800|12050

NG C 5882 9400 9700 9900|9550 14 500|13 750

M y Cn 18 7325 * 10 225|9450 12150|12100

“ Values before the dash are for recombination excitation contributions calculated adopting CEL

abundances; those after are from adopting ORL abundances.

2.4.2.4 Baimer discontinuity electron temperatures

Apart from temperatures derived from the CEL ratios, Table 2.5 also lists the mean nebular Baimer jump temperatures, Te(BJ), derived from the ratio of the nebular continuum discontinuity at 3646Â to H 11 A3770 [A(BJ)/H 11]. Values of the Baimer jump intensity were determined from the lower resolution 4.5 Â FWHM scanned spectra, apart from the cases of NGC 5882, NGC 6302 and NGC 6818 for which their higher resolution 1.5 Â AA3040-4040 hxed-slit spectra were used. The observed nebular continuum on either side of the Baimer discontinuity was mod­ elled with low-order polynomials. Redwards of 3646 Â the crowding of high order nebular Baimer lines frustrates the determination of the local continuum level on the lower resolution spectra, which was then estimated by linear extrapolation using continuum data points with wavelengths longer than 3800 Â. However, it is felt that no signihcant uncertainties were introduced from this. In the higher resolution spectra the high order Balmer lines are clearly resolved and the lo­ cal continuum is better estimated with no significant extrapolation from longer wavelengths. In practice, the Balmer jump temperatures were derived by comparing the observed and predicted values of the Balmer discontinuity to H 11 ratio, defined

as A(BJ)/H/3 = [7c(A3643) — /c(A3681)]//(H 11). By thus defining the Balmer 2.5. Elemental abundances from CELs 64

discontinuity we include in our computation two weak discontinuities of the He I

and He II continua which are present at 3678 and 3646 A, respectively and are inseparable from the Hi Balmer jump at the resolution of our observations. By definition then, the temperature thus deduced has a weak dependence on the He^/H*^ and He^"^/H"*" abundance ratios; these are presented in Tables 2.7 and 2.8. For the purpose of the above analysis, the emissivities of the Hi Balmer lines

and of the Hi, He I and He ll continua were taken respectively from Storey & Hummer (1995) and Brown & Mathews (1970).

2.5 Elemental abundances from CELs

In this section we present our results for ionic abundances derived from collisionally excited lines for the 15 Galactic and 3 Cloud PN observed, both from UV and optical transitions. These are presented in Tables 2.7 and 2.8 for the Galactic and Cloud nebulae respectively. The analysis of emission lines detected by the lUE satellite gave us access to important ionic stages such as: C'^-C ll] A2326, C^'^-C III] A1908, and G^"""-C IV

A1550 (resonantly excited doublet line); N^+-N III] A1750, N ^+-N iv] A1486, and

N^+_N V A1240 (resonant doublet); 0 + - [O ll] A2470, ill] A1663, and

0 ^ + - 0 IV A1401; Ne^+-[Ne iv] AA1601, 2423, and Ne^+-Ne v A1574. Especially with regard to carbon, the consideration of UV lines facilitated the calculation of forbidden line abundances for this element, which has no emission lines in the optical domain other than recombination lines (of C^'*', C^'*' and C^+). The nitrogen abundance estimate benefits also greatly from the inclusion of UV lines in the analysis, since in the optical region—in terms of CELs—this important element is represented only by the [N ll] AA5754, 6548-84 lines, which are typically indicative of just a small fraction of the total N abundance. Complementary analysis of the more frequently employed optical CELs allowed us to obtain total CEL elemental abundances for C, N, O, Ne, S, Cl, and Ar.

In order to derive ionic abundances we used the EQUIB code, which solves 2.5. Elemental abundances from CELs 65

Table 2.7; (a.) Helium and CEL O, N, C, Ne, Ar, S and Cl abundances in PN.“

NGC 3242 NGC 5882 NGC 5315 NGC 3918 NGC 2022 NGC 6818 4471 He+/H+ 0.0811 0.1102 0.1243 0.0721 0.0128 0.0488 5876 He+/H+ 0.0789 0.1064 0.1205 0.0655 0.0124 0.0482 6678 H e+ /H + 0.0762 0.1036 0.1174 0.0631 0.0121 0.0445 Avg. He+/H+ 0.0789 0.1066 0.1206 0.0663 0.0124 0.0476 4686 He^+/H+ 0.0208 0.0022 0.0000 0.0350 0.0921 0.0510 H e/H 0.100 0.109 0.121 0.101 0.105 0.099 2470 0 + /H + * 4.24e-5: ** 2.37e-5 3727 2.54e-6 1.31e-5 4.19e-5 2.23e-5 1.60e-6 2.28e-5 7320+7330 4.12e-6: 2.17e-5: 4.24e-5: 4.37e-5: 1.02e-6; 2.37e-5: 1663 0 ^ + /H + l.lle-4: 4.35e-4: 6.02e-4: 2.34e-4 5.72e-5 1.51e-4 4931 2.73e-4: 4.38e-4: 3.48e-4: 2.57e-4: 1.04e-4: 1.68e-4: 4959+5007 2.80e-4 4.72e-4 4.26e-4 2.90e-4 8.07e-5 2.52e-4 52+ 88 fj,m 2.20e-4 2.70e-4 8.81e-5 1.64e-4 * * 3.160-4 4.70e-4 4.95e-4 3.91e-4 * * 1401 0 ^ + /H + * * 1.55e-4 3.51e-4 2.86e-4 1.84e-4 %c/(0) 1.17 1.01 1.00 1.15,7.63 * * O /H 3.31e-4 4.90e-4 6.23e-4 7.28e-4 4.55e-4 5.32e-4 6548+6584 N+/H+ 3.20e-7 3.06e-6 3.15e-5 1.27e-5 1.66e-7 8.29e-06 1750 N 2+/H + 2.20e-5 l.lOe-4 1.93e-4 6.25e-5 1.29e-5 3.44e-5 57 /im 1.90e-5 5.08e-5 4.26e-5 2.92e-5 * * 4.13e-5 1.03e-4 2.76e-4 8.70e-5 * * 1486 N 3+/H + * * 1.09e-4 6.77e-6 1.33e-5 6.06e-6 1240 N ^+/H + ** * 2.35e-5 * 6.24e-6 ic/(N ) ** 1.00 * * 1.00 N/H 3.41e-5 1.52e-4 3.34e-4 2.54e-4 6.10e-5 8.48e-5 2326 C + /H + * * 1.54e-5 3.73e-5 * 2.63e-5 1908 C^+/H+ 1.18e-4 1.29e-4 1.92e-4 2.40e-4 8.69e-5 2.04e-4 1550 C^+/H+ * ** 1.33e-4 6.55e-5 2.21e-5 %c/(C) 1.01 1.03 * * 1.02 * C/H 1.39e-4 1.52e-4 2.12e-4 7.50e-4 2.12e-4 2.60e-4 3868+3967 Ne^+/H+ 6.61e-5 1.316-4 1.26e-4 2.03e-5 1.59e-5 5.57e-5 15.5 fj.m 3.71e-5: 1.47e-4: 9.61e-5: 5.75e-5: * 5.18e-5: 1601 N e^+/H + ** 7.35e-5 5.10e-5 5.35e-5 5.00e-5 2423 * * 1.04e-5: ** 1.45e-5; 4725 9.18e-6: ** 5.95e-5 4.47e-5 5.55e-5 1574 Ne'^+/H+ ** 1.60e-3: * 3.85e-6 2.05e-4: 3426 * * * 1.80e-5 * 9.57e-6 zc/(Ne) 1.18 1.04 1.00 1.00 1.00 1.00 Ne/H 7.80e-5 1.36e-4 2.00e-4 9.36e-5 6.89e-5 1.18e-4 7135 Ar^+/H+ 4.13e-7 1.37e-6 3.00e-6 8.87e-7 3.39e-7 8.41e-8 4711+4740 Ar^+/H+ 5.48e-7 6.31e-7 1.49e-7 5.88e-7 7.82e-7 6.08e-7 7005 Ar^+/H+ 7.16e-9 * * 1.15e-7 2.13e-7 9.57e-8 ïc/(A r) 1.01 1.02 1.16 1.08 1.00 1.11 Ar/H 9.79e-7 2.04e-6 3.65e-6 1.72e-6 1.33e-6 8.73e-7 4069+4076 S+/H+ 3.37e-9: 8.91e-8: 1.3e-6 2.89e-7 * 1.17e-7 6716+6730 1.45e-8 1.44e-7 1.24e-6: 2.83e-07: 1.50e-8: 2.73e-7: 6312 S^+/H+ 6.62e-7 3.50e-6 1.05e-5 1.93e-6 7.92e-7 1.54e-6 icf{S) 3.52 2.30 1.74 2.24 4.57 2.01 S/H 2.38e-6 8.3&&^ 2.04e-5 4.96e-6 3.69e-6 3.64e-6 5517+5537 Cl^+/H+ 2.44e-8 7.59e-8 1.31e-7 5.02e-8 1.76e-8 5.24e-8 2C/(C1) 3.59 2.40 1.95 2.57 4.66 2.37 Cl/H 8.77e-8 1.82e-7 2.55e-7 1.29e-7 8.20e-8 1.24e-7 2.5. Elemental abundances from CELs 66

Table 2.7: (b.)

NGC 3132 NGC 2440 NGC 6302 IC4406 IC4191 M y Cn 18

4471 He+/H+ 0.1229 0.0616 0.0850 0.1163 0.0971 0.1057 5876 He+/H+ 0.1151 * 0.0775 0.1180 0.1126 0.0984 6678 He+/H+ 0.1103 * 0.0118 0.1187 0.1100 0.0862 Avg. He+/H+ 0.1157 0.0616 0.0658 0.1178 0.1090 0.0974 4686 He^+/H+ 0.0031 0.0592 0.0696 0.0067 0.0110 0.0004 H e/H 0.119 0.121 0.135 0.124 0.120 0.098 2470 0 + /H + ** 1.67e-5 2.03e-4 ** 3727 3.03e-4 * 8.42e-6 1.79e-4 2.26e-5 2.08e-4 7320+7730 3.00e-4: * 2.03e-5: 1.70e-04: 2.81e-5: 4.78e-4: 1663 0^ + /H + 2.35e-4 5.86e-5 9.35e-5 3.86e-4 * * 4931 2.71e-4: 2.34e-4: 2.66e-4: 3.67e-4: 4.05e-4: 3.76e-4: 4959+5007 3.39e-4 9.71e-5 9.18e-5 3.80e-4 5.46e-4 3.54e-4 52+ 88 2.94e-4 * 4.20e-5 3.90e-4 1.90e-4 * 3.48e-4 * 9.90e-5 7.78e-4 8.36e-4 * 1401 0^ + /H + * 8.68e-5 8.22e-5 ** * ic /( 0 ) 1.02 * 1.37 1.04 1.07 O/H 6.55e-4 2.46e-4 2.50e-4 5.81e-4 6.08e-4 5.64e-4 6548+6584 N + /H + 1.29e-4 * 4.48e-5 5.90e-5 9.64e-6 8.11e-5 1750 N 2+/H + 9.95e-5 7.10e-5 1.03e-4 l.lO e-4 * * 57 8.68e-5 * 5.20e-5 9.05e-5 2.45e-5 * 1.33e-4 * 3.13e-4 2.63e-4 1.34e-4 * 1486 N 3+/H + * 2.71e-5 7.80e-5 *** 1240 N^+/H + * 1.22e-5 9.45e-5 * * * ic/(N ) * 1.00 1.00 * * N/H 2.36e-4 1.81e-4 3.34e-4 2.11e-4 3.89e-5 2.20e-4 2326 C + /H + ** * 1.34e-4 ** 1908 C^+/H+ 1.50e-4 1.13e-4 3.08e-5 2.11e-4 ** 1550 C^+/H + 1.67e-5 3.44e-5 1.43e-5 1.81e-5 * * 2C/(C) 1.89 1.28 1.72 1.00 ** C/H 3.15e-4 2.37e-4 7.74e-5 3.63e-4 ** 3868+3967 Ne^+/H+ 1.58e-4 3.36e-5 2.65e-5 1.39e-4 1.47e-4 7.74e-5 15.5 fim. 2.41e-5: 3.75e-5; 1.81e-5: * ** 1601 N e^+/H + ** 3.53e-5 * ** 2423 * 2.03e-5 7.59e-5: 1.Ole-5: ** 4724 * * 2.37e-5 * 7.21e-5: * 1574 N e^+/H + * 9.38e-5: * ** * 3426 * 5.59e-5 2.00e-5 ** zc/(Ne) 1.96 1.00 1.00 1.53 1.11 1.59 N e/H 3.10e-4 l.lO e-4 7.60e-5 2.13e-4 1.64e-4 1.23e-4 7135 Ar^+/H+ 2.37e-6 * 6.51e-7 1.64e-6 1.66e-6 1.88e-6 4711+4740 Ar^+/H+ 6.12e-8 5.12e-7 7.75e-7 1.80e-7 5.99e-7 * 7005 Ar^+/H+ ** 4.67e-7 * 7.37e-8 * zc/(Ar) 2.21 * 1.16 1.39 1.33 * A r/H 5.37e-6 * 2.19e-6 2.53e-6 3.10e-6 4069+4076 S+/H+ 2.43e-6: * 8.20e-7: 4.20e-7: 3.45e-7: 1.43e-6: 6716+6730 3.21e-6 * 1.04e-6 4.77e-7 5.59e-7 8.64e-7 6312 S^+/H+ 6.80e-6 * 1.57e-6 1.38e-6 5.49e-6 1.35e-5 ic/(S ) 1.06 * 2.14 1.14 2.10 1.20 S/H 1.06e-5 * 5.59e-6 2.12e-6 1.27e-5 4.43 5517+5537 Cl^+/H+ 1.48e-7 * 2.76e-8 9.37e-8 1.02e-7 2.09e-7 ic/(C l) 1.56 * 3.56 1.54 2.31 1.17 C l/H 2.31e-7 * 9.83e-8 1.44e-7 2.36e-7 2.45e-7

Values followed by have not been used in order to derive total abundances. 2.5. Elemental abundances from CELs 67

Table 2.8: He and CEL heavy-element abundances in Magellanic Cloud PN.“

SMC N87 LM CN66 LM CN141 4471 He+/H+ 0.0965 0.0572 0.1058 5876 H e+/H + * 0.0375 * 6678 He+/H+ 0.0981 0.0435 0.1119 Avg. He+/H+ 0.0972 0.0440 0.1085 4686 He:+/H+ 0.0000 0.0607 0.0003 H e/H 0.097 0.105 0.109 2470 0 + /H + *** 3727 1.84e-6 5.95e-6 4.48e-6 7320+7330 1.38e-5: 2.15e-5: 3.25e-5: 1663 0^ + /H + * 4.58e-5: * 4931 7.94e-5; * 1.04e-4: 4959 1.06e-4 7.81e-5 1.93e-4 1401 0 ^ + /H + *** ic /( 0 ) 1.00 3.77 1.00 O/H 1.08e-4 3.17e-4 1.97e-4 6548+6584 N+/H+ 6.10e-7 7.44e-6 2.03e-6 1750 N 2+/H + * 1.84e-5 * 1486 N 3+/H + * 4.46e-5 * 1240 N^+/H + * 2.70e-5 2.81e-5: ic/(N ) 58.7 1.00 43.97 N /H 3.58e-5 9.74e-5 8.93e-5 2326 C+/H+ 3.59e-5 * 3.02e-5 1908 C^+/H+ 3.22e-4 5.92e-6 1.49e-4 1550 C^+/H+ 1.92e-5 7.77e-6 2.04e-5 ic/(C ) 1.00 2.40 1.00 C/H 3.77e-4 3.29e-5 2.00e-4 3868+3967 Ne^+/H+ 1.04e-5 1.74e-5 2.35e-5 1601 Ne^+/H+ * ** 2423 * 1.39e-5 * 4724 * 1.90e-5 * 1574 Ne^+/H+ * * * 3426 * 8.24e-6 * icfÇNe) 1.02 1.00 1.02 Ne/H 1.06e-5 4.21e-5 2.40e-5 7135 Ar^+/H+ 1.76e-7 4.28e-7 6.39e-7 4711+4740 Ar^+/H+ 3.13e-8 5.43e-7 1.03e-7 7005 Ar^+/H+ * 1.57e-7 * zc/(Ar) 1.02 1.13 1.02 Ar/H 2.11e-7 1.27e-6 7.59e-7 4069 S+ /H + 4.10e-8: * 1.21e-7: 6716+6730 2.85e-8 2.43e-7 1.48e-7 6312 S^+/H+ * 7.69e-7 1.50e-6 TC/(S) * 2.11 2.47 S/H * 4.30e-6 4.06e-6

Values followed by have not been used in order to derive total abundances. 2.5. Elemental abundances from CELs 68

the statistical balance equation for each ion and yields level populations and line emissivities for a specified (T^, appropriate to the zones in a nebula where the ions are expected to exist. The following expression was then used in order to convert the observed line intensities to ionic abundance fractions (cf. Chapter 1, section 1.3.1);

iV(X"’+) _ /(A) N(H+) AijriiEii IÇÜ0) where /(A) are dereddened line fluxes, Aij are the radiative transition probabilities for 2— Hi is the fractional population of level 2, Eij is the excitation energy of

level 2 above j, aeff(H/3) is the effective recombination coefficient of H/? and hv/12 is the energy of an H/5 photon.

Generally when deriving CEL abundances, Te([N ll]) is assumed to represent the electron temperature appropriate for singly ionized species, while Te([0 III]) is used for higher excitation ions. However, as we showed in Section 2.4.2.2, in some cases the [N ll] and [O ll] nebular to auroral line ratios can be affected in ways that result in them being unreliable diagnostics of the temperature pertinent to the lower excitation zones of PN. Therefore, in our CEL abundance analysis we have adopted the following scheme: abundances of singly ionized species were derived using Te([N 11]) and N^iO ll), apart from the cases of NGC 3242, NGC 5882 and My Cn 18, where T^([0 ill]) was used instead; whenever Nq{0 II) was not available, Ng(S 11) was employed. For all doubly ionized species (C^"*", N^'*', 0^+, Ne^"'",

Cl^"*" and Ar^+), T^([0 III]) and A^e(Cl III) were used. For triply ionized species

(0^+, N^'*’, Ne^+, Ar^"*”), T^([0 Ill]) + 1000K with N q {K x iv ) were used, except for C^^, where Te([0 ill]) + 650K was used instead. Finally, ]A([0 III]) + 2270K, was used for all four-times ionizes species (N'^'*', Ne'^+, Ar“^+)—from ionization potential considerations as in Kingsburgh & Barlow (1994). Abundances of neutral species were not derived; it is assumed that fractions of neutrals relative to neutral hydrogen are equal to the abundance deduced for the same element when in ionic form, relative to H+. Therefore, a derivation of the total elemental abundance can be made using ionic fractions only. In Tables 2.7 2.5. Elemental abundances from CELs 69

and 2.8 the abundances for all observed species of C, N, O, Ne, S, Cl and Ar are presented, together with the adopted ionization correction factors (ICFs) and total elemental abundances, as derived using CELs only. The adopted fractions were obtained from the A3727 doublet only. The abundances derived from the [O ll] AA7320, 7330 lines are generally higher, as is evident from Table 2.6; the reason could be that a larger fraction of the line fluxes than of A3727 can be due to recombination excitation and/or the fact that the emission of AA7320, 7330 is biased towards higher density regions (cf. Section 2.4.2.2). Still, the ionic fractions deduced from A3727 may represent upper limits only, due to contributions from recombination.

In the case of we adopted ionic fractions obtained from the nebular [0 III] AA4959, 5007 lines only, in order to derive total O abundances. For comparison, in

Tables 2.7 and 2.8, we also present fractions derived from the O III] A1663

and [O III] A4931 lines. We further present abundances derived from the

far-IR [O III] 52- and 88-//m flne-structure lines for 8 nebulae, using the ISO LWS line fluxes published in Liu et al. (2001); two values per nebula are listed, since for comparison reasons we used both the low electron densities obtained from the 88 //m/52 //m line ratio and the higher obtained from our optical [Ar iv] and

[Cl III] density diagnostics (Table 2.4).

Similarly, we present ionic fractions from both the N III] A1750 line

and from the far-IR [N III] 57-//m line (using the ISO LWS 57-//m line fluxes of Liu et al. 2001); again abundance ratios from the latter line were derived adopting

both the low and high nebular electron densities just like for the [O III] far-IR lines.

Both the [0 III] and [N III] far-IR lines originate from atomic levels that have quite low critical densities, Ncr ~ (2-4) x 10^cm~^, lower than the average elec­ tron density of most PN in our sample. The ionic abundances derived from these lines are therefore acutely sensitive to our assumption of the actual density of the emitting medium. For example, adopting a density signiflcantly higher than 2.5. Elemental abundances from CELs 70

the Ncr of the collisional line will result in a proportional increase of the de­ duced abundance, since for the same flux to be emitted the presence of more ions is required, otherwise the emission line would be quenched. On the other

hand, the nebular [O III] AA4959, 5007 lines have much higher critical densities, A^cr X 10^cm~^, so that the abundances deduced from them are much less sensitive to the adopted nebular density; the same is true for the N III] A1750 line which has Ncr ~ H x 10® cm~^. The ionic abundance ratios and derived from far-IR lines are expected to be underestimated in the presence of density variations in nebulae, due to the low critical densities of the flne-structure levels from which the lines arise (Rubin 1989). The likelihood of systematic bias in abundance determinations from CELs becomes less when lines with fairly high critical densities are used, as long as Nq values in the nebula are well below those values. These expectations are confirmed by our observations. A comparison of the electron densities derived by the 88 //m/52 fim line ratio and the higher critical density [Cl III], [Ar iv] ratios in the eight PN common to our sample and Liu et al.’s (2001) ISO observations, shows that on average the latter diagnostics yield N^s about a factor of 6 higher than the far-IR ratio confirming the existence of density inhomogeneities in these nebulae (Table 2.4). In accord with the theoretical predictions of Rubin, the bias in the inferred ionic abundance ratios from the far-IR lines becomes less when the higher electron densities are adopted and the values returned are in good agreement, in most cases, with those derived by the [N III] A1750 and [O III] AA4959, 5007 lines (Table 2.7). A notable nebula is IC 4406, for which the and fractions derived from the far-IR lines agree with the corresponding fractions derived from the UV and optical lines in the low density case (see also the related discussion in Section 2.6.7.1).

We derived Ne^"'"/H'^ fractions from the optical [Ne III] AA3868, 3967 lines for all 15 PN, while for 9 of them we also determined this abundance ratio using the

IRAS [Ne III] 15.5-//m line fluxes by Pottasch et al. (1984). For the 15.5-//m line. 2.6. Ionic Abundances from ORLs 71

TVcr = 2 X lO^cm"^, and the resulting ionic abundance does not depend critically on the adopted TVg, given the inferred nebular densities listed in Table 2.4. In all these cases where IR collisionally excited lines were analyzed, we adopted the Te([0 lll])’s for abundance determinations. IR CELs have small excitation en­ ergies, Eex 1000 K, much smaller than UV and optical forbidden lines. Thus their emissivities have only a weak dependence on the adopted nebular electron temperature, very similar to that of Hi Balmer lines. Their intensities there­ fore relative to Yif3 are virtually insensitive to the assumed temperature, unless the emitting medium has 7^ 1000 K; however the electron temperatures of all nebulae in our sample are consistently much higher than that. Ne^'*‘/H+ ionic fractions were derived from the auroral [Ne iv] AA4724, 4726 lines in conjunction with those obtained from the A1601 line. Regarding the Ne^"^/H"*" ratio, whenever available the values obtained from [Ne v] A3426 were

preferred over those from the Ne V A1574 line. We adopted the abundance derived from the [S ll] AA6716, 6731 dou­ blet, rather than from the transauroral AA4068, 4076 lines, since the latter are potentially affected by recombination processes and density effects, as well (the and atomic levels of S''”, from which the AA6716, 6731 and AA4068, 4076 lines respectively arise, are directly analogous to the levels of O’*” from which the [O ll] nebular and auroral lines originate; especially, in terms of critical densities).

2.6 Ionic Abundances from ORLs

Since ionic abundances derived from optical recombination lines, relative to are almost completely insensitive to temperature and density variations, the stan­ dard [O III] temperature and mean [Cl III], [Ar iv] densities for each nebula, were adopted for all calculations. 2.6. Ionic Abundances from ORLs 72

2.6.1 Helium

The ionic and total He abundance fractions derived from the He I AA4471, 5876 and A6678, and He II A4686 recombination lines are given in Tables 2.7 and 2.8 for the galactic and Cloud PN, respectively. Case A recombination was assumed for the triplet lines A4471 and A5876, and Case B for the singlet A6678 line. The effective recombination coefficients, aefî(He^'^), were from Brocklehurst (1972).

The dereddened line strengths of the He I lines were first corrected for contributions to their upper levels by collisional excitation from the He*^2s^S metastable level, using the formulae of Kingdon & Ferland (1995a). If these corrections are not made, the resultant He abundance ratios would be systematically higher and in error. The expression which then yields the He fractions is

He‘+ A(Â) aeff(H/3) /(A) H+ 4861 «eff(Hei+) /(H/3)'

The He'^'/H’'' abundance values thus derived, were co-added, weighted accord­ ing to the dereddened intensity ratios of the three He I lines. In general, the agreement between the He'^/H'*' ratios derived from the three lines is within a few per cent. The He^‘*'/H'^ ratio was estimated from the He II A4686 line only, using the effective recombination coefficients of Storey & Hummer (1995). The total elemental He abundance relative to H is given by, He/H = He^/H'*' -f He^'^/H'*'.

2.6.2 C2+/H+, C3+/H+ and C^+/H+

We have detected recombination lines of C from all of the PN of our sample, apart from the carbon-poor nebula LMCN66. C II lines have been detected from all galactic PN, as well as from SMCN87 and LMCN141. C III lines were detected from the majority of them, excluding NGC 3132 and My Cn 18 only.

The strongest observed C II line is the A4267 (V 6 ) 3d-4f transition, which was consistently recorded with a high S/N ratio (> 45) in our high resolution deep spectra. We derived /VO' abundance ratios from it using the recently calculated effective recombination coefficients of Davey et al. (2000), which include 2.6. Ionic Abundances from ORLs 73

both radiative and dielectronic processes; abundance ratios derived using A4267 are case-insensitive. The A4267 C II line has been used in the past in abundance analyses of galactic PN (e.g. Rola & Stasihska 1994) with conflicting results. It has also been detected from a number of LMC and SMC PN (Barlow 1987, Meatheringham & Dopita 1991, and from SMCN42 only, by Vassiliadis et al. 1992). No detections of this line have been reported, however, for LMC N 141 and SMC N87, and carbon ORLs have not been used till now in abundance studies of Cloud PN. Due to the structure of the C'*’ ion, the configuration of the valence orbital gives rise to only one atomic term, compared to three atomic terms for N’*' and (e.g.

Kuhn 1969, Allen 1972). As a result there are fewer C II recombination lines than of O II or N II, so they are of greater intensity. As an observational consequence of this fact, in several PN of our sample we have also detected C II recombination lines from higher principal quantum numbers, originating from states above the 4f ^F° level. In Table 2.9 we compare the observed intensities of these lines of high excitation energy, normalized such that /(A4267) = 1.00, against the predictions of recombination theory from Davey et al. (2000). This permits us to check whether the 4f ^F°-ng transitions which populate the upper level of 3d-4f A4267 can be safely attributed to recombination only, or whether unidentified processes contribute as well. This is of importance in the light of results from this and previous works (e.g. Kaler 1986, Rola & Stasihska 1994, Liu et al. 1995, Liu et al. 2000), according to which the A4267 line abundances in PN are often significantly higher than those derived from the collisionally excited C III] A1908 line. This fact had some times been attributed in the past to erroneous recombination coefficients, inaccurate line detections, or blending of the A4267 transition with a line from an unknown ionic species. In the case of NGC 3242 the agreement between observations and theory is ex­ cellent for all detected lines. The 4f^F°-7g^G A5342 line is blended with a feature identified as [Kr iv] A5345.9 (see also HAF99, their Table 6) and its intensity was 2.6. Ionic Abundances from ORLs 74

retrieved though Gaussian profile fitting. In NGC 5315 the agreement is very good for both detected lines of high excitation. In NGC 5882 the A6462 line is stronger by 49 per cent and the A4802 line weaker by 19 per cent than predicted, relative to A4267. In IC4191, as measured on the fixed-slit spectrum, the former line is stronger than predicted by 40 per cent, while the A5342 line is within 8 per cent of the predicted value. It should be noted that apart from A4802, the other high-level

G II lines are covered in our lower resolution 4.5 Â FWHM spectra only and the modest discrepancies in NGC 5882 and IG4191 for the A6462 line are within the estimated error margin. In the above cases the A4802 line was deblended from

N II A4803.3 via Gaussian line fitting. In conclusion, the good consistency among the observed and predicted relative intensities of C II 4f^F°-ng^G transitions in this PN sample suggests that there is no competing mechanism other than recombination that could contribute to the excitation of the 3d-4f A4267 line. This fact along with the high S/N ratio

CCD detections of this line in the present PN sample, indicate that C II A4267 is a reliable abundance diagnostic.

A similar conclusion can be drawn for the C II recombination lines found in the spectrum of the Orion Nebula (M42). In Table 2.9 we extend the above comparison to 4f-ng C II transitions in M42, whose intensities were presented recently by Baldwin et al. (2000). These authors did not mark the listed lines as

C II transitions, but left them unidentified instead (see their table 1). However, the measured wavelengths and our examination of their relative intensities leave no doubt as to their identity. The agreement with theory is excellent in this case as well, establishing beyond reasonable doubt the interpretation of recombination excitation for these lines. The C^'^/H’*' abundance ratios were derived from the A4187 (V 18) and A4650 (VI) lines, using the effective radiative and dielectronic recombination coefficients of Péquignot, Petitjean & Boisson (1991) and Nussbaumer & Storey (1984), re­ spectively; they are insensitive to the assumption of Case A or B. 2.6. Ionic Abundances from ORLs 75

Table 2.9: High-excitation C ll recombination lines.

A o ( Â ) Trans. NGC 3242 NGC 5315 N G C 5 8 8 2 IG4191 Orion (M 42) Theory

lo b s lo b s lo b s 7obs lo b s 7pred

4267.15 3 d -4 f 1.000 1.000 1.000 1.000 1.000 1.000

6 2 5 8 . 7 8 4p-5d * * * * * 0.012

6151.43 4d -6 f 0.030 * * *10.055“ 0.038 0.040

6461.95 4f-6g 0.102 0.095 0.153 0.166|0.144 0.087 0.103 ** 5 3 4 2 . 3 8 4f-7g 0.065 *10.056 0.049 0.053 * 4 8 0 2 . 2 3 4f-8g 0.031 0.020 0.025 0.034 0.031

“ The value before the dash is for the whole nebula; those after are from a fixed-slit.

In the PN NGC 2022, NGC 3242, NGC 5882 and NGC 6818, the G ill (V16) triplet at A4069 is also detected. This multiplet is primarily excited by radiative recombination and is usually seriously blended with [S ll] and O II V 10 lines in medium resolution spectra; its overall intensity was retrieved via multiple Gaussian fitting as described in Section 2.3. As was noted by Liu (1998), the observed relative intensities of G III V 16, V 18 and V 1 multiplets are not in accord with theoretical predictions; in an analysis of NGG4361, that author found that the

abundance ratio derived from the three different G III multiplets spans a range of 0.4 dex, which he attributed to probable uncertainties in the effective recombination coefficients. The four PN for which all V 16, V 18 and V 1 multiplets have been detected, offer a further testing ground for the reliability of current ORL

G III recombination coefficients. An instructive case is that of the high excitation nebula NGC 2022. As is appar­ ent from our high resolution FWHM 1.5 Â spectra of this object, the transauroral

[S ll] AA4068, 4075 lines are virtually absent and the observed intensity of C III V 16 is most accurately recorded, since it is less affected by blending effects. In this nebula, the observed G III AA4069, 4187 and A4650 multiplets have intensity ratios of 1.80 : 0.34 : 1.00 (cf. Table 2.13), versus the theoretical 0.59 : 0.21 ; 1.00 calculated using the effective radiative and dielectronic recombination coefficients 2.6. Ionic Abundances from ORLs 76

of Péquignot, Petitjean k. Boisson (1991) and Nussbaumer k Storey (1984), re­ spectively. In the case of NGC 5882, the observed multiplet intensities show ratios of 2.38 : 0.22 : 1.00. Similarly for IC4191, where C III A4069 is not detected, the observed intensities of A4187 (V18) and A4650 (VI) have a ratio of 0.20 : 1.00 (Table 2.13). We see that the theoretical predictions for the relative intensities of A4187 and A4650 appear to be more secure than that for A4069; it would seem that the effective recombination coefficient for the A4069 V 16 multiplet has been underestimated by a factor of ~2-4, or else that it is blended with an unknown line from an ion of similar excitation. As a result of the above discussion, in our ORL abundance analysis for we have discarded values derived from the A4069 triplet; instead, we have adopted

abundance ratios as derived from an intensity weighted mean of the C III A4187 and A4650 lines.

Regarding we have derived /Li'^ fractions from the detected C IV A4658 line for a number of galactic PN, using the Case A effective recombination coefficients of Péquignot, Petitjean k Boisson (1991). The ionic and total C abundances derived from recombination lines are pre­ sented in Tables 2.13 and 2.16 for the galactic and Cloud PN, respectively.

2.6.3 N2+/H+, N^+/H+ and N^+/H+

All the abundances were derived using effective recombination coefficients from Escalante k Victor (1990), assuming Case A for singlets and Case B for triplets.

In Table 2.10 we present our N II results for all galactic PN; nitrogen ORLs were not detected from any of the three Cloud PN. Apart from the abundances derived from individual transitions, we also list in Table 2.10 the mean abundances derived for each PN, by weighting each result according to the predicted intensity of individual transitions. For all galactic PN apart from NGC 5315 and My Cn 18, the A4379 (V 18) 2.6. Ionic Abundances from ORLs 77

Table 2.10: (a .) Recombination line N^"""/H abundances in PN. N2+ N2+ N^+ N2+ Ao M ult. Jobs lohs lohs lohs H+ H+ H+ H+ (Â) (10-4) (10-4) (10-4) (1 0 -4 )

NGC 3242 NGC 5882 NGC 5315 NGC 3918 V 3 3s^ P ° - 3 p 5666.63 V3 .03213 2.34 * ***** 5676.02 V3 .01336 2.20 .01885 3.10 *** * 5679.56 V 3 .04278 1.68 .06036 2.36 ** * * V 5 3s^ P ° - 3 p ^ P 4601.48 V5 .004988 1.19 * ** * * * 4607.16 V5 * ** * ** * * 4621.39 V5 * ** ** *** 4630.54 V5 .01787 1.43 * ** * * * 3 d -4 f 4041.31 V39b .01276 0.88 .02727 1.79 .05729 3.71 .01338 0.94 4237.05 V48b .01728 1.71 .01393 1.38 .01658 1.64 .01630 1.65 4241.78 V48a * + * * .04215 4.16 .01628 1.75 4530.41 V58b * * .01795 1.75 * * .01476 1.55 4678.14 V61b * * * * .03123 4.80 ** S u m .141 1.71 .138 2.24 .147 3.43 .0607 1.42

NGC 6302 IC 4191 NGC 6818 NGC 2022 V 3 3s^ P'’- 3 p 5666.63 V3 * * .09206 6.90 **** 5676.02 V3 * * .04003 6.58 .03215 5.29 .02877 4.73 5679.56 V3 * * .12817 5.02 .10295 4.03 .09213 3.61 V 5 3s^ P ° - 3 p ^P 4630.54 V5 .07433 5.55 .05494 4.50 * * * * V 2 8 3s ^ P °-3p ^ D ° 5927.81 V28 .02049 8.26 ** ** ** 5931.78 V28 .03605 6.47 .04601 5.73 * * ** 3 d -4 f 4041.31 V39b * * .06388 4.26 * * .05798 4.27 S u m .131 5.64 .425 5.50 .135 4.27 .179 4.38 2.6. Ionic Abundances from ORLs 78

Table 2.10: (b.) Recombination line + in PN. N2 + N2+ N2+ N2+ Ao Mult. lo h s 7obs lo b s 7obs H+ H+ H+ H+ (Â) (1 0 -') (1 0 -') (1 0 -') (1 0 - ')

NGC 3132 NGC 2440 IC 4406 My Cn 18 V 3 3s^ P ° - 3 p ^D 5676.02 V3 .02495 4.10 * * **** 5679.56 V3 .07989 3.13 * * * *** V 5 3s^P°-3p ^P 4607.16 V5 ** ** * * .08417 26.73 4621.39 V5 * * * *** .07307 23.28 4630.54 V5 .04107 3.40 * * .2318 1.90 .2099 17.86 3 d -4 f 4041.31 V39b * * .06184 4.28 * * * * 4237.05 V48b ** .04951 4.90 * *** 4241.78 V48a ** .05579 5.89 * ** 4530.41 V58b ** .04152 4.28 *** * S um .146 3.40 .209 4.77 .232 1.90 .283 20.35

N III recombination line has been detected. We derived abundances from it using the effective radiative and dielectronic recombination coefficients of Péquignot, Petitjean & Boisson (1991) and Nussbaumer & Storey (1984), respec­ tively. We have also detected the N IV 5g-6h A4606 line from the high excitation PN NGC 3918, NGC 2022, NGC 6818, NGC 2440 and NGC 6302; we used it to derive abundances, employing effective recombination coefficients as in the case of N III A4379.

2.6.4 Q2+/H+

These observations present us with one the most extensive records of O II recom­ bination spectra in gaseous nebulae thus far, obtained from PN possessing a wide range of physical conditions. For the first time also, O II lines have been detected and measured from Magellanic Cloud objects (PN SMC N87, LMC N66, and LMC N141). This allowed us to obtain reliable ORL abundances for all the PN in our sample and allowed a comprehensive investigation on the occurrence of the discrepancy between ORL O^"*" abundances and those derived in the usual manner 2.6. Ionic Abundances from ORLs 79

from CEL lines of 0^"*". Furthermore, having such an extensive inventory of O II recombination line intensities, we are able to perform a thorough comparison with

current theoretical predictions of the O II recombination spectrum. In Table 2.11 we present ORL ionic abundance ratios for the whole PN sample. Effective recombination coefficients are from Storey (1994) for 3s-3p transitions (LS'-coupling) and from LSBC for 3p-3d and 3d-4f transitions (inter­ mediate coupling), assuming Case A for doublet and Case B for quartet lines. In several PN we have also detected lines arising from doubly excited spectral terms such as the multiplet V 15 3s'^D-3p^ ^F° at 4590Â (cf. Appendix). The excitation of this multiplet is dominated by dielectronic recombination, but we have not derived ionic abundance ratios from it since the existing recombination coefficients are not of the desired accuracy (P. J. Storey, private communication).

Table 2.11: (a.) Recombination line O /H abundances of PN

0^+ Ao(Â) Mult fobs (10-") Lbs (10-") H+ H+ NGC 3242 NGC 5882 4638.86 VI .1692 15.86: .1690 16.21: 4641.81 VI .2152 8.00 .2522 9.59 4649.13 V I .2231 4.36 .3447 6.89 4650.84 VI .09785 9.18 .1141 10.95 4661.63 VI .08697 6.38 .1188 8.92 4673.73 VI .03040 13.11: .01878 9.10 4676.24 VI .05740 5.01 .09064 8.10 V 1 3s "^P-3p 0.681 6.02 0.939 8.29 4317.14 V2 .04539 5.39 .06375 8.43 4319.63 V2 .02270 2.74 .03331 4.08 4325.76 V2 .03107 20.28: ** 4345.56 V2 .08813 11.09: .09113 14.64 4349.43 V2 .07123 3.71 .09113 4.82 4366.89 V2 .09092 10.61: * * o 00

CM * 50 N CM * 50 50 00 b- O 00 (0 CD 50 CM 0 5 0 0 CM 50 00 00 (05 50 (05 D- CD 0 0 o 50 (05 T f CM(05 (05 50 00 I> CM 0 5 05 00 CO 00 0 5 50 00 (05 (05 CD 50 0 0 (05 o6 cô 50 05 CÔCD N CD (05 CM o d o rH c ô CD h-: O d d d d rH rH 2 rH tH

o (05 CD CM CO o 50 o CM lO CM T-l 00 00 N (05 50 CO (05 CM 50 0 0 (05 00 4* 50 (05 (05 50 5 0 l> (05 co 0 5 00 CO CM CD o 00 CM 0 0 50 T f I > CD CM (05 OO (05 00 CD 00 50 00 O OO 50 CM N o 00 CD 50 50 50 15- CO CD CM 1 ^ 50 I > oo oo 00 CM CM CM 50 (05 tr- CM M 00 05 CD OO CD CM CO 50 CM CM CM CD 50 CD CM (05 q o O O CD q o * d * O d O * o O o O 6 O O O O O d O * * P * * p

00 o t - 0 0 50 (05 CM rH O CD 00 (05 b - CD O CD (D 5 0 f - (05 00 CM CD o (05 !>■ 0 0 50 00 (05(05 t - 00 O CO CO (05 CM CM (05 CM CD CM 50 50 0 5 00 00 t - 0 0 OO OC) CM O CD d 5 0 (05 CO CO CD d d d 5 0 d 50 50 d O d d c ô c6 N cô d (05 d o6 oô oô

CO CM 0 5 50 0 0 CD 0 0 0 0 CD (05 CM 50 CO CD t ^ 0 0 00 CO CO b - 5 0 5 0 CO (05 b - 5050 CD500050 OO 00 00CO CMO O CO CD o CO CM 50 CD CO (D 50 CD 50 (05 50 CM CD 50 CO b - CO CM 50 CD 5 0 15- o h - CM (D 15- N 0 0 (05 50 CD 00 00 OO 00 (05 00 § CO CO 0 0 CD CM CO CM CM 50 CM CO co CM CM CO CM OO O O CO 50 O o CD O CD ( 3 d d CM d C5 OO d O q q q q q O (D (D CD CD O CD CD O o CD CD o O CD CD CD

ci rO iS OOOO^pOoO oo OO 0 0 CD oo CO 1—4 1-4 1-4 1-4 1-4 t -4 CM CM CM CM CM Tf^TfCDCDWCDCDf^lOCyr^LO 50 50 CD S > > > > > > > > > >>>>>> > > > > > >>>>>>>>>>>>> > > > > > & Ig Oh 'Ü T3 T3 13 w 00 e o a a ? ” ” O h O h g w, (05 CD CD CO CO CD CM CD o CO CM 00 CM CD CO CO o 50 (05 50 50 CO CD CO (05 CM 50 0 0 CO CO (05 50 CO CO CO Tf b- w W OO 00 0 0 (05 d (CO 0 0 CO 50 CM b- CM 00 0 0 50 05 CM 50 b- (05 CD 0 0 P CM b- 0 0 CM 50 CM d CM 50 oô 50 CM O d CM cô d d 05 d d o d d Tf 00 cô d d 50 d t > CM cô 50 oô CM cô 1> d cô M lO CD 15- N 0 0 (05 CM CM CO 50 50 CD rH CM 05 CD CM CM 0 0 0 0 00 D- t - N 00 0 0 0 0 0 0 (05 (05 (05 o CD 50 50 %) CD O o o CD CD 0 0 (05 (05 (D o o CM CM CM CM CM CM CM CM CM CM CO CO CO CO CO K > > T f )> T f T f Tfl 2.6. Ionic Abundances from ORLs 81

0^+ Ao(Â) Mult Iq6s lobs ■ (1 0 - ') H+ 4366.53 V75a * * .06046 7.13 4466.42 V86b .02597 23.78: .03047 26.73: 4477.90 V88 .01448 15.72 * * 4489.49 V86b .00551 7.86 ■¥ * 4491.23 V86a .02351 16.01 .02034 13.27 4609.44 V92a .04394 7.26 .06087 9.64 4669.27 V89b .01193 27.64: ** 3d-4f 0.493 8.40 0.569 9.67

Adopted abundance 6.28 9.70

Table 2.11: (b.) Recombination line 0 ^ ^ /H abundances of PN.

q 2 + Ao(Â) Mult lobs (1 0 -') H+ NGC 5315 4638.86 VI .1740 16.54: 4641.81 V I .2367 9.86 4649.13 V I .3574 7.08 4650.84 VI .09002 8.56 4661.63 VI .08969 6.67 4673.73 V I .01367 6.56 4676.24 VI .07894 6.99 VI 3s^P-3p^D° 0.866 7.58 4317.14 V2 .06940 8.23 4319.63 V2 .04911 5.94 4325.76 V2 .01911 11.82: 4345.56 V2 .06347 8.34 4349.43 V2 .07439 9.77 4366.89 V2 .05887 6.87 V 2 3s ^P-3p "^P° 0.315 6.06 4414.90 V5 .07206 13.00 4416.97 V5 .06200 20.16 4452.37 V5 .01455 29.77: V 5 3s 2p-3p ^D° 0.134 15.55 4069.62 VIO .4263 16.45: 4072.16 VIO .1879 7.78 4075.86 VIO .1412 4.05 4085.11 VIO .04371 9.69 VIO 3p^D°-3d^F 0.329 5.58 4083.90 V48b .02411 7.80 4087.15 V48c .03185 10.85 4089.29 V48a .08986 8.30 4275.55 V67a .10053 8.11 4303.82 V53a .03307 6.50 2.6. Ionic Abundances from ORLs 82

A o ( Â ) Mult lob s 4609.44 V92a .04763 7.79 3 d -4 f 0.332 8.08

Adopted abundance 8.57

NGC 3918 4638.86 V I .2200 20.37; 4641.81 VI .2168 7.96 4649.13 VI .2108 4.07 4650.84 VI .04146 3.84 4661.63 VI .07432 5.39 4673.73 VI .03101 14.50: 4676.24 VI .06560 5.66 V 1 3s ^P-3p 0.609 5.28 4317.14 V2 .04170 4.93 4319.63 V2 .01701 2.04 4325.76 V2 .01414 9.17: 4349.43 V2 .06383 3.31 V 2 3s"‘P -3 p ^ P ° 0.123 3.39 4414.90 V5 .05178 8.88 4416.97 V5 .02900 8.96 4452.37 V5 .06525 100.: V5 3s^P-3p^D° 0.081 8.91 4121.46 V19 .09197 33.08: 4132.80 V19 .01800 3.32 4153.30 V19 .02600 3.35 4156.53 V19 .03701 29.94: 4169.22 V19 .004548 1.72 V19 3p^P°-3d^P 0.049 3.07

Adopted abundance 5.36

NGC 2022 4638.86 VI .1051 9.41: 4641.81 VI .2538 9.01: 4649.13 VI .3674 6.86 4650.84 VI .08923 8.00 4661.63 VI .1093 7.67 4676.24 VI .07717 6.44 V 1 3s ^P-3p ^D° 1.002 7.69 4414.90 V5 .1247 19.28 V5 3s2p-3p^D° 0.125 19.28 4072.16 VIO .3164 13.06 4075.86 VIO .2949 8.43 VIO 3p^D°-3d^F 0.611 10.32 4089.29 V48a .1726 16.95 4276.75 V67b .1856 13.79 4477.90 V88 .09358 106.7: 2.6. Ionic Abundances from ORLs 83

A o ( Â ) Mult lob s 4609.44 V92a .05776 10.04 3 d -4 f 0.416 14.75

Adopted abundance 13.01

NGC 6818 4638.86 VI .2160 19.833: 4641.81 VI .4956 18.039: 4649.13 V I .1687 3.229 4650.84 VI .02617 2.499 4661.63 VI .08505 6.113 V 1 3s ^P-3p 0.280 3.65 4366.89 V2 .09975: 11.52: V2 3s^P-3p^P° 0.099: 11.52: 4414.90 V5 .05944 9.90 4452.37 V5 .06432 105.: V5 3s2p-3p^D° 0.059 9.90 VIO 3p^D°-3d^F 0.468 6.41 4089.29 V48a .09842 9.37 3d-4f .09842 9.37

Adopted abundance 7.33

NGC 3132 4638.86 VI .1043 10.01 4641.81 VI .1706 6.49 4649.13 VI .1649 3.30 4650.84 VI .1226 11.77 4661.63 V I .1157 8.70 4676.24 VI .07618 6.82 VI 3s^P-3p^D° 0.754 7.06 4069.62 VIO .4086 15.82: 4072.16 VIO .1484 6.17 4075.86 VIO .3592 10.34 4085.11 VIO .04328 9.64 VIO 3p^D°-3d^F 0.551 8.70

Adopted abundance 8.15

NGC 2440 4638.86 VI .2740 24.0: 4641.81 VI .2051 7.14: 4649.13 V I .1266 2.32 4650.84 VI .05099 4.48 4661.63 VI .03461 2.38 4676.24 VI .07103 5.81: VI 3s^P-3p^D° 0.212 2.63 4414.90 V5 .04715 6.91 4416.97 V5 .03795 10.02 4452.37 V5 .08022 106.: 2.6. Ionic Abundances from ORLs 84

Ao(Â) Mult lobs

V5 3s^P-3p^D° 0.085 8.02 4072.16 VIO .1458 5.99 V 10 3p ^D°-3d 0.146 5.99 4121.46 V19 .03467 12.42: 4129.32 V19 .06667 102.: 4153.30 V19 .03327 4.28 V19 3p^P°-3d^P 0.033 4.28

Adopted abundance 5.23

IC 4406 4638.49 VI .1817 17.34; 4641.44 VI .2627 9.94 4649.13 VI .2585 5.14 4650.84 VI .09590 9.16 4661.63 VI .09884 7.39 4676.24 VI .04808 4.28 VI 3s^P-3p^D° 0.764 6.84 4069.62 VIO .4106 15.90: 4072.16 VIO .3776 15.70: 4075.86 VIO .2530 7.28 VIO 3p^D°-3d^F 0.253 7.28

Adopted abundance 7.06

SMC N87 4649.13 VI .07814 1.51 4661.63 VI .02428 1.77 4676.24 VI .02481 2.15 V I 3s^P-3p^D° 0.127 1.72

Adopted abundance 1.72 LMC N66 4641.81 VI .3072 10.25 4649.13 VI .3286 5.77 4650.84 VI .08736 7.35 4661.63 VI .1316 8.67 V I 3s^P-3p^D° 0.855 8.53

Adopted abundance 8.53 LMC N141 4638.86 VI .02549 2.38 4641.81 VI .1202 4.45 4649.13 VI .1267 2.47 4650.84 VI .08360 7.81 4661.63 V I .05177 3.79 4676.24 VI .04646 4.04 V I 3s^P-3p'^D° 0.0757 3.08 4069.62 VIO .1879 7.27 4072.16 VIO .1238 5.15 2.6. Ionic Abundances from ORLs 85

'^o(A) Mult lobs 1 (10-^) 4075.86 VIO .2536 7.30 VIO 3p^D°-3d^F 0.565 6.68

Adopted abundance 4.96

Table 2.11; (c.) Recombination line abundances of PN.

I 0 '+ Ao(Â) Mult lobs lobs y + (10-")

IC 4191 whole nebula fixed slit 4638.86 VI .1912 17.99: .2426 22.83: 4641.81 VI .3218 12.01 .2020 12.17 4649.13 V I .5439 10.67 .5757 11.93 4650.84 VI .09970 9.38 .1168 14.49 4661.63 VI .1539 11.34 .1442 10.62 4676.24 VI .1226 10.75 .1084 9.29 VI 3s^P-3p'‘D° 1.242 10.95 1.147 10.09 4317.14 V2 * * .06755 8.77 4319.63 V2 * * .05508 6.62 4345.56 V2 * * .1210 15.17 4349.43 V2 * * .2324 12.07 4366.89 V2 * * .1064 12.37 V 2 3 s^ P -3 p '‘P ° * * 0.582 11.07 4414.90 V5 .07566 13.15 .05461 9.49 4416.97 V5 .05307 16.62 .04438 13.90 4452.37 V5 .04022 63.3; * * V5 3s2p-3p^D° 0.129 14.40 0.099 11.07 4069.62 VIO .7056 27.25; .5998 23.17; 4072.16 VIO .3654 15.16 .3565 14.79 4075.86 VIO .2883 8.28 .3452 9.91 4085.11 VIO .05028 10.99 .02658 5.81 VIO 3p^D°-3d^F 0.704 11.09 0.728 11.47 4083.90 V48b .04179 13.81 .07012 23.17 4087.15 V48c .05799 18.56 .06847 21.91 4089.29 V48a .1475 13.94 .2059 19.09 4275.55 V67a .2096 16.15 .2434 18.76 4282.96 V67c .05873 12.61 * * 4303.82 V53a .09296 18.12 .06839 13.33 4466.42 V86b * * .03869; 35.86; 4609.44 V92a .08519 14.25 .06032 10.09 3 d -4 f 0.694 15.24 0.717 17.56

Adopted abundance 12.92 12.25 2.6. Ionic Abundances from ORLs 86

In Table 2.12 we present a comparison of the dereddened O II line intensities detected from all objects, against the predicted intensities from recombination theory. The comparison is relative to the strongest expected line within each multiplet for the 3-3 transitions, while for the collection of 3d-4f transitions it is relative to the strongest expected line at 4089.3 Â. The bracketed figures are the estimated absolute errors in the values, arising from the line fitting method only described previously in Section 2.3, and do not include any possible systematic errors (for instance, those arising from the calibration process or the corrections for atmospheric and interstellar extinction).

Table 2.12: Comparison of the observed and predicted relative intensities of planetary nebula O II lines.

Ao(Â) Mult Termj-Terniu gl~gu ïpred Io6s ^obs A p red

NGC 3242 raa-gp; 4638.86 VI 3s 4P-3p 4D* 2-4 0.21 0.76[.08] 3.6[0.4] 4641.81 VI 3s 4P-3p 4D* 4-6 0.53 0.97[.07] 1.8[0.1] 4649.13 VI 3s 4P-3p 4D* 6-8 1.00 1.00 1.0 4650.84 VI 3s 4P-3p 4D* 2-2 0.21 0.44[.05] 2.1[0.2] 4661.63 VI 3s 4P-3p 4D* 4-4 0.27 0.39[.04] 1.4[0.1] 4673.73 VI 3s 4P-3p 4D* 4-2 0.04 0.14[.01] 3.5[0.3] 4676.24 VI 3s 4P-3p 4D* 6-6 0.22 0.26[.02] 1.2[0.1] 4317.14 V2 3s 4P-3p 4P* 2-4 0.44 0.64[.08] 1.5[0.2] 4319.63 V2 3s 4P-3p 4P* 4-6 0.43 0.32[.06] 0.7[0.1] 4325.76 V2 3s 4P-3p 4P* 2-2 0.08 0.44[.07] 5.5[0.9] 4345.56 V2 3s 4P-3p 4P* 4-2 0.42 1.24[.25] 3.0[0.6] 4349.43 V2 3s 4P-3p 4P* 6-6 1.00 1.00 1.0 4366.89 V2 3s 4P-3p 4P* 6-4 0.45 1.28[.2l] 2.8[0.5] 4414.90 V5 3s 2P-3p 2D* 4-6 1.00 1.00 1.0 4416.97 V5 3s 2P-3p 2D* 2-4 0.56 0.95[.15] 1.7[0.3] 4452.37 V5 3s 2P-3p 2D* 4-4 0.11 0.51[.17] 4.6[1.5]

4069.89 VIO 3p 4D*-3d 4P 4-6 0.74 1.10[.22] 1.5[0.4] 4072.16 VIO 3p 4D *-3d 4P 6-8 0.69 1.29[.21] 1.9[0.3] 4075.86 VIO 3p 4D*-3d 4P 8-10 1.00 1.00 1.0 4078.84 VIO 3p 4D*-3d 4P 4-4 0.11 0.13[.03] 1.2[0.3] 4085.11 VIO 3p 4D*-3d 4P 6-6 0.13 0.21[.04] 1.6[0.3] 4092.93 VIO 3p 4D*-3d 4P 8-8 0.09 0.14[.04] 1.6[0.4] 4121.46 V19 3p 4P*-3d 4P 2-2 0.36 0.90[.59] 2.5[1.6] 4129.32 V19 3p 4P*-3d 4P 4-2 0.08 1.70[.24] 21.[3.0] 4132.80 V19 3p 4P *-3d 4P 2-4 0.70 0.69[.15] 1.0[0.2] 4153.30 V19 3p 4P*-3d 4P 4-6 1.00 1.00 1.0 4156.53 V19 3p 4P*-3d 4P 6-4 0.16 1.23 [.20] 7.7[1.3] 4169.22 V19 3p 4P*-3d 4P 6-6 0.34 0.41[.07] 1.2[0.2] 2.6. Ionic Abundances from ORLs 87

Ao(Â) Mult Termj-Terrriu g l-g u Ipred lofcs ^obs/^pred

4083.90 V48b 3d 4F-4f G4* 6-8 0.29 0.19[.03] 0.7[0.1] 4087.15 V48c 3d 4F-4f G3* 4-6 0.27 0.18[.03] 0.7[0.1] 4089.29 V48a 3d 4F-4f G5* 10-12 1.00 1.00 1.0 4275.55 V67a 3d 4D-4f F4* 8-10 0.65 0.28[.05] 0.4[0.1] 4276.75 V67b 3d 4D-4f F4* 8-10 0.31 0.25[.04] 0.8[0.1] 4277.43 V67c 3d 4D-4f F2* 2-4 0.23 0.15[.02] 0.7[0.1] 4282.96 V67c 3d 4D-4f F2* 4-6 0.15 0.08[.01] 0.5[0.1] 4283.73 V67c 3d 4D-4f F2* 4-4 0.10 0.09[.01] 0.9[0.1] 4285.69 V78b 3d 2F-4f F3* 6-8 0.19 O.llj.Ol] 0.6[0.1] 4288.82 V53c 3d 4P-4f Dl* 2-4 0.10 0.05[.02] 0.5[0.2] 4291.25 V55 3d 4P-4f G3* 6-8 0.16 0.15[.04] 0.9[0.2] 4292.21 V78c 3d 2F-4f F2* 6-6 0.09 0.08[.03] 0.9[0.3] 4294.78 V53b 3d 4P-4f D2* 4-6 0.29 0.28[.03] i.ojo.i] 4303.83 V53a 3d 4P-4f D3* 6-8 0.47 0.52[.09] 1.1[0.2] 4307.23 V53b 3d 4P -4f D2* 2-4 0.11 0.08 [.02] 0.7[0.2] 4313.44 V78a 3d 2F-4f F4* 8-10 0.12 0.08[.02] 0.7[0.2] 4353.59 V76c 3d 2F-4f G3* 6-8 0.10 0.09[.02] 0.9[0.2] 4357.25 V63a 3d 4D-4f D3* 6-8 0.06 0.16[.02] 2.7[0.3] 4466.42 V86b 3d 2P-4f D2* 4-6 0.10 0.23[.06] 2.3[0.6] 4477.90 V88 3d 2P-4f G3* 4-6 0.09 0.13[.03] 1.4[0.3] 4489.49 V86b 3d 2P-4f D2* 2-4 0.07 0.05[.02] 0.7[0.3] 4491.23 V86a 3d 2P-4f D3* 4-6 0.14 0.21[.02] 1.5[0.1] 4609.44 V92a 3d 2D-4f F4* 6-8 0.56 0.40[.23] 0.7[0.4] 4669.27 V89b 3d 2D-4f D2* 4-6 0.04 0.11[.02] 2.8[0.5] NGC 5882

4638.86 VI 3s 4P-3p 4D* 2-4 0.21 0.49[.05] 2.3[0.1] 4641.81 VI 3s 4P-3p 4D* 4-6 0.53 0.73[.05] 1.4[0.1] 4649.13 VI 3s 4P-3p 4D* 6-8 1.00 1.00 1.0 4650.84 VI 3s 4P-3p 4D* 2-2 0.21 0.33[.02] 1.6[0.1] 4661.63 VI 3s 4P-3p 4D* 4-4 0.27 0.35[.02] 1.3[0.1] 4673.73 VI 3s 4P-3p 4D* 4-2 0.04 0.05[.0l] 1.3[0.3] 4676.24 VI 3s 4P-3p 4D* 6-6 0.22 0.26[.02] 1.2[0.1] 4317.14 V2 3s 4P-3p 4P* 2-4 0.40 0.70[.20] 1.8[0.5] 4319.63 V2 3s 4P-3p 4P* 4-6 0.43 0.37[.13] 0.9[0.3] 4345.56 V2 3s 4P-3p 4P* 4-2 0.40 1.21[.50] 3.0[1.3] 4349.43 V2 3s 4P-3p 4P* 6-6 1.00 1.00 1.0 4414.90 V5 3s 2P-3p 2D* 4-6 1.00 1.00 1.0 4416.97 V5 3s 2P-3p 2D* 2-4 0.56 0.91[.22] 1.6[0.4]

4069.89 VIO 3p 4D*-3d 4F 4-6 0.74 0.49[.16] 0.7[0.2] 4072.16 VIO 3p 4D*-3d 4F 6-8 0.69 0.87[.14] 1.3 [0.2] 4075.86 VIO 3p 4D*-3d 4F 8-10 1.00 1.00 1.0 4085.11 VIO 3p 4D*-3d 4F 6-6 0.13 0.11[.04] 0.9[0.3] 4092.93 VIO 3p 4D*-3d 4F 8-8 0.09 0.08[.02] 0.9[0.2] 4121.46 V19 3p 4P*-3d 4P 2-2 0.36 0.30[.23] 0.8[0.6] 4132.80 V19 3p 4P*-3d 4P 2-4 0.70 0.56[.08] 0.8[0.1] 4153.30 V19 3p 4P*-3d 4P 4-6 1.00 1.00 1.0 4156.53 V19 3p 4P*-3d 4P 6-4 0.16 0.76[.05] 4.8[0.3] 4169.22 V19 3p 4P*-3d 4P 6-6 0.34 0.29[.06] 0.9[0.2] 2.6. Ionic Abundances from ORLs

Ao(Â) Mult Term;-Term u gi-gu Ipred Io6s lobs/Ipred

4110.78 V20 3p 4P*-3d 4D 4-2 0.27 0.39[.07] 1.4[0.3] 4119.22 V20 3p 4P*-3d 4D 6-8 1.00 1.00 1.0 4890.86 V28 3p 4S*-3d 4P 4-2 0.27 0.13[.07] 0.5[0.3] 4906.83 V28 3p 4S*-3d 4P 4-4 0.59 0.55[.17] 0.9[0.3] 4924.53 V28 3p 4S*-3d 4P 4-6 1.00 1.00 1.0

4083.90 V48b 3d 4F-4f G4* 6-8 0.29 0.15[.05] 0.5[0.2] 4087.15 V48c 3d 4F -4f G3* 4-6 0.27 0.25[.04] 0.9[0.1] 4089.29 V48a 3d 4F-4f G5* 10-12 1.00 1.00 1.0 4275.55 V67a 3d 4D-4f F4* 8-10 0.65 0.40[.07] 0.6[0.1] 4276.75 V67b 3d 4D-4f F3* 6-8 0.28 0.16[.07] 0.6[0.3] 4277.43 V67c 3d 4D-4f F2* 2-4 0.23 0.18[.04] 0.8[0.2] 4291.25 V55 3d 4P-4f G3* 6-8 0.16 0.15[.03] 0.9[0.2] 4292.21 V18c 3d 2F-4f F2* 6-6 0.09 0.05[.03] 0.6[0.4] 4294.78 V53b 3d 4P-4f D2* 4—6 0.27 0.11[.02] 0.4[0.1] 4303.83 V53a 3d 4P-4f D3* 6-8 0.47 0.39[.04] 0.8[0.1] 4353.59 V76c 3d 2F-4f G3* 6-8 0.10 o.ioj.oi] i.ojo.i] 4366.53 V75a 3d 2F -4f D3* 6-8 0.75 0.42[.05] 0.6[0.1] 4466.42 V86b 3d 2P -4f D2* 4-6 0.09 0.21[.05] 2.3[0.6] 4491.23 V86a 3d 2P-4f D3* 4-6 0.14 0.14[.03] 1.0(0.2] 4609.44 V92a 3d 2D -4f F4* 6-8 0.57 0.42[.08] 0.8[0.2]

NGC 5315

4638.86 VI 3s 4P-3p 4D* 2-4 0.21 0.49[.04] 2.3(0.2] 4641.81 VI 3s 4P-3p 4D* 4-6 0.53 0.66[.05] 1.3(0.1] 4649.13 VI 3s 4P-3p 4D* 6-8 1.00 1.00 1.0 4650.84 VI 3s 4P-3p 4D* 2-2 0.21 0.25[.03] 1.2(0.1] 4661.63 VI 3s 4P-3p 4D* 4-4 0.27 0.25[.02] 0.9[0.1] 4673.73 VI 3s 4P-3p 4D* 4-2 0.04 0.04[.0l] 1.0(0.3] 4676.24 VI 3s 4P-3p 4D* 6-6 0.22 0.22[.0l] 1.0(0.1] 4317.14 V2 3s 4P-3p 4P* 2-4 1.00 1.00 1.0 4319.63 V2 3s 4P-3p 4P* 4-6 0.98 0.71[.10] 0.7(0.1] 4325.76 V2 3s 4P-3p 4P* 2-2 0.19 0.28[.04] 1.5(0.2] 4345.56 V2 3s 4P-3p 4P* 4-2 0.90 0.92[.20] 1.0(0.3] 4349.43 V2 3s 4P-3p 4P* 6-6 2.27 1.07[.20] 0.5[0.1] 4366.89 V2 3s 4P-3p 4P* 6-4 1.02 0.85[.40] 0.8[0.3] 4414.90 V5 3s 2P-3p 2D* 4-6 1.00 1.00 1.0 4416.97 V5 3s 2P-3p 2D* 2-4 0.56 0.86[.50] 1.5(0.9] 4452.37 V5 3s 2P-3p 2D* 4-4 0.11 0.20[.ll] 1.8(1.0]

4069.62 VIO 3p 4D*-3d 4F 4-6 0.74 3.02[2.2] 4.1(3.0] 4072.16 VIO 3p 4D*-3d 4P 6-8 0.69 1.33(1.0] 1.9(1.4] 4075.86 VIO 3p 4D*-3d 4F 8-10 1.00 1.00 1.0 4085.11 VIO 3p 4D*-3d 4F 6-6 0.13 0.31[.21] 2.4(1.6] rad-#; 4083.90 V48b 3d 4F-4f G4* 6-8 0.29 0.27[.05] 0.9(0.2] 4087.15 V48c 3d 4F-4f G3* 4-6 0.27 0.35[.04] 1.3(0.2] 4089.29 V48a 3d 4F-4f G5* 10-12 1.00 1.00 1.0 4275.55 V67a 3d 4D-4f F4* 8-10 1.20 1.17[.20] 1.0(0.2] 4303.82 V53a 3d 4P -4f D3* 6-8 0.47 0.37[.06] 0.8[0.1] 4609.44 V92a 3d 2D -4f F4* 6-8 0.56 0.53[.04] 1.0(0.1] 2.6. Ionic Abundances from ORLs 89

Ao(À) Mult Termj-Termu gj-gu Ipred lofcs Io6s/Ipred

NGC 3918

4638.86 V I 3s 4P-3p 4D* 2-4 0.21 1.04[.23] 5.0[1.1] 4641.81 VI 3s 4P-3p 4D* 4-6 0.53 1.03[.40] 1.9[0.7] 4649.13 VI 3s 4P-3p 4D* 6-8 1.00 1.00 1.0 4650.84 VI 3s 4P-3p 4D* 2-2 0.21 0.20[.07] 1.0[0.4] 4661.63 VI 3s 4P-3p 4D* 4-4 0.27 0.35[.04] 1.3[0.2] 4673.73 VI 3s 4P-3p 4D* 4-2 0.04 0.15[.03] 3.8[0.8] 4676.24 VI 3s 4P-3p 4D* 6-6 0.22 0.31[.02] 1.4[0.1] 4317.14 V2 3s 4P-3p 4P* 2-4 1.00 1.00 1.0 4319.63 V2 3s 4P-3p 4P* 4-6 0.98 0.41 0.4 4325.76 V2 3s 4P-3p 4P* 2-2 0.18 0.34 1.9 4349.43 V2 3s 4P-3p 4P* 6-6 2.28 1.53 0.7 4414.90 V5 3s 2P-3p 2D* 4-6 1.00 1.00 1.0 4416.97 V5 3s 2P-3p 2D* 2-4 0.56 0.56 1.0 4452.37 V5 3s 2P-3p 2D* 4-4 0.11 1.26 11.

4069.62 VIO 3p 4D*-3d 4P 4-6 0.74 8.11[4.1] 11[5.6] 4072.16 VIO 3p 4D*-3d 4P 6-8 0.69 3.62[1.8] 5.3[2.6] 4075.86 VIO 3p 4D*-3d 4P 8-10 1.00 1.00 1.0 4085.11 VIO 3p 4D*-3d 4P 6-6 0.13 0.30[.22] 2.3[1.7] 4121.46 V19 3p 4P*-3d 4P 2-2 0.36 3.54 9.8 4132.80 V19 3p 4P*-3d 4P 2-4 0.70 0.69 1.0 4153.30 V19 3p 4P*-3d 4P 4-6 1.00 1.00 1.0 4156.53 V19 3p 4P*-3d 4P 6-4 0.16 1.42 8.9 4169.22 V19 3p 4P*-3d 4P 6-6 0.34 0.18 0.5

4083.90 V48b 3d 4P-4f 04* 6-8 0.29 0.50[.18] 1.7[0.6] 4087.15 V48c 3d 4P-4f 03* 4-6 0.27 0.43[.17] 1.6[0.6] 4089.29 V48a 3d 4P-4f 05* 10-12 1.00 1.00 1.0 4275.55 V67a 3d 4D-4f P4* 8-10 1.24 1.15 0.9 4291.25 V55 3d 4P-4f 03* 6-8 0.26 0.44 1.7 4294.78 V53b 3d 4P-4f D2* 4-6 0.29 0.45 1.6 4303.83 V53a 3d 4P-4f D3* 6-8 0.47 0.50 1.1 4466.42 V86b 3d 2P-4f D2* 4-6 0.10 0.43 4.3 4610.20 V92c 3d 2D-4f P2* 4-6 0.57 0.67 1.2 4669.27 V89b 3d 2D-4f D2* 4-6 0.04 0.22 5.5

NOG 2022 r^a-3p; 4638.86 VI 3s 4P-3p 4D* 2-4 0.21 0.29[.ll] 1.4[0.5] 4641.81 VI 3s 4P-3p 4D* 4-6 0.53 0.69[.16] 1.3[0.3] 4649.13 V I 3s 4P-3p 4D* 6-8 1.00 1.00 1.0 4650.84 VI 3s 4P-3p 4D* 2-2 0.21 0.24[.21] 1.1[1.0] 4661.63 VI 3s 4P-3p 4D* 4-4 0.27 0.30[.06] 1.1[0.2] 4676.24 VI 3s 4P-3p 4D* 6-6 0.22 0.21[.08] 1.0[0.4]

4072.16 VIO 3p 4D*-3d 4P 6-8 0.69 1.07[.17] 1.6[0.3] 4075.86 VIO 3p 4D*-3d 4P 8-10 1.00 1.00 1.0

4089.29 V48a 3d 4P-4f 05* 10-12 1.00 1.00 1.0 4276.75 V67b 3d 4D-4f P3* 6-8 1.32 1.08 0.8 o en

CO q r q r es CN TT’ f—1 l>- r-t 05 l> CO es CO eo es q ? q q q q q q d . â â d . à à â â d . à â à N en q q o T-i q O) 00 00 q q 00 CO 0 0 *-| CO Tt* q CD 1-H Ot) CO O d es d d d d CO es d

eô’ eo’ I n ’ oT oT oT cô’cô' cT 'io' d " d " OÔ* o T d ’ d*' ôô"Td CO CO 00 CO CO q q q q q q q q q q q q g g o ’-j q g g w § § s o - ' î o oT oô" oô" cô" es" IC^ “d" "in' oT o CO es ■o< CN CN 00 t-- Tf q 00 i> ïO CN 7 S § d d d d d d d d d d d CO d d d CN d eo

CO o t^ CO o t - CN o 00 o CD o CD 05 CO eo es lO q es es es es lO q CN CN CN q 05 05 (N q q iD q ID 7 CD § 7 s § d d d d d d d d d d d d CN d d d C) d d d d i eo

"9 3 O O CD 00 00 T 9 9 7 7 7 7 9 °T7 T 7 T 7 7 Y 7 7 -s 7 7 7 7 7 7 7 7 7 I I I es eo CS eo es CD CN CD CN CD CD CN CN 'et* CD CD Tl* eo § g oô 00 eo i nil Tf CO CO O y i g CO Tf q q q q q q q Û Q Q Û Q Ph PL, PU a a Û û QÛQ CO lO CN CO Tf a o k 'Cf Tf -ef CN CN CN t3 73 T3 13 ü ü ü q q Ü a q q q q q q a a a a A a PL a a a a a a a a 7 7 7 7 7 7 7 7 ^ ^ ^ ^ ^ ^ ^ TT 7 7 7 7 7 7 ? ? ? ? 7 ?? 7 ? ? 7 7 7 7 7 * * * * * eu q q q q q q q q PL, Oh Pu PL, PU PL, PL, a PU a a a a a a Pu tu q q q § (N es Tf Tf Tf TT Tf CN CN CN CN CN Q Tf Tf Tf Tf a CO CO CO CO CO CO CO CO CO CO CO CO CO CO CO CO CO CO CO CO CO c^c^ eo eo e^ eo I -O x> Ci a ci o <à d 00 es o o o o 00 00 00 t> i> CO es 00 05 -—irH,—irHi-H,—I cseseseses koin loicm Tf Tf eo eo lO 05 > > > > > > > > >>>>>>>>>>>>> >>> > > > > > > > > >>>>>>> > g O -N CO CD CO CO Tf CO CD CO a o N o I> CN CD CD CN CD CD O LO 05 LO CD CN o 05 0 0 CN r-H °< OC q 00 0 0 00 CD q CD LO q 05 05 q CO Go CD q 00 q CD q 00 s 05 CN LO q 00 05 05 00 3 d d oô 05 d CD 05 LO 05 CD CD CD CN 05 CN lÔ lÔ 05 CN LÔ LO 1 CÔ 05 LÔ CN cô (35 cô O «0 CO lO CD N CO Tf LO CD N CD LO a CD l> 00 CD N 00 00 00 00 N 00 o O 00 'o CD GO CD CD CD CD CD CD CD CD CD CD CD CD CO CO CO CO CO Tf GO OOOOO O o o GO o o O CN CN CO CD O Tf Tf Tf Tf Tf Tf 2.6. Ionic Abundances from ORLs 91

Ao(A) Mult Termj-Termu g l-g u ^pred lobs ^obs/^pred

4087.15 V48c 3d 4F -4f G3* 4-6 0.30 0.33[.07] 1.1[0.2] 4089.29 V48a 3d 4F-4fG5* 10-12 1.00 1.00 1.0 4275.55 V67a 3d 4D -4f F4* 8-10 1.24 1.18[.18] 1.0[0.2] 4303.82 V53a 3d 4P-4fD3* 6-8 0.48 0.33[.07] 0.7[0.2] 4466.42 V86b 3d 2P-4fD2* 4-6 0.10 0.19[.03] 1.9[0.3] 4609.44 V92a 3d 2D -4f F4* 6-8 0.56 0.29[.08] 0.5[0.1]

IC 4406 ras-apl 4638.86 VI 3s4P-3p 4D* 2-4 0.21 0.70[.ll] 3.3[0.5] 4641.81 V I 3s 4P-3p 4D* 4-6 0.53 1.02[.12] 1.9[0.2] 4649.13 VI 3s4P-3p4D* 6-8 1.00 1.00 1.0 4650.84 VI 3s 4P-3p 4D* 2-2 0.21 0.37[.09] 1.8[0.4] 4661.63 VI 3s 4P-3p 4D* 4-4 0.27 0.38[.09] 1.4[0.3] 4676.24 VI 3s 4P-3p 4D* 6-6 0.22 0.19[.10] 0.9[0.5] ('gp-ad; 4069.62 VIO 3p 4D *-3d 4P 2-4 0.74 1.62[.60] 2.2[0.8] 4072.16 VIO 3p4D*-3d 4F 6-8 0.69 1.49[.51] 2.2[0.8] 4075.86 VIO 3p 4D*-3d 4F 8-10 1.00 1.00 1.0 NGC 6818

4638.86 VI 3s 4P-3p 4D* 2-4 0.21 1.28[.30] 6.1[1.4] 4641.81 V I 3s 4P-3p 4D* 4-6 0.53 2.94[.41] 5.5[0.8] 4649.13 VI 3s 4P-3p 4D* 6-8 1.00 1.00 1.0 4650.84 VI 3s 4P-3p 4D* 2-2 0.21 0.16[.ll] 0.8[0.6] 4661.63 VI 3s 4P-3p 4D* 4-4 0.27 0.50[.07] 1.9[0.3] ("gp-ad; 4069.78 VIO 3p 4D*-3d 4F 2-4 0.74 0.30[.26] 0.4[0.4] 4072.16 VIO 3p 4D*-3d 4F 6-8 0.69 0.40[.14] 0.6[0.2] 4075.86 VIO 3p 4D*-3d 4F 8-10 1.00 1.00 1.0 4085.11 VIO 3p 4D*-3d 4F 6-6 0.13 0.16[.06] 1.2[0.5] NGC 2440

4638.86 VI 3s 4P-3p 4D* 2-4 0.21 2.16[.56] 10.[2.6] 4641.81 VI 3s 4P-3p 4D* 4-6 0.53 1.62[.37] 3.1[0.7] 4649.13 V I 3s 4P-3p 4D* 6-8 1.00 1.00 1.0 4650.84 VI 3s 4P-3p 4D* 2-2 0.21 0.40[.26] 1.9[1.2] 4661.63 V I 3s 4P-3p 4D* 4-4 0.27 0.27[.06] 1.0[0.2] 4676.24 VI 3s 4P-3p 4D* 6-6 0.22 0.56[.08] 2.5[0.4] 4414.90 V5 3s 2P-3p 2D* 4-6 1.00 1.00 1.0 4416.97 V5 3s 2P-3p 2D* 2-4 0.56 0.80[.22] 1.4[0.4] 4452.37 V5 3s 2P-3p 2D* 4-4 0.11 1.70[.37] 15.[3.3]

NGC 3132

4638.86 V I 3s 4P-3p 4D* 2-4 0.21 0.63[.47] 3.0[2.2] 4641.81 VI 3s 4P-3p 4D* 4-6 0.53 1.03[.58] 1.9[1.1] 4649.13 V I 3s 4P-3p 4D* 6-8 1.00 1.00 1.0 4650.84 VI 3s 4P-3p 4D* 2-2 0.21 0.74[.38] 3.5[1.8] 4661.63 V I 3s 4P-3p 4D* 4-4 0.27 0.70[.34] 2.6[1.3] 4676.24 V I 3s 4P-3p 4D* 6-6 0.22 0.46[.23] 2.1[1.1] ('gp-gd; 4069.62 VIO 3p 4D*-3d 4F 2-4 0.74 1.14[.78] 1.5[1.0] 2.6. Ionic Abundances from ORLs 92

Ao(Â) Mult Term(-Termu gl-gu Ipred lobs lobs/Ipred

4072.16 VIO 3p 4D*-3d 4F 6-8 0.69 0.41[.34] 0.6[0.5] 4075.86 VIO 3p 4D*-3d 4F 8-10 1.00 1.00 1.0 4085.11 VIO 3p 4D*-3d 4F 6-6 0.13 0.12[.03] 0.9[0.2] LMC N141

4638.86 VI 3s 4P-3p 4D* 2-4 0.21 0.20[.10] 1.0[0.5] 4641.81 VI 3s 4P-3p 4D* 4-6 0.53 0.95[.28] 1.8[0.5] 4649.13 VI 3s 4P-3p 4D* 6-8 1.00 1.00 1.0 4650.84 VI 3s 4P-3p 4D* 2-2 0.21 0.66[.41] 3.1[1.9] 4661.63 VI 3s 4P-3p 4D* 4-4 0.27 0.41[.14] 1.6[0.6] 4676.24 VI 3s 4P-3p 4D* 6-6 0.22 0.37[.09] 1.7[0.4]

4069.62 VIO 3p 4D*-3d 4F 2-4 0.74 0.74[.36] 1.0[0.5] 4072.16 VIO 3p 4D *-3d 4F 6-8 0.69 0.49[.18] 0.7[0.3] 4075.86 VIO 3p 4D*-3d 4F 8-10 1.00 1.00 1.0 SMC N87

4638.86 VI 3s 4P-3p 4D* 2-4 0.21 0.95[.17] 4.5[0.8] 4641.81 VI 3s 4P-3p 4D* 4-6 0.53 0.80[.26] 1.5[0.4] 4649.13 VI 3s 4P-3p 4D* 6-8 1.00 1.00 1.0 4650.84 VI 3s 4P-3p 4D* 2-2 0.21 0.73[.15] 3.5[0.7] 4661.63 VI 3s 4P-3p 4D* 4-4 0.27 0.34[.04] 1.3[0.2] 4676.24 VI 3s 4P-3p 4D* 6-6 0.22 0.34[.ll] 1.6[0.5]

4069.62 VIO 3p 4D*-3d 4F 4—6 0.74 2.50[2.4] 3.4[3.3] 4072.16 VIO 3p 4D*-3d 4F 6-8 0.69 1.20[1.1] 1.7[1.6] 4075.86 VIO 3p 4D*-3d 4P 8-10 1.00 1.00 1.0 4078.84 VIO 3p 4D*-3d 4F 4-4 0.11 0.50[.62] 4.5[5.6] 4085.11 VIO 3p 4D*-3d 4F 6-6 0.13 0.24[.39] 1.8[2.9] 2.6. Ionic Abundances from ORLs 93

2.6.5 Total C, N and O Abundances from ORLs

In this Section we discuss the derivation of the total abundances of C, N and O for our PN sample, as derived from ORLs. The results are summarized in Ta­ bles 2.13, 2.14, and 2.15 respectively. We will discuss each PN individually in order to highlight differences in the adopted ICF scheme.

N G C 3242: the C II A4267, C III AA4069, 4187 and A4650 lines are detected; we adopt = 2.03 x 10“^, as an intensity weighted mean from the AA4187, 4650 lines only. The unseen C"'" is corrected for using an icf{C) = 1 + = 1.01. We calculate the C/H fraction using,

C /H = 2 c/(C) X (C^+ + C^+)/H+. (2.3)

This yields C/H =8.45x10“'^; this value is only 2per cent more than the one found if = 1.88 x 10“ ^, as derived from C IV A4658 is not added in. The latter value is an upper limit to the fraction, since A4658 is blended with

[Fe III] A4658. From an intensity weighted mean of N^'^/H'^ ratios from seven N ii lines (cf.

Table 2.10), we find = 1.71 x 10“'^. From the N III A4379 line,

= 6.85 X 10“ ^. We use the CEL ratio N’^/N^'*' = 0.0145 (cf. Table 2.7) to account for the unseen N'*'; the errors introduced should be negligible, since N'^’/N = 0.008. Hence, summing three ionic stages, we find N/H = 2.42 x 10“^.

From our extensive ORL O II line survey, we derive G^'^'/H'*' = 6.28 x 10“ “^, from six O II multiplets. A recombination line 0"'"/H"'" abundance is not available, so the unseen O"'" is corrected for assuming that the CEL ratio = 0.0091 holds for the ORLs too; in general this is not true, as we saw from the analysis of C and N. However, in view of the minor concentration of O’*' in this object (about 1 per cent), the errors introduced should be negligible. We then employ the stan­ dard correction factor to account for the unseen O^"*", icf{0) = (1 + He^"*"/He"'")^/^ = 1.17, and calculate the total 0/H abundance, using. 2.6. Ionic Abundances from ORLs 94

0 /H = %c/(0) X (0+ + 0^+)/H+. (2.4)

This yields, 0 /H = 7.41 x Alternatively, we can adopt 0^^/H"^ =

2.05 X 10“^, as derived by LD93a from the 0 III V 8 multiplet at 3265 A; hence, summing all three ionic stages we find 0/H = 8.39 x 10“'^, i.e. 13 per cent larger than yielded by the ICF method. We adopt this latter value.

NGC 5882: the C II A4267, C ill AA4069, 4187 and A4650 lines are detected; we adopt C^'^’/H ’*' = 5.27 x 10“^, as an intensity weighted mean from the AA4187, 4650 lines only. No is expected to exist, since only 2 per cent of He is in the He^^ form. Using Eq. 2.3 and icf{C) = 1.03, we find, C/H = 4.43 x 10“ '^. From an intensity weighted mean of five N^'*'/H+ ratios, N^+/H+ = 2.24 x 10"^ is found. From the N III A4379 line, N^'^'/H"^ = 5.75 x 10~^. As in NGC 3242, we use the CEL ratio, N'^'/N^'*' = 0.0278 in this case (cf. Table 2.8), to account for the unseen N+; we obtain N /H = 2.88 x 10“'^.

From eight O II ORL multiplets we derive O^'^/H'^ = 9.70 x lO”'^. We account for unseen O'*" as in the case of NGC 3242, using the CEL ratio 0"'"/0^''" = 0.029 (in this object the concentration of O^ is about 3per cent); we have icf{0) = 1.01, so Eq. 2.4 yields, 0 /H = 10.08 x 10~^^.

N G C 5315: the C II A4267, C III A4187 lines are detected; no C^"*" exists since the He^'*"/H"'" fraction is negligible. We correct for unseen C'*’ using the CEL ratio ofC'^’/C^’*' = 0.08 (given the small concentration of C"*", about 7per cent, no significant uncertainty is introduced) ; hence we sum three ionic stages and find

C /H = 7.29 X 10-4. From an intensity weighted mean of four N^+/H+ ratios, N^'*'/H+ = 3.43 x 10~^ is found; the N III A4379 line is not detected. If we assume that the N'*'/N^''" = 0.163 and N^"'"/N^"'" = 0.565 ratios, as given by the CELs are also valid for the ORLs, we can account for both N+ and N^+, hence, N /H = 5.93 x 10“ ^. This 2.6. Ionic Abundances from ORLs 95

Table 2.13: Carbon abundances from optical recombination lines.

NGC 3242 NGC 5882 NGC 5315 NGC 3918 NGC 2022 NGC 6818 /(A4267) 0.620 0.399 0.706 0.497 0.820 0.449 10^ X C^+/H+ 6.15 3.77 6.60 5.02 8.83 4.69 /(A4069) 0.755 0.412 * * 1.75 0.957 10^ X C^+/H+ 3.64 2.11 ** 7.49 3.97 /(A4187) 0.204 0.0376 0.0107 0.152 0.332 0.118 10^ X C^+/H+ 2.83 0.552 0.161 2.03 4.07 1.41 /(A4650) 0.590 0.173 * 0.418 0.974 0.389 10^ X C^+/H+ 1.75 0.522 * 1.22 2.77 1.10 Adopted 10^ X C^+/H+ 2.03 0.527 0.161 1.44 3.10 1.37 /(A4658) 0.0776 * * * 0.109 1.03 0.180 ** 10'^ X C^+/H+ 0.188 0.262 2.62 0.447 ic/(C ) 1.01 1.03 * * 1.02 * 10^ X C /H 8.45 4.43 7.29 7.50 14.84 7.71

NGC 3132 NGC 2440 NGC 6302 IC4406 IC4191 My Cn 18

/(A4267) 0.697 0.403 0.163 0.805 0.546 0.252 10^ X C^+/H+ 6.60 4.47 2.00 7.72 5.16 3.80 /(A 4187) * 0.222 * 0.0912 0.0711 * * * * 10'^ X C^+/H+ 2.58 1.32 1.04 /(A4650) * 0.337 * 0.238 0.359 * * 10^ X C^+/H+ * 1.06 * 0.715 1.08 Adopted * * 10^ X C^+/H+ * 1.66 0.883 1.07 /(A4658) * 0.722 *** * 10^ X C^+/H+ * 1.82 * * * * 2C/(C) 1.93 1.28 2.72 1.31 1.04 1.59 10^ X C/H 12.74 10.21 5.45 11.27 6.48 6.04 2.6. Ionic Abundances from ORLs 96 result may be quite uncertain, since probably the CEL 7^ ORL (e.g. as for N G C 3918 below).

From five O II ORL multiplets, 0^"^/H"*" = 8.57 x although this is a rel­ atively low excitation object, a substantial O^"^ concentration has however been derived from the A1401 CEL line. If we account for unseen O'*" using the CEL ratio O’^'/O^'*’ = 0.098, we find 0/H = 9.41 x 10“^^, ignoring 0^+/H+; however, if we further assume that O^'^’/O^'^ = 0.364, as given by CELs, holds for ORLs too, then this brings the total ORL O abundance up to 0 /H = 12.53 x lO'"^, i.e. 33 per cent higher. We adopt this latter value.

NGC 3918: the C II A4267, C III AA4187, 4650 lines are detected; C^+/H''" =

1.44 X 10“'^ is derived, as an intensity weighted mean from the AA4187, 4650 lines.

Using the [Fe III] A4702 line, we estimate that 67 per cent of the detected A4658 in­ tensity is due to C'^'*', i.e. /(C IV A4658) = 0.1093; hence, C^'^/H"^ = 2.62xl0~^. An ORL abundance is not available for C"^, so it is estimated using the ratio C+/C2+ — 0.155 (cf. Table 2.7) derived from CELs (given the small ionic con­ centration of C'*', 9 per cent maximum, the error introduced should be negligible) ; adding this to the remaining three ionic stages we find, C/H = 7.50 x 10"^ (if instead we correct for C"*" using an icf{C) = 1.09, we arrive at a value which is smaller by 3 per cent only). Prom an intensity weighted mean of four N^+/H+ ratios, N^"""/H+ = 1.42 x 10"^ is found; from the N III A4379 line, N^+/H"^ = 6.41 x 10“ ^; all of the measured in­ tensity of the A4606 line is attributed to N iv (no contribution from N ll A4607.2 is expected, since the stronger predicted N II A4601 line is absent); therefore,

N^'^/H'^ = 1.90 X 10"^. We account for the missing N"'" using the CEL ratio N^/N^""" = 0.203; summing all four ionic stages, we have, N /H = 2.56 x 10~^. For this PN, we find that just as is the case for relative ionic fractions of carbon, the CEL N^+/N2+ ratio (= 1.22) 7^ the ORL N^+/N^+ ratio (= 0.45).

We have derived 0^"*"/H+ and 0^"^/H+ abundances from five O II multiplets 2.6. Ionic Abundances from ORLs 97

and the O IV A4632 line, respectively; = 5.36 x 10“'^ and =

4.74 X 10“^. We account for unseen O’*" and using the CEL ratios, = 0.086 and = 1.35; hence, summing all four ionic stages we have 0 /H

= 13.53 X 10-^.

N G C 2022: the C II A4267, C ill AA4069, 4187 and A4650 lines as well as the 0 iv A4658 line are detected; = 3.10 x 10“^, is the intensity weighted mean

from the AA4187, 4650 lines only. We attribute the total A4658 intensity to C IV,

since other [Fe III] lines are not present; furthermore this is a very high excitation object - He^'^/He = 0.88 - and we expect a substantial concentration of we find — 2.62 x 10“^. Using icf{C) = 1.02, to correct for unseen and summing the three higher ionic stages, we obtain C/H = 14.84 x 10“^. We derive N^"*"/H"^ = 4.38 x lO”'^ from an intensity-weighted mean of three

N^'^’/H'^ ratios; from N III A4379, N^'^/H'*' = 2.08 x 10“ '^ and from N IV A4606,

N'^'^’/H'^ = 3.68 X 10“ ^ (as for NGC 3918, no contribution by N II A4607.2 is found for the A4606 line). We correct for unseen N""" using the CEL ratio N"*"/N^"*" = 0.013 and sum a total of four ionic stages to obtain, N/H = 6.88 x 10“^.

From four O II multiplets and the O IV A4632 line, respectively, we derive

Q2+/H+ = 13.01 X 10-4 and Q4+/H+ = 1.10 x 10~4; we also adopt 0^+/H + =

3.62 X 10-4, as derived by LD93a from the O III V 8 multiplet at 3265 A. The un­ seen Q4" is accounted for using the CEL ratio O^'/O^^- — 0.020; hence we obtain

0 /H = 17.99 X 10-4.

N G C 6818: the C II A4267, C ill AA4069, 4187 and A4650 and C iv A4658 lines are detected; C^4-yH+ — 1.37 x 10-4, is the intensity-weighted mean from the AA4187, 4650 lines only. Using exactly the same argument as in the case of NGC 2022, we attribute the A4658 line to C4+ only; we obtain C4+/H+ = 4.47 x 10-^. If we assume that the ratio C^’/C^’*- = 0.129 derived from the CELs is also valid for the ORLs, we deduce a total abundance of C/H = 7.71 x 10-4. 2.6. Ionic Abundances from ORLs 98

Table 2.14: Nitrogen abundances from recombination lines in Galactic PN.

NGC 3242 NGC 5882 NGC 5315 NGC 3918 NGC 2022 NGC 6818

10^ X N ^+/H + 1.71 2.24 3.43 1.42 4.38 4.27 /(A4379) 0.158 0.136 * 0.146 0.472 0.203 10^ xN^+/H+ 6.85 5.75 * 6.41 20.80 &87 /(A 4606) * * * .0660 0.112 .0460 * 10^ X N ^+/H + ** 2.07 3.68 1.46 CEL N +/N ^+ 0.008 0.028 0.163 0.203 0.013 0.241

10^ X N/H 2.42 2.88 5.93 2.56 6.88 6.33

NGC 3132 NGC 2440 NGC 6302 IC 4406 IC4191 My Cn 18

10"^ xN^+/H+ 3.40 4.77 5.64 1.90 5.00 20.35 /(A 4379) .0571 0.324 0.182 0.172 0.270 * * 10^ X N^+/H+ 2.47 14.30 8.07 7.31 11.50 7(A4606) * 0.475 0.668 ** * * * * 10® X N ^+/H + * 17.0 2Z8 CEL N+/N^+ 1.296 0.475 0.435 0.536 0.393 *

10^ X N/H 8.05 10.17 11.18 3.65 8.12 > 2 0 .3 5

The N II AA5676, 5679 lines from multiplet V3 are detected; they yield

= 4.27 X 10“ ^; from N III A4379 we find = 8.87 x 10“^, while from N IV A4606 we deduce — 1.46 x 10"^; the unseen N+ is corrected for using the CEL ratio = 0.241. Summing a total of four ionic stages N/H =

6.33 X 10“ ^ is obtained. We have derived = 7.33 x 10^^ and = 6.41 x 10“ ^ from four

O II multiplets and the O IV A4632 line, respectively; we further adopt

= 2.37 X 10“^, as derived by LD93a from the O III V 8 multiplet at 3265 Â. The unseen O"*" is accounted for using the CEL ratio = 0.090; hence, 0/H —

11.00 X 10-^.

N G C 3 1 3 2 : only C II A4267 is detected. If we then use the standard ICF to account for the unseen C^"*", icf {C) = O/ = 1.93, and calculate the carbon abundance as C/H = icf{C) x then we find C/H = 12.74 x 10"^. If instead, we assume that the ratio C^'^/C^'*' = 0.111 given by the CELs is also 2.6. Ionic Abundances from ORLs 99 valid for the ORLs and use it to estimate the ORL abundance, we find,

C /H = 13.86 X 10~^, i.e. 9 per cent larger. However, from observations of 7 PN it is found that, in general, ORL ^ CEL hence, we adopt the former C/H value.

From three N II ORL lines and the N III A4379 line we obtain, N^'^'/H'*' =

3.40 X 10“^ and N^"*"/H"'' = 2.47 x 10"^, respectively. We correct for the unseen N+ using the CEL ratio N'*'/N^‘^ = 1.296 and sum a total of three ionic stages to find N /H = 8.05 x 1 0 "!

From the O II V 1 and V 10 multiplets we have derived 0^"'"/H'^ = 8.15 x 10“ ^; we account for unseen O^'*’ using an icf{0) = 1.02 from Eq. 2.4; we account for the unseen 0"*" using the CEL ratio O'^/O^"*' = 0.894, hence, 0 /H = 15.74 x 10“ '^. The concentration of O’*' in this PN is about 46 per cent and it may be that some uncertainty has been introduced by assuming that the CEL 0+/0^+ ratio equals the ORL O'^'/O^'"' ratio.

N G C 2440: the C II A4267, C III AA4187, 4650 and C IV A4658 lines are detected; we attribute the A4658 intensity solely to thus C^"*"/H"*" = 1.66 x 10“'^ and

C^"'"/H'^ = 1.82 X 10~"^. Using an icf {C) = l + O'^’/O^'^ = 1.28 to account for unseen C'*' and summing the three higher ionic stages, we find C/H = 10.21 x 10“ ^.

Four N II lines are detected; their intensity weighted mean yields N^^/H”^ =

4.77 X 10“ '^; from N III A4379 we find N^'^/H'*' = 1.43 x 10“^; both N IV AA4606,

4707 lines are detected, yielding a mean N'^'^/H'*' = 1.70 x 10“'^. No N+ abundance is available either from ORLs or CELs (our wavelength coverage did not include the [N ii] lines). Thus, we adopt the CEL ratio N'^/N^'^ = 0.475, from KB94, and use it to correct for the unseen N+; summing four ionic stages we obtain, N/H =

10.17 X 10-4.

From four O II multiplets and the 0 IV A4632 line, respectively, we derive

0^"*"/H"'" = 5.23 X 10-4 and O^'^/H"'" = 1.01 x 10“4 from our optical spectra; we further adopt = 2.56 x 10~4, as derived by LD93a from the O III V 8 2.6. Ionic Abundances from ORLs 100

multiplet at 3265 Â, together with the CEL ratio 0 + /0 ^ + = 0.285 by the same authors (as reported for their PA = 270^ slit); we assume that this ratio is valid for ORLs too and hence derive, 0 /H = 10.29 x 10“^.

NGC 6302: only C II A4267 is reliably detected; using the same ICF scheme as for NGC 313, yielding icf{C) = 2.72, from Eq. 2.3 we obtain C /H = 5.45 x 10“ ^. Three N II lines are detected; their intensity weighted mean yields N^"*"/H"*" =

5.64 X 10-4; ^ jji A43?9 yields N^+/H+ = 8.07 x lO'^; N iv A4606 yields N4+/H+

= 2.28 X 10~4 (no contribution is expected from N ll A4607). N+ is accounted for using the CEL ratio = 0.435; summing four ionic stages we find N /H = 11.18x10-4. We obtain 0^"'"/H''' = 3.28 x lQ-4 from the O il V I multiplet and the 3d-4f

A4089 line; 04+/H"'' = 7.84 x 10“ ^ is derived from O IV A4632; we assume that the CEL O’^'/O^'^ = 0.092 and /O^'^ = 0.895 ratios are valid for ORLs as well, and deduce 0 /H = 7.30 x 10-4.

IC 4406: C II A4267, C III AA4187, 4650 are detected; the He^"^ concentration is only 6 per cent, so it is assumed that the C4+/H"^ abundance is negligible; C^'*‘/H'^

= 8.83 X 10-5 is derived. Using Eq. 2.3, with icf{C) = 1.31 to account for unseen C+, we find C/H = 11.27 x 10"4.

From the sole detected N II A4630 line, N^"""/H+ = 1.90 x 10“4; from N ill A4379, N^^/H’*' = 7.31 x lOr^. N"*" is accounted for using the CEL ratio N’^’/N^'^ = 0.536; summing three ionic stages we find N/H = 3.65 x 10-4, We obtain 0^"*"/H"'" = 7.06 x 10-4 and use an icf{0) = 1.04 to correct for unseen the missing 0 ’*' is corrected for using the CEL ratio 0*^/0^+ = 0.471 (some error may be introduced since in this object, the O'*' concentration is about 31 per cent of all oxygen). Hence, from Eq. 2.4 we obtain 0 /H = 10.80 x 10-4.

IC 4191: C II A4267, C III AA4187, 4650 are detected; C^+/H+ = 1.07 x 10-4 is 2.6. Ionic Abundances from ORLs 101 found. No significant is present. Using Eq. 2.3 with icf{C) = 1.04 to account for unseen C'*', we find C/H = 6.48 x lO”'^.

Six N II lines have been detected; an intensity weighted mean yields =

5.00 X lO”'^; from N III A4379, we derive = 1.15 x 10"^. N+ is accounted for using the CEL ratio = 0.393, where = 2.45 x 10~^ from the

[N III] 57-/im line, as given by Liu et al. (2001) and N +/H ’*' from our optical observations; hence, summing three ionic stages we obtain N/H = 8.12 x 10"^.

We derive G^'^’/H'*' = 12.92 x 10“ “^ from four O II multiplets and use an icf(0) = 1.07 to correct for unseen the missing O"'" is corrected for using the CEL ratio 0+/0^+ = 0.041 (the error introduced in this case should be negligible, since the 0 ’*' concentration in this PN is only about 4per cent). Hence, we deduce 0/H

= 14.39 X 10-^.

My Cn 18: only C II A4267 is seen; no significant amounts of C^'*' or are expected; using Eq. 2.3, with icf{C) = 1.59 to account for the unseen C"*" only, we find C/H = 6.04 x 10"^. Three N li lines of multiplet V5 have been detected; from an intensity weighted mean N^'^’/H'^ = 20.35 x lO"'^. Apart from optical CELs no other information on N exists for this PN and the standard ICF scheme cannot account for unseen ionization stages. Therefore, as a lower limit we adopt N/H >20.35 x 10“^.

We derive O^^/H"*" = 6.43 x 10"'^ from the O II V 1 AA4649, 4650 lines for this low excitation nebula; the unseen O’*” is corrected using the CEL ratio 0^/0^^ = 0.59 (O’’' represents a significant amount of the total oxygen, about 37per cent). The resulting 0/H = 10.24 x lO”'^ may be somewhat uncertain.

S M C N 8 7 : C II A4267, C ill AA4187, 4650 are detected; C^+/H+ = 6.75 x lO'^ is found. No C^’*' is present; using Eq. 2.3, with icf{C) = 1.02 to account for the unseen C"*" only, we find C/H = 7.55 x 10“^. If instead, we assume that the ratio C’^'/C^^ = 0.095 derived from the CELs is also valid for the ORLs and can be 2.6. Ionic Abundances from ORLs 102

Table 2.15: Oxygen abundances from recombination lines in Galactic PN

NGC 3242 NGC 5882 NGC 5315 NGC 3918 NGC 2022 NGC 6818

10^ x Q2+/H+ 6.28 9.70 8.57 5.36 13.01 7.33 * 10^ X Q 2+/H + 2.05 * * 3.62 2.37 /(A4632) * * * 0.150 0.334 0.200 10^ X 0 ^ + /H + * * * 4.74 11.00 6.41 CEL G +/Q 2+ 0.0091 0.029 0.098 0.086 0.020 0.090 icf{0) * 1.01 ** * * 10^ X O/H 8.39 10.08 12.53 13.53 17.99 11.00

NGC 3132 NGC 2440 NGC 6302 IC 4406 IC4191 M y Cn 18

10'^ X Q2+/H+ 8.15 5.23 3.28 7.06 12.92 6.43

10^ X 0^ + /H + * 2.56 * * ** 7(A4632) * 0.300 0.228 * ** * 10^ X 0^ + /H + * 10.1 7.84 * * CEL Q +/Q 2+ 0.894 0.285 0.092 0.471 0.041 0.590 ic /( 0 ) 1.02 * * 1.04 * 1.003 10^ X O/H 15.74 10.29 7.30 10.80 14.39 10.24

used to estimate the abundance, we arrive at the same total C/H value. We derive 0^"'"/H^ = 1.72 x 10“'^ from the O ll V 1 multiplet; the unseen O'*" is corrected for using the CEL ratio 0^/0^^ = 0.018 (this is a medium excita­ tion PN and the concentration of O’*" is only about 2 per cent); we obtain 0/H =

1.75 X 10-4.

LM C N66: this is a highly ionized nebula reminiscent of the galactic Type-I PN NGC 6302; no ORLs of carbon are detected in agreement with the carbon-poor, nitrogen-rich nature of this nebula, hence no ORL C/H is available for this object.

From an intensity weighted mean of O II AA4641, 4649, 4650 and A4661, we derive = 8.53 x 10“4; we adopt icf{0) = 3.77 [from eq. (A7) by KB94 coupled with our ionic nitrogen abundances for this PN]; we account for the un­ seen O'*" using the CEL ratio O^'/O^'^ — 0.076 and obtain 0/H = 3.46 x 10“^.

L M C N 141; C II A4267, C l i i AA4187, 4650 are detected; C^+/H+ = 5.82 x 10"^ 2.6. Ionic Abundances from ORLs 103

Table 2.16: Recombination line Carbon and Oxygen in Magellanic Cloud PN.

SMC N87 LMC N66 LMCNI4I

/(A4267) 0.672 * 0.683 10^ X C^+/H+ 6.73 * 6.79 /(A 4I87) 0.108 * 0.0522 10^ X C^+/H+ 1.46 * 0.733 /(A4650) 0.I4I * 0.I8I 10^ X C^+/H+ 0.413 * 0.538 Adopted * 10^ X C^+/H+ 6.75 * 5.82 /(A 4658) ** * 10"^ X C^+/H+ ** * ic /(C ) 1.02 * 1.02 * 10^ X C/H 7.55 7 .5 4

10^ X 0^ + /H + 1.72 8.53 4.96 CEL 0+/0^+ 0.018 0.076 0.023 ic /( 0 ) I.OO 3.77 I.OO 10^ X O/H 1.75 3 4 .6 5 .0 8

is found. No is present; using Eq. 2.3, with icf{C) = 1.02 to account for the unseen O'*", we find C/H = 7.54 x lO”'^. If instead, we assume that the ratio C+/C2+ = 0.203 derived from the CELs is also valid for the ORLs and can be used to estimate the C'*'/H"^ abundance, we obtain C/H = 8.75 x 10“'^. In view of the error this may introduce, we adopt the former value.

We derive 0^"'"/H+ = 4.96 x 10“ ^ from O II V 1 and V 10 multiplets; the unseen is corrected for using the CEL ratio 0'*"/0^"'" = 0.023; we obtain 0 /H =

5.08 X 10-^. Nitrogen ORLs have not been detected from either SMCN87, LMCN66 or LMCN141, hence no inference can be made about their N abundances from optical recombination lines.

2.6.6 Total elemental abundances

The total elemental abundances obtained for the whole PN sample from CELs and ORLs are presented in Table 2.17. For comparison, we also list in the same 2.6. Ionic Abundances from ORLs 104

Table 2.17: Elemental abundances (by number) in Planetary Nebulae, derived from CELs and ORLs, in units where log 77(H) = 12.0.“

PN He C N O Ne S Cl Ar ORL ORLCEL ORL CEL ORLCEL CEL CEL CEL CEL

NGC 3242 11.00 8.93 8.14 8.38 7.53 8.92 8.52 7.89 6.38 4.94 5.99 NGC 5882 11.04 8.65 8.18 8.46 8.18 9.00 8.67 8.13 6.92 5.26 6.31 NGC 5315 11.08 8.86 8.33 8.77 8.52 9.10 8.79 8.30 7.31 5.41 6.56 NGC 3918 11.00 8.88 8.88 8.40 8.40 9.13 8.86 7.97 6.70 5.11 6.24 NGC 2022 11.02 9.17 8.33 8.84 7.79 9.26 8.66 7.84 6.57 4.91 6.12 NGC 6818 11.00 8.89 8.41 8.80 7.93 9.04 8.73 8.07 6.56 5.09 5.94 NGC 3132 11.08 9.11 8.50 8.91 8.37 9.20 8.82 8.49 7.03 5.36 6.73 NGC 2440 11.08 9.01 8.37 9.01 8.26 9.01 8.39 8.04 ** * NGC 6302 11.13 8.74 7.89 9.05 8.52 8.86 8.40 7.88 6.75 4.99 6.34 IC4406 11.09 9.05 8.56 8.56 8.32 9.03 8.76 8.33 6.33 5.16 6.40 IC4191 11.08 8.81 * 8.91 7.59 9.16 8.78 8.21 7.10 5.37 6.49 M y Cn 18 10.99 8.78 * 9.31 8.34 9.01 8.75 8.09 7.24 5.39 * SM C N 87 10.99 8.88 8.58 * 7.55 8.24 8.03 7.03 ** 5.32 LMC N66 11.02 * 7.52 * 7.99 9.54 8.50 7.62 6.63 * 6.10 LMC N141 11.04 8.88 8.30 * 7.95 8.71 8.29 7.38 6.61 * 5.88 Type I 11.11 * 8.48 * 8.72 * 8.65 8.09 6.91 * 6.42 non-Type I 11.05 * 8.81 * 8.14 * 8.69 8.10 6.91 * 6.38 Solar*’ 10.99 * 8.55 * 7.97 * 8.69 8.09 7.21 5.50 6.56

“ All values are for the whole nebulae, except NGC 6302, 6818, My Cn 18 and IC 4406 where values are from fixed-slit observations; Solar photospheric abundances from Anders & Grevesse (1989) and Grevesse & Noels(1993), except from oxygen which is from Allende Prieto et al. (2001).

Table the average abundances of Galactic PN (for both Type I and non-Type I nebulae) derived by Kingsburgh & Barlow and the solar photospheric abundances compiled by Anders & Grevesse (1989) and Grevesse & Noels (1993).

2.6.7 Comparison of ORL and CEL abundances

2.6.7.1 Abundance grids and ionic discrepancy ratios

The ionic abundances obtained from the UV, optical and IR collisionally excited lines (Table 2.7, 2.8) and those from ORLs (Tables 2.10, 2.11, 2.13, 2.14, 2.15 and 2.16) are presented graphically in Fig 2.2. Note that the plotted fractions derived from optical CELs are those from the [O III] AA4959, 5007 lines only. Fig 2.2 shows that in all nebulae and for all ions where ionic abundances from 2.6. Ionic Abundances from ORLs 105 both ORLs and CELs (UV, optical or IR) have been derived, the values from the optical recombination lines are consistently higher than those derived from the collisionally excited lines. This includes the Magellanic Cloud nebulae SMC N87, LMC N66 and LMC N141; no such comparisons have been reported so far regarding extragalactic PN. In Table 2.18 we present a listing of the ionic dis­ crepancy ratios, in terms of ORL versus UV, optical (OPT) and IR CEL abundances; for comparison, the previously studied NGC 7009 (LSBC), NGC 6153 (Liu et al. 2000), the galactic bulge PN M 1-42, M 2-36 (Liu et al. 2001), as well as NGC 6644 and 6572 (from unpublished ESO 1.52 m observations) are also in­ cluded. Possible exceptions regarding ionic discrepancies involve in NGC 3918, where values derived from the C III AA4187, 4650 ORLs and the C IV A1550 resonant doublet almost coincide; in the same nebula obtained from the

N IV A4606 ORL is 10 per cent less that the value derived from the N V A1240 CEL; also in NGC 6302, where identical values are derived from the

N III A4379 ORL and the N iv] A1486 intercombination line. This could be due to an underestimation on our part of the actual electron temperature pertaining to highly ionized species in these two high-excitation nebulae; adopting a higher Te in order to derive and would result in lower CEL abundances and produce a discrepancy with the ORL values as well. Alternatively, it could mean that at least in some nebulae the mechanisms that are responsible for the abundance discrepancies do not affect species of different ionization degree in the same way. This work has thus revealed a further two nebulae, NGC 2022 and LMC N66, which exhibit extreme discrepancies - a factor of 16 and 11, respectively, for

02+ / h + _ in the heavy element ionic abundance ratios deduced from ORLs as opposed to those from the traditionally employed CELs; these nebulae are added to a rare class of PN, along with NGC 6153 and Ml-42, that show similar, excessive discrepancies; we also expose another object, the Type-I nebula NGC 2440, which 2.6. Ionic Abundances from ORLs 106 shows discrepancy factors similar to those of NGC 7009 and M2-36. Regarding the and fractions plotted in Fig 2.2 that were derived from far-IR lines and taking into account that in Table 2.7 we presented both ‘low’ (adopting [O ill] 88 //m/52 fim densities) and ‘high’ values (adopting densities from [Ar iv], [Cl III]), we note the following: for IC 4406 the ‘low’ values are plotted for both and 0^+; for NGC 5315, 3918, and 3132, a median (half­ way) N^”*" value is plotted; for NGC 3242 and NGC 6302 the ‘low’ N^+ value is plotted, while for NGC 5882 and IC 4191 the ‘high’ N^"*" value is plotted; similarly regarding 0^+, for N G C 3242, 3918, 6302 and IC4191 the median IR value is plotted, while for NGC 5882, 5315 and 3132 the ‘high’ value is plotted. As we discussed in Section 2.5 following on the ideas of Rubin (1989), the far-IR N^”*” and 0^“^ lines have low critical densities so that their emission is biased towards lower-than-average nebular density regions, while the opposite is true for the UV and optical lines whose emissivity ratios to Hj3 are not acutely sensitive to variations in density (up to certain medium-high values, depending on the critical densities of the lines in question). The median IR N^'*' and abundances, as defined above, agree excellently with the values obtained from the N III] A1750 and [O III] AA4959, 5007 lines. This shows that even though density variations are present in the nebulae, in the medium density regime— which probably represents a weighted mean of random densities along the line of sight through the nebula—no significant bias in the inferred CEL abundances is present, since very satisfactory agreement amongst the derived values is achieved— but note also the discussion below regarding O III] A1663. Alternatively, this might mean that whatever bias there is, it affects all UV, optical and IR CEL abundances in a collective, rather similar manner. A matter of interest here, in the context of density variations and their ef­ fect on abundance determinations, is this: in 8 out of 11 PN for which both the

O III] A1663 and [O lll] AA4959, 5007 lines have been used to derive abundances, we find that on average the 0^+(A1663)/H"'" ratio is about 40per cent less that 2.6. Ionic Abundances from ORLs 107 the ratio. This could be construed as follows; in the presence of temperature fluctuations along the line of sight within PN, the A1663 line will be preferentially emitted from a hotter medium than the optical AA4959, 5007 and A4363 lines due to the dissimilar dependence of the line emissivities on Tg; there­ fore, the tacit assumption on our part that Te(A1663) = Tg(A4959 + A5007)/A4363

[— Te(opt)], which is the temperature derived from the nebular to auroral [O III] line ratio and is adopted to derive abundances, should result in 0^"""(A1663) being overestimated and possibly make it bigger than 0^+ (opt). The fact that the opposite is observed may be interpreted according to the model analysis of Mathis et al. (1998) as due to the existence of dense clumps within the nebulae at densities of ~ 10^- 10® cm"®. Within such dense condensations the nebular [O III] lines would be collisionally suppressed, leading to an overestimation of Tg(opt}, such that T(opt) > r(A1663); this could result in the observed 0^""'(A1663)/H+ abundance ratio being less than the 0^"*"(opt)/H''" ratio. In Section 2.7 we will look further into the arguments in favour of or against the existence of density and temperature fluctuations in PN as revealed from our observations. .3 C2+ C3+ C4+ N+ N2+ N3+ 0 + 02+ 03+Ne2+Ne3+ S+ S2+ CI2+Ar2+Ar3+Ar4+ 10 !0 o ! ^ i 10 □ é

10

10 □

-7 NGC 3242 10 r : O - Opt recomb, lines ■ ; A - UV coll. ex. lines -8 □ - Opt coll. ex. lines 10 # - IR coll. ex. lines

C2+ C3+ C4+ N+ N2+ N3+ 0+ 02+ 03+Ne2+Ne3+ S+ S2+ C12+Ar2+Ar3+Ar4+

C2+ C3+ N+ N2+ N3+ 0+ 02+ Nc2+ S+ S2+ C12+ Ar2+ Ar3+

10 <> ■ À o 0 10 4

10

-6 NGC 5882 10 0 - Opt recomb, lines ■ A . UV coll. ex. lines

- □ - Opt coll. ex. lines -7 # - IR coll. ex. lines 10

C2+ C3+ N+ N2+ N3+ 0+ 02+ Ne2+ S+ S2+ CI2+ Ar2+ Ar3+

Figure 2.2: Comparison of ionic abundances in PN derived from optical recombination lines, and from UV, optical and IR collisionally excited lines. In these and subsequent graphs the plotted values are for the entire nebulae, originating from ground-based scanning optical spectroscopy and wide aperture, lUE, ISO and IRAS observations, except for NGC 6302, NGC 6818, IC 4406 and My Cn 18, where abundances from ORLs and optical CELs have been derived from fixed-slit spectra. 2.6. Ionic Abundances from ORLs 109

C+ C2+ C3+ N+ N2+ N3+ 0 + 02+ 03+ Ne2+ Ne3+ Ne4+ S+ S2+ C12+ Ar2+ Ar3+

10

▲ 10

5 10 □ X

NGC 5315 10 O - Opt recomb. lines A - UV coll. ex. lines □ - Opt coll. ex. lines • - IR coll. ex. lines 10

C+ C2+ C3+ N+ N2+ N3+ 0 + 02+ 03+ Ne2+ Ne3+ Ne4+ S+ S2+ CI2+ Ar2+ Ar3+

C+ C2+ C3+ C4+ N+ N2+ N3+ N4+ 0+ 02+ 03+ 04+ Ne2+Ne3+Ne4+ S+ S2+ C12+Ar2+Ar3+Ar4+ 10 <> 0 ▲ 1 $ 10 ? : 4 <5> o * : è à 10 i

□ 10 NGC 3918 P i O - Opt recomb. lines A - UV coll. ex. lines □ - Opt coll. ex. lines 10 • - IR coll. ex lines

C+ C2+ C3+ C4+ N+ N2+ N3+ N4+ 0+ 02+ 03+ 04+ Ne2+Ne3+Ne4+ S+ S2+ C12+Ar2+Ar3+Ar4+

Figure 2.2: — Continued. 2.6. Ionic Abundances from ORLs 110

C2+ C3+ C4+ N+ N2+ N3+ N4+ 0+ 02+ 03+ 04+ Ne2+Ne3+Ne4+ S+ 82+ C12+Ar2+Ar3+Ar4+

10 O

10 À

À A 10

10

NGC 2022 P- O - Opt recomb. lines 10 A - UV coll, ex. lines □ - Opt coll. ex. lines

10 C2+ C3+ C4+ N+ N2+ N3+ N4+ 0+ 02+ 03+ 0 4 + Ne2+Ne3+Ne4+ S+ 82+ CI2+Ar2+Ar3+Ar4+

C+ C2+C3+C4+ N+ N2+N3+N4+ 0 + 0 2 + 0 3 + 04+Ne2-dSe3HNe4+S+ 82+C124Ar2+\r3HAr4+ 10 r <> Â 10 i À

10 a À A

10 NGC 6818 O - Opt recomb. lines A - UV coll. ex. lines □ - Opt coll. ex. lines • - IR coll. ex. lines 10 0-- C+ C2+C3+C4+ N+ N2+N3+N4+ 0 + 0 2 + 0 3 + 04+Ne2-flVe3-ff^e4+8+ 82+CI2+Ar2HAr3+\r4+

Figure 2.2: — Continued. 2.6. Ionic Abundances from ORLs 111

. 3 C2+ C3+ C4+ N+ N2+ N3+ N4+ 0+ 02+ 03+ 0 4 + Ne2+Ne3+Ne4+ S+ S2+Ar2+Ar3+Ar4+ 10 o : <> : <>

10 4

P

10

NGC 2440 0 - Opt recomb. lines 10 - ▲ - UV coll. ex. lines □ - Opt coll. ex. lines • - IR coll. ex. lines

C2+ C3+ C4+ N+ N2+ N3+ N4+ 0+ 02+ 03+ 04+ Ne2+Ne3+Ne4+ S+ S2+ Ar2+Ar3+Ar4+

C2+ C3+ N+ N2+ N3+ N4+ 0+ 02+ 03+ 0 4 + Ne2+Ne3+Ne4+ S+ S2+ C12+Ar2+Ar3+Ar4+ 10

▲ 10 * 4 <> é

10

10 NGC 6302 Q 0 - Opt recomb. lines ▲ - UV coll. ex. lines 10 □ - Opt coll. ex. lines • - IR coll. ex. lines

C2+ C3+ N+ N2+ N3+ N4+ 0+ 02+ 03+ 0 4 + Ne2+Ne3+Ne4+ S+ S2+ CI2+Ar2+Ar3+Ar4+

Figure 2.2: — Continued. 2.6. Ionic Abundances from ORLs 112

.3 C+ C2+ C3+ N+ N2+ N3+ 0+ 02+Ne2+Ne3+ S+ S2+ C12+Ar2+Ar3+ 10

10

IC 4406 10 O - Opt recomb. lines ▲ - UV coll. ex. lines □ - Opt coll. ex. lines • - IR coll. ex. lines □ 10 C+ C2+ C3+ N+ N2+ N3+ 0+ 02+Ne2+Ne3+ S+ S2+ CI2+Ar2+Ar3+ C2+ N+ N2+ 0+ 02+ Ne2+ S+ S2+ C12+ Ar2+

10

10

X 10

10 My Cn 18 O - Opt recomb. lines □ - Opt coll. ex. lines

10 C2+ N+ N2+ 0+ 02+ Ne2+ S+ S2+ C12+ Ar2+

Figure 2.2: — Continued. 2.6. Ionic Abundances from ORLs 113

C2+ C3+ N+ N2+ N3+ 0+ 02+ Ne2+ S+ S2+ C12+ Ar2+ Ar3+ 10 : O É ▲

10

10 ; : : ; ; ; ; Ù ■ : D -6 10 NGC 3132 - ; O - Opt recomb. lines A - UV coll. ex. lines -7 □ - Opt coll. ex. lines P : 10 i : # - IR coll. ex. lines □ :

C2+ C3+ N+ N2+ N3+ 0+ 02+ Ne2+ S+ S2+ C12+ Ar2+ Ar3+ C2+ C3+ N+ N2+ N3+ 0+ 02+Ne2+Ne3+ S+ S2+ C12+Ar2+Ar3+Ar4+

<> : 10 <> * i

é 10

10

10 IC 4191 □ o - Opt recomb. lines □ - Opt coll. ex. lines • • IR coll. ex. lines 10

C2+ C3+ N+ N2+ N3+ 0+ 02+Ne2+Ne3+ S+ S2+ C12+Ar2+Ar3+Ar4+

Figure 2.2: — Continued. 2.6. Ionic Abundances from ORLs 114

.3 C2+C3+ N+ N2+N3+N4+ 0+ 02+04+Ne2+Ne3+Ne4+ S+ S2+Ar2+Ar3+Ar4+ 10

10

5 10 X

10 LMC N66 O - Opt recomb. lines □ A . UV coll. ex. lines □ - Opt coll. ex. lines 10 C2+C3+ N+ N2+N3+N4+ 0+ 02+04+Ne2+Ne3+Ne4+ S+ S2+Ar2+Ar3+Ar4+

C+ C2+ C3+ N+ N4+ 0+ 02+ Ne2+ S+ S2+ Ar2+ Ar3+ 10 o <>

4 g 10 : <> 4 ▲ □ 10

: : i : 10 LMC N141 O - Opt recomb. lines 4 - UV coll. ex. lines □ 10 □ - Opt coll. ex. lines p -

C+ C2+ C3+ N+ N4+ 0+ 02+ Nc2+ S+ S2+ Ar2+ Ar3+

Figure 2.2: — Continued, 2.6. Ionic Abundances from ORLs 115

C+ C2+ C3+ N+ 0+ 02+ Nc2+ S+ Ar2+ Ar3+ 10 <> i

10

10

10 Û

SMC N87 10 O - Opt recomb. lines A - UV coll. ex. lines □ - Opt coll. ex. lines □

10 C+ C2+ C3+ N+ 0+ 02+ Ne2+ S+ Ar2+ Ar3+

Figure 2.2: — Continued. Table 2.18: Ionic ORL/CEL abundance discrepancies in PN.

PN 3%(C^+) %(C^+) %(N^+) %(N^+) %(N^+) %(0^+) %(0^+) 3%(Ne^+)

ORL/UV ORL/UV ORL/UV ORL/IR ORL/UV ORL/UV ORL/OPT ORL/UV ORL/IR ORL/UV ORL/OPT

NGC 3242 5.2 (5.5)“ * 7.8 9.0 * * 2.2 5.7 2.3 * 2.5

* * * * N G C 5 8 8 2 2 .9 1.5 1.6 2.1 2 .2 2.1 2.3 o- * * * * NGC 5315 3.4 1.8 2.1 2.0 1.4 1.7 1.3 i PL * D NGC 3918 2.1 (3.4) 1.1 2.3 2.4 9.5 0.9 1.8 2.3 1.9 8.7 3 NGC 2022 10.(9.1) 4.7 34. * 15. * 16. 23. * 1.3 33. CnR * * NGC 6818 2 .3 ( 1 . 3 ) 6.2 12. 14. 2.3 2.9 4.9 1.3 13.

* * * * N G C 3 1 3 2 4.4 3.4 3.1 2.4 3.5 2.0 *

* * NGC 2440 4.0(2.5) 4 .8 6.7 5.3 14. 5.4 8 .9 * 2.9 g Cn NGC 6302 6.5(3.3) * 5.5 3.1 1.0 2.4 3.6 3.5 4.7 * 36.

IC 4406 3.7 4.9 1.7 2.1 * * 1.9 1.8 1.8 * *

IC4491 * * * 3.7 * * 2.4 * 2.5 * 3.0

My Cn 18 * * * * * * 1.8 ** * *

SM C N 87 2.1 3.5 * *** 1.6 * ** *

LM CN66 * ** * ** 11. 19. ** *

LM CN141 4.6 2.9 * * * * 2.6 * * * *

NGC 4361 15. 4.0 * * ** ** * * *

NGC 6572 2.1 (1.1) * * * * * 1.5 * 3.9 * * * * * ** NGC 6644 2 .8 ( 2 . 8 ) ** * 1.9 * ** * * * NGC 7009 4.1 (5.3) 3.2 3 .2 7.1 5.0 5.9 * * * * * NGC 6153 9 . 0 ( 9 . 7 ) 3.3 9.0 9.2 7.3 11.

* * M l-4 2 23. ** 8 .2 22. * 11. * 17.

M 2-36 4.8 * * 5.6 * * 6.9 * 5.6 * 7.7

“ Parenthesized values are from Rola & Stasihska (1994). 2.6. Ionic Abundances from ORLs 117

2.6.7.2 C/O and N/O elemental ratios

In Tables 2.19 and 2.20 we present, respectively, the C/O and N/O ionic and elemental abundance ratios derived using pure ORL and pure CEL line ratios. In both tables the last column contains two values: the one before the vertical dash is computed with the total oxygen CEL abundance found if we take into account the O^+(A4959+A5007)/H+ ionic ratio in the ICF method; that after the dash uses the oxygen abundance found if we take into account the 0^'*'(A1663)/H''‘ ratio; in the case of the CEL N/O ratio the value after the dash is for both nitrogen and oxygen CEL abundances computed by adopting the N^"^(A1751)/H"'" and 0^"*"(A1663)/H'^ ionic ratios in the ICF method. Two exceptions are SMC N87 and LMC N141 for which the values after the dash, in both tables, include oxygen CEL abundances computed with the 0^"^(A4931)/H"^ ionic ratios. In Figure 2.3 we show the CEL C^'^'/O^'^ and N^"*'/0^'^ ratios compared with the corresponding ratios derived from ORLs. The revised criterion of KB94 states that Type-I PN are those objects that have experienced envelope-burning conversion to nitrogen of dredged-up primary carbon, i.e. PN with a N/O ratio greater than the (C-l-N)/0 ratio of Hll regions in the host galaxy, the latter being 0.8 for our own Galaxy. According to this criterion three objects in our sample, NGC 2440, NGC 6302 and My Cn 18, all bipolar, qualify as Type-I PN by virtue of their ORL N/O ratios; the first two qualify on account of their CEL N/O ratios too. 2.6. Ionic Abundances from ORLs 118

OSMCN8

0.4

ON2440

j 0.2 ON2022 ON68I8 !.. OLMCN141

u -0.2 ON3132 b£ >IC4406 o •• . -0.4 O - Xl908/k1663 Sn630^jvj53j5 ON5882 • -X.1908/(H959 + ;^5007) -0.6 L-L. -0.4 - 0.2 0.0 0.2 0.4 0.6 Log CIO^ (C II U 267/0 II ORLs)

...... - 1 ]— I ■ I —1— I I —1—7 1 f — N2440# # 0.0 -

N63020' -0.2 ;

-0.4 -

-0.6 0 " IC41910 Ç • • N6818 • z W) -0.8 O

-1.0 • - X1750/X1663

N32420 0 - 57gm/(52pm + 88gm) -1.2 1 . . . 1 ...... 1 -0.6 -0.4 -0.2 0.0 0.2

Log n"(N II/O II ORLs)

Figure 2.3: Comparison of the {top) and (bottom) ionic abundance ratios derived from optical recombination lines, and from UV, optical and far-IR collisionally excited lines. Data points with identical x-axis values share the same label. The dashed lines have a slope of unity. 2.6. Ionic Abundances from ORLs 119

Table 2.19: Carbon to Oxygen abundance ratios from CELs and ORLs.

PN C 2+ /o2+ C 2 + /q 2+ C 2+ /q 2+ ORL C/O CEL C/O A 4267/0 II A1908/A1663 A1908/A4959+A5007

NGC 3242 0.98 1.06 (0.66)“ 0.42(0.62)'’ 1.01 0.42|1.05 NGC 5882 0.39 0.30 0.27 0.44 0.31 NGC 5315 0.77 0.32(0.31) 0.45 (0.27) 0.58 0.34 NGC 3918 0.94 1.03(1.15) 0.83(0.90) 0.55 0.57 NGC 2022 0.68 1.52 1.08(1.21) 0.82 0.47 NGC 6818 0.64 1.35(1.45) 0.81(1.28) 0.70 0.49|0.64 NGC 3132 0.81 0.64 (0.63) 0.44(1.68) 0.81 0.48 NGC 2440 0.85 1.93(1.23) 1.16(0.94) 0.99 0.96 NGC 6302 0.61 0.33 (0.38) 0.34 (0.39) 0.75 0.31 IC 4406 1.09 0.55 0.58 1.25 0.62 IC4191 0.40 * * 0.45 * My Cn 18 0.59 ** 0.59 * SMC N87 3.91 * 3.04 4.31 3.49|4.64 LM CN66 * 0.13 0.08 * 0.10 LMCN141 1.38 * 0.77 1.48 1.0211.85 Solar * * ** 0 .7 2 “

“ Parenthesized values are from Rola & Stasihska (1994); Rola & Stasihska (1994) values are from the A1908/A5007 line ratio; With solar oxygen abundance taken from Allende Prieto et al. (2001) and solar carbon as in Table 2.17. 2.6. Ionic Abundances from ORLs 120

Table 2.20: Nitrogen to Oxygen abundance ratios from CELs and ORLs.

PN N 2 + / o 2+ N 2 + / o 2+ N 2 + / q 2+ N 2 + / q 2 + ORL N/O CEL N/O

N ll/O II A1750/A1663 A1750/A4959+A5007 57 /im/52+88 /im

NGC 3242 0.27 0.20 0.08 0.07 0.29 0.10|0.26 NGC 5882 0.23 0.25 0.23 0.22 0.29 0.33 NGC 5315 0.40 0.32 0.45 0.32 0.47 0.54|0.42 NGC 3918 0.26 0.27 0.22 0.21 0.19 0.23 NGC 2022 0.34 0.23 0.16 * 0.38 0.13 NGC 6818 0.58 0.22 0.14 * 0.58 0.16|0.21 NGC 3132 0.42 0.42 0.29 0.32 0.51 0.36|0.43 NGC 2440 0.91 1.21 0.73 * 0.99 0.74|1.09 NGC 6302 1.72 1.10 1.12 0.74 1.53 1.34 IC 4406 0.27 0.28 0.29 0.23 0.34 0.36|0.34 IC4191 0.39 * * 0.26 0.56 0.06|0.17 My Cn 18 3.16 * * * >1.99 0.39 SM C N 87 ** * * * 0.33|0.44 LMC N66 * 0.55 0.34 * * 0.41 LM CN141 * * * * * 0.45|0.83 Solar * * * * * 0 .1 9 “

With solar oxygen abundance taken from Allende Prieto et al. (2001) and solar nitrogen as in Table 2.17. 2.1. Discussion 121

2.7 Discussion

2.7.1 Correlations

2.7.1.1 Discrepancy factors versus temperature gradients

In Fig. 2.4 {top) we plot the discrepancy ratio of ionic abundances de­ rived from the C II A4267 ORL over those from the C III] A1908 CEL against the difference between the nebular electron temperatures derived from the [O III] nebular to auroral forbidden line ratio and from the nebular Baimer jump, for a sample of 16 planetary nebulae. The data for six PN are from other studies: NGC 7009 (LSBC), N GC4361 (Liu 1998; who quotes C^+(A1908)/H+ values from Torres-Peimbert, Peimbert & Pena 1990), NGC 6153 (Liu et al. 2001), M 1-42 and M 2-36 (Liu et al. 2001) and N G C 6644 (unpublished observations). We see that a very tight, positive linear correlation exists between

A(C2+/H+) = log(C2+/H+)oRL - log(C2+/H+)A1908CEL. and

A T = Te([0 III]) - Te(BJ).

A linear fit to the sixteen planetary nebulae plotted in Fig. 2.4 (top) yields,

A(C^+/H+) = (0.408 ± 0.053) -k (15.5 ± 2.1) x 10"^ AT, (2.5) with a linear correlation coefficient of 0.89. Fig. 2.4 {bottom) shows that a very similar linear correlation exists between A(02+/H+) = Iog(02+/H+)ORL - log(O2+/H+)A4959+5007CEL and A T too (cf. also Liu et al. 2001 for a similar result). A linear fit to the sixteen planetary nebulae of the bottom panel of Fig. 2.4, which includes IC4191 but not NGC 4361 yields,

A(Q2+/H+) = (0.224 ± 0.053) -F (21.6 ± 2.4) x 10"^ AT, (2.6) with a linear correlation coefficient of 0.92. 2.7. Discussion 122

In Fig. 2.5 we plot A(N^+/H+) = log(N^+/H+)oRL - log(N^+/H+)Ai 75iCEL, versus AT for twelve PN, where it is clear that a similar, positive trend is present as well. We have taken A (0^'^/H +) and A(N^+/H+) from Table 2.18, while AT is from Table 2.5. The slope of the A(0^'^/H'^) fit is steeper than that of A(C^’''/H'^) by a factor of 1.39 ±0.34, suggesting that the former abundance discrepancy ratio may be marginally more sensitive to changes in temperature than A(C^"'"/H'^) is; this may be giving us some clues as to the physical causes behind the abundance discrepancies (see Section 2.7.2). These readily apparent correlations of all three discrepancy ratios for doubly ionized C, N and O with the difference between the [O III] and BJ nebular tem­ peratures are significant, since they provide a strong observational confirmation of the importance of temperature variations, however those may be caused, on the frustrating problem of discordant—ORL versus CEL—abundance determinations.

2.7.1.2 Discrepancy factors versus PN intrinsic surface brightness

In Fig. 2.6 the A(0^'^/H+) and A(C^'*'/H‘^) discrepancy ratios (Table 2.18) are plotted against the nebular H/3 surface brightness, for samples of 21 and 19 PN respectively. Here the surface brightness, ^(H/?), is defined as the flux received per square arcsec of the nebula corrected for interstellar extinction. For all galactic PN we use nebular angular radii, H/3 integrated fluxes and logarithmic extinction coefficients, c(H/3)'^^'^, from CKS92; for the Cloud PN integrated fluxes are from Meatheringham et al. (1988); for LMC N66 we adopt an angular radius of 1.5 arcsec (Dopita et al. 1993), while for LMC N141 and SMC N87 we adopt 0.30 arcsec (Shaw et al. 2001) and 0.50 arcsec (upper limit value; quoted by Bofh & Stanghellini 1994), respectively. Discrepancy ratios for NGC 6572 are derived from ESO 1.52-m unpublished observations. We see that both discrepancy factors are clearly anticorrelated with the nebular 2 . 7. Discussion 123 surface brightness. A linear fit to the 21 PN plotted in Fig. 2.6 {top) yields,

A (0^+ /H +) = (-3.52 ± 1.41) - (0.327 ± 0.113) x log5(H/3), (2.7) with a linear correlation coefficient of —0.55; while a linear fit to the 19 PN of Fig. 2.6 (bottom) yields,

A(C^+/H+) = (-2.50 ± 1.03) - (0.253 ± 0.083) x logS(H ^), (2.8) with a linear correlation coefficient of —0.60. The surface brightness can be considered as an evolutionary parameter since for an expanding nebula S'(H^) decreases as the nebula ages.

2.7.1.3 Discrepancy factors versus PN absolute radii

In Fig. 2.7 the A(0^"*"/H+) and A(C^'*'/H+) discrepancy ratios are plotted against the nebular absolute radii. For all galactic PN we have used the absolute radii quoted by CKS92, except for NGC3132 where the radius was derived assuming a distance to the nebula of 600 pc (Sahu & Desai 1986) and an angular radius of 22.5 arcsec from CKS92. For the Cloud PN we adopted the following: for LMC N66 a radius of 0.324pc (Dopita et al. 1993); for LMC N141 a radius of 0.074pc, adopting a nebular angular radius of 0.30 arcsec (Shaw et al. 2001) and assuming a distance to the LMC of 50.6 kpc (Dopita et al. 1996); for SMC N87 we derived an upper limit radius of 0.141 pc, adopting an angular radius of 0.50 arcsec (upper limit value; quoted by Boffi & Stanghellini 1994) and a distance to the SMC of 58.3 kpc (Dopita et al. 1996). We see that a positive linear correlation exists between A(C^"'"/H"'") and the absolute nebular radii of 18 PN [Fig. 2.7 (bottom)] excluding SMC N87 whose radius is an upper limit], which can be fitted by,

A (C^^/H +) = (0.299 ± 0.119) + (4.84 ± 1.36) x R, (2.9) with a linear correlation coefficient of 0.67. No such correlation is clearly obvious regarding the A(0^"*"/H+) discrepancy factor [Fig. 2.7 (top)]. 2.1. Discussion 124

2.7.1.4 Temperature gradients versus PN radii and surface brightness

So far we saw that the abundance discrepancy ratios are: i) positively corre­ lated with the difference between the [O III] forbidden-line and BJ temperatures, ii) anticorrelated with the intrinsic nebular surface brightness, and iii) positively correlated with the absolute PN radii mostly]. Therefore it is not surprising to find that the aforementioned temperature difference is also, anti­ correlated with S(HP) and positively correlated with the absolute nebular radii (Fig. 2.8). A linear fit to the 25 nebulae plotted in Fig. 2.8 (left) yields,

A T = (-17.6 ± 6.1) X 10^ - (1517 ± 481) x log5(H/?), (2.10) with a linear correlation coefficient of —0.55. The relation between A T and the nebular radii for the same 25 objects [Fig. 2.8 (right)] can be fitted by,

A T = (-853 ± 744) + (28.1 ± 7.6) x 10^ R, (2.11) which has a linear correlation coefficient of 0.61. In Fig. 2.8 apart from the 11 nebulae whose BJ temperatures were measured in the context of this study, we also include the 6 objects published previously by Liu and colleagues (see Section 2.7.1.1 for details on these), along with a further

8 nebulae whose BJ and [O III] temperatures were presented by Liu & Danziger (1993b). We present the integrated, intrinsic H/3 fluxes, intrinsic surface bright­ nesses and the angular and absolute radii of all these PN in Table 2.21. 2 . 7. Discussion 125

Table 2.21: Total H/? intensities, intrinsic surface brightnesses, angular and absolute radii of PN.

PN log/(H/?) log5(H/3) e R

(cgs) (cgs) (arcsec) (pc)

NGC 3242 -9.59 -12.63 18.6 0.098

NGC 5882 -9.95 -12.14 7.0 0.057

NGC 5315 -9.85 -11.30 3.0 0.018

NGC 3918 -9.64 -12.08 9.4 0.046

NGC 2022 -10.59 -13.06 9.7 0.117

NGC 6818 -10.12 -12.54 9.1 0.082

NGC 3132 -10.12 -13.32 22.5 0.065

NGC 2440 -9.98 -12.91 16.4 0.107

NGC 6302 -10.55 -12.35 22.3 0.057

IC 4406 -10.45 -12.94 10. 0.114

IC4491 -10.29 -12.48 7.0 0.085

My Cn 18 -10.47 -12.57 6.3 0.096

SM C N 87 -12.43 -12.33 0.5 0.141“

LMC N66 -12.54 -13.39 1.5 0.324

LMCN141 -12.34 -11.79 0.3 0.074

NGC 6644 -10.57 -11.30 1.3 0.016

NGC 4361 -10.40 -14.11 40.5 0.173

NGC 6153 -9.67 -12.54 15.3 0.076

M l-42 -11.15 -12.87 4.1 0.109

M2-36 -11.18 -12.74 3.4 0.100

NGC 7009 -9.63 -12.43 14.1 0.082

NGC 2392 -10.26 -13.46 22.4 0.135

IC 2448 -10.75 -12.65 5.0 0.096

NGC 2792 -10.60 -12.72 6.5 0.095

NGC 2818 -11.14 -14.24 20.0 0.192

NGC 2867 -10.08 -12.38 8.0 0.071

NGC 3211 -10.76 -13.06 8.0 0.111

Me 2-1 -11.24 -12.69 3.0 0.096

NGC 6741 -10.30 -11.98 3.9 0.039

“ Upper limit derived from an upper limit 6 by BofR & Stangellini

(1994); 2. 7. Discussion 126

1.4

1.2 1 i 1.0 I o 0.8

S 0.6

0.4

0.2 0 1000 2000 3000 4000 5000 6000

Te([0 III]) - Te(BJ) (K)

Figure 2.4: The ionic discrepancy ratios for plotted against the difference between the [O ill] and BJ electron temperatures (top); and same for the O^'*’ ratios (bottom)] the dashed lines are the linear fits of Eqs. 2.5 and 2.6; see text for details. 2. 7. Discussion 127

, ...... 1 1 1 ■ 1 1 r~ ' ’— '— ' ■■ 1.6 a N2022

1.4 -

1.2 - ^N6818 5 1.0 - • a N3242 (► a N2440 0.8 î a N6302

i 0.6 .N 3132 s - (, N7009 ( ,N6153 &£ © 0.4 *N3918 -

• ^ « ’[^4406 0.2 0 N5882 ■ 1 1__ 1__ 1__ 1__ 1__ 1__ 1__ 1__ 1__ ■ ■__ 1 ■ ■ ■__ 1000 2000 3000 4000

Te([0 III]) - Te(BJ) (K)

Figure 2.5: Same as Fig. 2.4, but for the ORL to N^'*'(A1751)/H'^ CEL discrepancy ratio; arrows point to ORL/57-/xm values; see also Table 2.18. 2.7. Discussion 128

1.4 • M l-42

1.2 - N2022

• LMCN66 1.0 • N6153

• M2-36 'o 0.8 - . # N2440 I • N7009

0.6 - • N6302

W) • N68iK 0.4 - • LMCN141 2 • N3132 • IC4191 • N3242 • N5882 \ • IC4406^ MyCnl8 •N39I8 ^ ^ 0.2 • SMCN87 , N6572 * ‘ ‘ * « « « ■ > 1 É É î ■ iN. ■ 1.4 • M l-42

1.2 • N4361

1.0 I N2022 ' N6153 u ï 08 • N6302 X # N3242 • M&# • LMCN141 • N3132 'u 0 6 I N2440 • ^57009 • IC4406 s • N53k o • N5882 • N664 0.4 'N6818 \ • N3918 •N6572 ''x • SMCN87 0.2 -14.0 -13.5 -13.0 -12.5 -12.0 -11.5 -11.0 -2 -1 -2 Log S(HP) (erg cm s arcsec )

Figure 2.6: The discrepancy ratios of (top) and (bottom) ionic abundances derived from O ll and C II A4267 ORLs over optical and UV CELs, respectively, plotted against the nebular H/5 surface brightness. The dashed lines are the linear fits discussed in the text. 2.7. Discussion 129

—1—1—1—1—1—1—I—1—I—I—1—1—1—1 1 1—1—1—1—1—[—r 1 1 1 1 1 1 1 1 1 1 1 M l-42#

1.2 N2022#

LMCN66I 1.0 N6153#

M2-36# 0.8 N2440# N7009#

0.6 'g N6302# N6818# 0.4 LMCN141# 2 N3130#191#

ffia ji# * ® * N 3 9 1 8 . u , a W ^ 0.2 "N6572# SMCN87# . . . 1 .... 1 . . 0.00 0.05 0.10 0.15 0.20 0.25 0.30

*mU2 ■ ...... ' '

1.2 •N 4561

1.0 # N 2 0 2 2 /' •N 6I53 U I 0.8 •N 6302 «N3242

'u 0.6 , •N7009#n 2440 0£ •IC4406 o • N5315 / •N 5882 •N6644 0.4 ►N6818 0.3 ►N657?N3918 »SiyiCN87 _i I I 1 I 0.05 0.10 0.15 0.20 0.25 Radius (pc)

Figure 2.7: The discrepancy ratios of (top) and /YD (bottom) ionic abundances derived from O ll and C II A4267 ORLs over optical and UV CELs, respectively, plotted against the nebular radius. The dashed line is a linear fit to all data points, excluding SMC N87 whose radius corresponds to an upper limit value; see text for more details. 2 . 7. Discussion 130

6000

5000

4000 g ” 3000 a 2000 — .J»' O dn5334î * N5R82 • SÏ5S.18 • N6818 ^1000 ______•IC4Î^K • IC4191 • N5315 •N531^5^*N3918 0 • N6644’" ,X3132 0^2792^^^ mm? o N2792 OMe2-l O Me2-1

-1000 pN6741_ I _ ONy41 J I_ -14.5 -14.0 -13.5 -13.0 -12.5 -12.0 -11.5 -11.0 0.02 0.05 0.10 0.15 0.20 -2 -1 -2 Log S(HP) (erg cm s arcsec ) Radius (pc)

Figure 2.8: The difference between the nebular electron temperatures derived from the [O ill] nebular to auroral forbidden line ratio, Te([0 III ]), and from the nebular Baimer continuum discontinuity, Te(BJ), is plotted against the intrinsic nebular H/5 surface brightness (left) and absolute radius {right). Open circles data are from LD93b. The dashed lines are the linear fits discussed in the text. ^.7. Discussion 131

2.7.2 Temperature fluctuations

Peimbert (1967) and Peimbert & Costero (1969) first proposed that in the pres­ ence of temperature fiuctuations within nebulae, adoption of the [O ill] (A4659 -f A5007)/ A4363 line ratio as a standard thermometer will result in the underestima­ tion of elemental abundances as derived from CELs, since that line ratio is biased towards high-temperature nebular regions. Generally speaking, when variations in electron temperature exist along the line of sight within a nebula, abundance ratios derived from the ratio of a collisionally excited line to a recombination line (e.g. the most commonly used line H/5) will be underestimated, while those de­ rived from a pure recombination line ratio are almost insensitive to Te and largely unaffected by any incorrect assumption regarding the true nebular temperature. Peimbert (1971) found possible evidence regarding the existence of such tem­ perature fiuctuations, by comparing the temperature derived from the ratio of the

Baimer continuum over H/5 to that obtained from the [O III] nebular to auroral CEL ratio in three PN; he discovered that the latter temperatures were higher than Te(BJ). This has since been supported by further observations of more plan­ etary nebulae (Liu et al. 1995b). In this work we also present Baimer jump temperatures for 12 PN and find them lower on average than Te([0 III]) values (Table 2.5). The concept of temperature fiuctuations has been thus invoked several times as a promising cause of the ORL/CEL abundance discrepancy problem, pertain­ ing to both discrepant C^"^(A4267)/H"'" and ORL abundance ratios. For instance, Peimbert, Storey & Torres-Peimbert (1993) proposed that spatial tem­ perature fluctuations with an amplitude of ~ 20 per cent around a mean value would reconcile the factor of 1.55 discrepancy between the ORL and CEL 0^'^/H'^ abundances in the planetary nebula NGC 6572. Peimbert, Torres-Peimbert & Luridiana (1995) derived electron temperatures for a sample of nebulae from the

C III] A1908/C II A4267 ratio and showed that they are generally lower than

Te([0 III]). On account of this they argued that in order to explain the discrep- 2 . 7. Discussion 132 ancy between the C^+(A4267)/O^+(A5007) and C^+(A1908)/0^+(A5007) values in planetary nebulae published by Rola & Stasihska (1994), large temperature fluc­ tuations need to be called for; they also recommended the adoption of Te{C‘^'^) rather than 7A([0 III]) when deriving ionic abundances from CELs and proposed that due to their insensitivity to temperature changes, abundances derived from pure recombination line ratios are more reliable.

The validity however of the C III] A1908/C II A4267 ratio (or even that of the

C IV A1549/C III A4650 ratio) as a reliable thermometer is ambiguous, since there is evidence that at least in some PN that exhibit non-uniform, inhomogeneous abundances the C II A4267 ORL and C III] A1908 CEL do not originate in the same gaseous component; hence the derived temperatures do not have a physical meaning. Harrington & Feibelman (1984) argue that such is the case for the hydrogen-deficient material which exists at the centre of the planetary nebula A 30 in the form of condensations. We may well ask what evidence there is in our results about the role of tem­ perature fluctuations as a possible cause of the abundance discrepancies. The major premise of this theory is that ionic abundances derived from the intensity ratio of two lines with very different temperature dependencies—e.g. /([O III] A1663)//(H/3)—will be severely biased if an incorrect Tg is adopted in order to calculate the line emissivities. Since the realization that in PN, on average Te(BJ)

< Te([0 III]), and since the standard practice is to use the [0 III] optical line ratio temperature in order to calculate both the emissivities of CELs and of H/3, the possibility of systematically underestimating CEL abundances is there; should then ORL abundances that are insensitive to temperature bias be trusted more and if so do they point towards real heavy element overabundances in most neb­ ulae? One way to try to answer this question is first by examining the effects of changing Tg on ionic ratios derived from CELs having different temperature dependencies. As we mentioned in Section 2.5, in this study the standard temperature from ^.7. Discussion 133

the [0 III] (A4659 + A5007)/A4363 line ratio was adopted in order to derive all ionic abundances of doubly ionized species. If we suppose however that the Tg pertinent to the nebular zones where doubly ionized species exist is lower than that, e.g. that it may be as low as the BJ temperature, then the ionic abundances derived from CELs would all be underestimated by our adoption of Tg = Tg([0 III]), but not all by the same factor. Abundances derived from high excitation energy transitions, such as the O III] A1663 (E^gx — 86kK), N III] A1751 (Egx —82kK), or the C III] A1908 (Egx —75kK) lines, would be underestimated more relative to those derived from transitions that have a lower temperature sensitivity, such as the [O III] AA4959, 5007 nebular lines (Egx —29kK). Our results show however (Table 2.18), that all CEL C^*^, and abundances are systematically lower than those derived from ORLs, by rather similar factors. For instance, the average discrepancy factor C^"^ A4267/A1908 for eighteen nebulae is 5.4, while the 0^"^(0RL/A4959 + A5007) factor is 5.1; also the O^"*" ORL/A1663 discrepancy factor for ten nebulae is 5.7 (Table 2.18). This apparent uniformity cannot be reconciled with the different behaviour that one would expect amongst the discrepancy ratios in the case of temperature fluctuations of the type envisaged by Peimbert. We would expect the C^"^ ORL/CEL overabundance to be greater than that for O^"^, since Egx(A1908) > Egx(A4959 + A5007); however, in the extreme nebula NGC 2022 the discrepancy factor for C^"^ is 10, but that for 0^+ is 16, i.e. the opposite of that expected; in the other extreme nebulae NGC 6153 (Liu et al. 2000a) and M 1-42 (Liu et al. 2001) the corresponding factors are almost equal. Similarly, we would expect the C^+ ORL/CEL overabundance to be greater than that for C^"^, due to the higher excitation energy of the C IV A1549 resonant doublet (Egx —93kK), relative to that of C III] A1908 (Egx —75kK); however, this is not the case for the nine PN with both C^"^ A4267/A1908 and C^“^ ORL/A1549 ratios documented, where on average, these discrepancy factors are 5.3 and 3.9, respectively. Having just seen that in the presence of temperature fluctuations we would 2. 7. Discussion 134 expect A(C^+/H+), as defined in the previous section, to be affected more by errors in Tg than our fits (Eqs. 2.5 and 2.6) show that actually it may well be that the opposite is true, judging from the slightly steeper A(0^+/H"'") function (Fig. 2.4). This again may be suggesting a conflicting behaviour of the two discrepancy factors with that expected under simple temperature fluctuations. Probably, however, the most serious obstacle to the temperature fluctuation hypothesis lies with the IR abundance results. An inspection of Fig. 2.2 and Table 2.18 shows that the ORL/CEL discrepancy factors for CEL abundances obtained from the far-IR [O III] and [N III] lines are in line with those from the optical and UV CELs, i.e. there is no correlation whatsoever with the very low Egx (< 1000 K) of IR transitions. As we mentioned in Section 2.6.7.1, very satisfactory agreement is reached amongst UV, optical and IR CEL ionic ratios, while the same is also true for and Ne^'^/H’'"; this would not be the case in the presence of strong temperature fluctuations.

2.7.3 Density inhomogeneities

In section 2.4.2.1 the nebular electron densities derived from various diagnostic ratios were presented. It was shown that the low critical density [O III] 52- and 88-/im lines yield values that are, on average, a factor of 6 lower than those derived from the optical [Ar iv] and [Cl III] doublets, which have much higher critical den­ sities. Given that all [O III], [Ar iv] and [Cl III] lines arise from similar ionization regions^, this result suggests the presence of strong density variations within the nebulae. In sections 2.5 and 2.6.7.1 the effects of these variations on the and abundance ratios derived from the far-IR lines were discussed. Rubin’s (1989) examination of the effects of varying Ne on nebular abundances were men­ tioned; the general conclusion from Rubin’s work is that when there are variable

^The ionization potential of (= 35.1 eV), falls between those of CU (= 23.8 eV) and Ap''"

(= 40.7 eV). 2. 7. Discussion 135 density conditions, line ratios relative to H/3 (or any recombination line with a volume emissivity oc generally underestimate the true ionic abundances. We saw that in accord with these theoretical predictions, abundance ratios from far-IR lines are, in general, underestimated relative to UV and optical values (that are subject to smaller bias due to their higher Ncr), when the ‘low’ [0 III] 52/zm/88/xm densities are adopted for the calculations (Table 2.7). On the contrary, we saw that when higher N^s are adopted (intermediate with or equal to the [Ar iv] and

[Cl III] densities), the inferred abundances from the far-IR lines become compat­ ible with those derived from the UV and optical indicators (Fig. 2.2); with the exception of the high critical density [Ne III] 15.5-/im line which is not affected by such modest variations. We can therefore draw the conclusion that in the density regime probed by the IR and optical density diagnostics used here, satisfactory agreement is found among the abundances derived from UV, optical and IR CEL indicators. One interesting exception, is that of the disagreement between the 0^"""(A1663) and 0^'*‘(A4959 -f A5007) abundances which was mentioned in section 2.6.7.1. The conclusion is that much stronger density contrasts than those inferred from the far-IR and optical diagnostics may need to be called for in order to explain this result; the existence of dense clumps {Nq ~ 10^ - 10^ cm“^) embedded within a lower-density medium could be responsible for this result. In all cases however, and even after accounting for modest density variations, the CEL abundances (from UV, optical or IR lines) remain consistently lower than the ORL values. As was mentioned in Chapter 1, section 1.3.3, Viegas & Clegg (1994) showed that dense clumps with Nq > 10^ cm~^ can have a substantial effect on the [O III] forbidden-line temperature via collisional suppression of the nebular

AA4959, 5007 lines; in this way the observationally derived T^([0 III]) would be significantly overestimated and the CEL abundances accordingly underestimated. In their analysis of the extreme nebula NGC 6153, Liu et al. (2001a) discuss the case for the existence of dense condensations of a very small filling factor as 2. 7. Discussion 136 a potential cause of the ORL/CEL abundance discrepancies. They conclude that such dense clumps would have a dissimilar effect on the various CEL abundances. In particular, they find that in the increased density environment of the conden­ sations, the IR and optical abundance indicators would in the best of cases yield significantly higher abundance ratios, thus covering part of the distance from the ORL values; on the other hand however, the UV CELs (which possess critical den­ sities much higher than the far-IR and optical lines) would become so enhanced, so as to overcome the ORL abundance results. The conclusion in this case being that the existence of plain dense condensations by itself, is not a sufficient solution to the problem, since it requires the discrepancies to be correlated with the critical densities of the various CELs, something not indicated by the observations. The newly exposed extreme nebula NGC 2022 merits some attention in this regard. In that nebula, the O^^ ORL/OPT discrepancy is a factor of 1.4 less than the ORL/UV discrepancy and a factor of 2 less than the ORL/UV discrepancy

(Table 2.18). The O III] A1663 and N III] A1751 lines have much higher critical densities (~ 10^^ cm"^) than the optical [O III] nebular lines. In this case therefore it could be that within dense condensations of a small filling factor, the various UV and optical CELs abundance indicators would reach better agreement with the ORL indicators, something apparently not feasible in the case of NGC 6153; i.e. some kind of correlation of the discrepancies with the critical densities of the and N^'*' CELs may be present. Similar arguments may apply in the case of NGC 2440 and NGC 3242 (see Table 2.18). A more detailed analysis however will have to await another opportunity. Instead, in this thesis we present models for NGC 5882 that incorporate both temperature and density variations through an empirical approach. That nebula displays mild ORL/CEL abundance discrepancies (about a factor of 2 for most ions; Table 2.18) that have been extensively documented from UV, optical and IR observations. Its optical spectrum is of very good quality and high S/N ratio, exhibiting very prominent C, N, 0 , Ne ORLs (see Fig ??). We also possess f.7. Discussion 137 high resolution spectra of the region near the Baimer discontinuity that allowed the accurate recording of the high-order Hi Baimer lines, which are sensitive to ionized high-density condensations (Chapter 1, section 1.2.2). The next Chapter discusses these issues and the limitations presented by both our observations and the models. Ch a pter 3

Empirical Composite Models of the Planetary Nebula NGC 5882

3.1 Introduction

A code was developed by P. J. Storey and first used by Liu et al. (2000) to empirically model the emission line spectrum and chemical abundances of the ‘unusual’ PN NGC 6153. That object belongs to a rare class of planetary nebulae that exhibit extreme discrepancies (a factor of 10 in this case) in the heavy ele­ ment ionic ratios deduced from ORLs as opposed to those from the traditionally employed CELs. Other relevant objects include NGC 2022, 2440, LMC N66 (ex­ posed through this work), NGC 7009 (LSBC) and the galactic bulge PN M 2-36 and M 1-42 (Liu et al. 2001b). The code basically performs an elaborate optimization process; the input con­ sists of available spectroscopic information on the nebula, in terms of observed line intensities and the standard derived thermal/density diagnostics, as well as elemental abundances and relative abundance ratios. The user defines a physical nebular model, i.e. specifies initial guesses regarding Tg, Nq and abundances, and commences the optimization. These initial conditions are then used as a basis

138 3.2. NGC 5882: overview of nebular properties 139 for a minimization process that searches for the optimal values of the observed parameters. In practice, the routine minimizes the sum of the squares of the frac­ tional differences between prediction and observation for all input lines, except H 11-H24. The latter are given equal weight and a combined weight of unity in the optimization procedure (lines that are affected by blending problems can be excluded from the fitting). The choice of input lines is guided by the need to include important CELs and ORLs, as well as lines arising from more than one ionization stage, in order to treat the effect of recombination excitation (e.g. on lines such as [O ll] AA7320, 7330 and C ill] A1908). The code does not involve a consistent model of the nebular physics (e.g. such as in photoionization modelling), so the results can be physically unrealistic. The model solutions however offer valuable insight into the potential physical realities that may need to be invoked in order to explain the picture that emerges from the observations. In this chapter we will present empirical models for the PN NGC 5882. In Figs. 3.1, 3.2, 3.3, we present the lUE large-aperture SWF spectrum and the optical spectrum of the nebula from the blue atmospheric cutoff to 4960 Â. In the spectrum of Fig. 3.2 a broad emission feature (FWHM = 16 Â) has been identified as a O IV 3409.70 Â (multiplet 3) line from the central star of the PN. This may indicate that the nucleus of NGC 5882 is hydrogen-deficient (Mendez 1991).

The lUE spectrum of the nebula (Fig. 3.1) shows also a very high excitation O V line from the central star’s wind.

3.2 NGC 5882: overview of nebular properties

In Table 3.1 and Table 3.2 we summarize for reference the physical parameters of NGC 5882 of importance to the analysis that follows; these were derived using a variety of techniques from UV, optical and IR transitions (Chapter 2, Table 2.7). In Fig. 3.4 we plot the observed intensities of high-order Baimer lines (n —> 2, for n = 10, 11,...,24) versus the principal quantum number, n, of their upper 3.2. NGC 5882: overview of nebular properties 140

6

5 He II

CIV CHI]: ] stellar O V stellar 4 stellar

OIV] OIII] NIH]

S 2

1

1200 1300 1400 1500 1600 1700 1800 1900 X(A)

Figure 3.1: The lUE large-aperture, SWP spectrum of NGC 5882 from 1200 to 1950 Â.

levels. The Baimer decrement intensities become sensitive to the nebular electron density for n > 10, so they are a valuable diagnostic tool. In the same figure we have overplotted predicted Baimer intensity progressions for varying values of TVe using emissivities from Storey & Hummer (1995). In the case of NGC 5882 the observed Balmer decrement points towards the existence of an Hi emitting medium with an Nq less than IC^cm”^; this point is discussed further in the context of the model solutions that follow. A list of selected lines is given in Table 3.4, where the intensities are in units where /(H/5) = 100. The measurements are from the whole nebula, except for the high order Balmer lines, where high-resolution data taken with a fixed slit are used.

The intensities of the [0 III] far-IR lines are from Liu et al. (2001), while that of

[Ne III] 15.5-/im is from Pottasch et al. (1984); these have been expressed in units of /(H^) = 100, adopting an integrated dereddened flux of log/(H/5) = —9.95. For all models we adopt the following relative abundance fractions: for oxygen 3.2. NGC 5882: overview of nebular properties 141

Table 3.1: Plasma diagnostics of NGC 5882.

Diagnostic 7 ;(K )

[0 in] (A4959 + A5007)/A4363 9400

[0 III] (88/xm + 52/im)/(A4959 +A5007) 7450

[Ne III] 15.5/xm/3868 8850

[N II] (A6548 + A6584)/A5754 10 800“

[0 II] (A7320 + A7330)/A3727 15 300“

B J /H ll 7800

To 8270

e 0.034

Ne (cm~^)

[0 III] 88/im/52/im 1170'’

[Ar IV] A4740/A4711 5000

[Cl III] A5537/A5517 2700

[0 II] A3729/A3726 4750

[S II] A6731/A6716 4000

Balmer decrement 4430

“ Neglecting any recombination excitation contribution

to the component lines

Table 3.2: Ionic abundances of NGC 5882 derived from ORLs and CELs.

ORLCEL %

10^ X C^+/H+ 37.7 12.9u 2.9

10^ X N^+/H+ 16.3 ll.O u 1.5

10.3r 1.6

10^ X 0^+/H+ 97.0 43.5u 2.2

47.2o 2.1

47.0r 2.1

10^ X Ne^+/H+ 31.9 13.10 2.4

14.7r 2.2

The symbols u, o, r denote, respectively, values derived from UV, , optical and IR lines. s.2. NGC 5882: overview of nebular properties 142

0.6 OIV s te lla r

3100 3150 3200 3250 3300 3350 3400 3450 3500 3550

1

0.1 [O 11] [Neill] [Neill]

3550 3600 3650 3700 3750 3800 3850 3900 3950 4000 4050 X(A)

Figure 3.2: Log-inten.sity, fixed-slit spectrum of NGC 5882 from the blue atmospheric cutoff to 4050A, showing the broad stellar O IV 3409.16 (M3) line, the weak He I discontinuity at 3421 A (upper graj)h), the Hi Balmer discontinuity at 3646A and the interstellar Ca ll K line at 3934 À (lower graph). High-order Hi Balmer lines up to H 24 are resolved. Apart from the O IV and Ca II K lines, all other lines present are of nebular origin. The spectrum has not been corrected for extinction and is in intensity units such that F(H/3) = 100.

we have 0"*"/0 = 0.028, 0^"^/0 = 0.962 and 0 ^ + /0 = 0.01, while for carbon we adopt C^+/C = 0.851 and C^"""/C = 0.119. We further adopt N e^+/0 — 0.267 and A r^+/0 = 0.0014 throughout. The H 14 line is excluded from the minimization due to serious blending problems, while the remaining 13 hydrogen Balmer lines are treated as mentioned above. The optimal value of the sum for each of the empirical models is given in Table 3.3, where it is referred to as 3.3. Results and discussion 143

3.3 Results and discussion

We first model the emission by assuming a homogeneous nebula of uniform elec­ tron temperature and density and treat these two quantities plus the 0/H, He/H and C /0 ratios as free parameters. The resultant model line intensities are listed in Table 3.4 while the derived physical parameters are presented in Table 3.3, under the label YO. The best-fitting electron temperature and density are 8299 K and 5523 cm"^ respectively, while 0 /H = 8.94 x 10“^^, He/H = 0.108 (He II lines are not included in the fitting, hence only the He^/H”** fraction is determined) and C /0 = 0.51. An inspection of Table 3.4 shows that apart from the Balmer jump and the high-order Balmer lines, the differences between the predicted and observed intensities are generally larger than the expected uncertainties in the measurements. For instance, the [O III] A4363 and 88-//m lines are underesti­ mated by about 30 and 50 per cent respectively, while the Ne II A4392 perm itted recombination line is predicted to be weaker than observed by in excess of 40 per cent, all ample reasons to rule out the model as physically unsound. In model Y0 we simulate the effects of temperature fluctuations in the nebula by assuming that the nebular medium has a Gaussian distribution in InTg. We define x — logTe, so that the fraction of the material, f{x) in the temperature range dx is

f{x)dx = — î^ e x p dx o-x V ^ CTx )

In this way the temperature distribution in the nebula is characterized by xq and (Jx, for which we find xq = 3.906 and

Table 3.3: Parameters of empirical models.

Parameter YO Y 0 YM Y E l Y E 2

Component 1

Ne (cm “ ^) 5523 5426 6521 1656 1878

?;(K) 8299 8054 9397 8913 8610

O’lo g T e 0 0.07 0 0 0

10^ X O/H 8.94 9.29 4.64 7.39 9.25

H e/H 0.108 0.107 0.085 0.096 0.109

C/O 0.51 0.47 0.37 0.46 0.51

filling factor 1.0 1.0 0.325 %1.0 « 1 .0

fr a c (0 )“ 1.0 1.0 0.260 0.993 0.995

Component 2

Ne (cm “ ^) - - 265 3.61 X 10^ 4.15 X 10^

Te(K ) - - 690 7423 7876

10^ X O/H -- 419 13.4 as C .l

H e/H - - 1.76 0.120 as C .l

C/O-- as C .l as C.l as C.l

filling factor -- 0.791 1.77x10“^ 2.11 xlO“^

x ' " 0.782 0.716 0.050 0.304 0.366

“ Fraction of the total number of oxygen atoms in Component 1;

^ Goodness-of-fit measure (see text for definition); almost 30 per cent too. The prediction for the Balmer discontinuity is worse due to the presence of material at a temperature much lower than the average. Fi­ nally, the metal ORLs are severely underestimated even in this solution. Again the differences cannot be accommodated within the measurement errors. Next we consider a composite model (YM) which invokes the possibility that the nebula might be made of two components that differ in Tg, iVg as well as in elemental abundances from each other. The model is characterized by two electron temperatures and densities, two oxygen abundances and a volume filling factor for each component. We further assume that in the first component the He/H fraction is about 0.1, while the helium abundance is a free parameter in 3.3. Results and discussion 145

Table 3.4: Comparison of observed line intensities with those from empirical models".

Line A(Â) /(A) YO Y 0 YM Y E l YE2

H 11 3770.6 3.66 3.97 3.96 3.98 4.19 4.22

H 12 3750.1 3.05 3.06 3.05 3.07 3.30 3.33

H 13 3734.3 2.29 2.41 2.41 2.42 2.67 2.70

H 15 3712.0 1.61 1.59 1.59 1.60 1.84 1.86

H 16 3703.9 1.24 1.33 1.33 1.33 1.56 1.58

H 17 3697.1 0.96 1.12 1.12 1.13 1.34 1.35

H 18 3691.5 0.84 0.97 0.97 0.97 1.16 1.17

H 19 3686.8 0.74 0.84 0.84 0.84 1.01 1.02 H 20 3682.8 0.63 0.74 0.74 0.74 0.89 0.90

H 21 3679.5 0.58 0.65 0.66 0.65 0.79 0.80

H 22 3676.3 0.53 0.59 0.59 0.58 0.71 0.71

H 23 3673.7 0.49 0.53 0.53 0.53 0.64 0.64

H 24 3671.2 0.46 0.48 0.48 0.48 0.58 0.58

BJ/H /? 3646 0.48 0.49 0.51 0.49 0.50 0.49

He I 4471.5 5.65 5.57 5.55 5.33 5.59 5.64

He I 5875.7 16.0 15.84 15.81 16.3 15.8 15.94

He I 6678.2 4.31 4.41 4.41 4.48 4.41 4.44

C II 4267.2 0.40 0.42 0.40 0.41 0.43 0.42

C III] 1908 31.4 30.0 31.3 30.6 28.1 30.0

[O II] 3726.0 10.2 12.1 11.9 10.7 8.95 8.32

[O II] 3728.8 4.99 5.62 5.55 5.10 5.65 5.73

O II 4089.3 0.15 0.10 0.11 0.14 0.12 0.10 ** * O II 4661.6 0.12 0.12 0.11

[O II] 7320.0 0.94 1.00 1.02 1.03 1.05 1.03

[O II] 7329.7 0.86 0.81 0.82 0.82 0.84 0.82

[O III] 4363.2 5.57 3.99 4.25 5.29 4.91 5.14

[O III] 5006.8 1057. 1189. 1136. 980. 1008. 986.

[0 III] 52 /xm 152. 172. 176. 155. 167. 177.

[O III] 88 /zm 52.6 26.4 27.1 52.4 45.7 45.0

Ne II 4392.0 .034 .019 .020 .034 .028 .025

[Ne III] 3868.8 93.1 86.2 82.5 99.0 107. 108.

[Ne III] 15.5 iim 223. 279. 283. 210. 215. 223.

[Ar iv] 4711.4 2.61 2.52 2.53 2.46 2.39 2.45

[Ar iv] 4740.2 3.00 3.09 3.08 3.14 3.21 3.15

Mean difference (per cent) 14.02 13.02 4.20 9.16 9.53

“ Boldface type indicates, for representative lines, the most discrepant results

amongst the various models; 3.3. Results and discussion 146 the other component. We assign initial values to Tg, and 0/H for the first component similar to those diagnosed by the CELs. The best fit is obtained with the ‘normal’ component occupying 32.5 per cent of the nebular volume and having an electron density of 6521 cm“^, a temperature of 9397 K and 0/H = 4.64 x 10~'^, i.e. very similar to those obtained from the standard CEL analysis. The remaining 67.5per cent of the nebular volume is occupied by metal-rich (0/He = 0.024), hydrogen-deficient (He/H = 1.76), rarefied and cool material with Nq = 265 cm~^ and Te = 690 K. The agreement with observation in this case is rather excellent for most lines. The average percentage difference between predicted and observed intensities is 6.2 per cent. According to this scenario, the majority of CEL emission originates from the hotter, denser component (apart from the [O III] far-IR lines which are collisionally suppressed there), while most of the heavy-element ORL emission comes from the rarer, cooler medium. This model reproduces particularly well the [Ne III] 15.5-/im line, since the temperature in the metal-rich region is now low enough, so that the line is not efficiently excited. The total oxygen abundance is 3.5 times solar, while the C/O ratio is a factor of 2 less than solar and the total He/H fraction is 0.136. In model YM almost all of the HI emission (96 per cent) comes from component 1 that occupies only about a third of the volume of the nebula; the remaining 79 per cent of the ionized volume is extremely H-deficient, but contributes 74 per cent of all the heavy elements (Table 3.3). However, such a nebula would be far from pressure equilibrium, since due to the much lower temperature and density of the H-deficient component with respect to component 1, there exists a very large pressure contrast—a factor of 335—between the two regions, which is clearly not easy to support; diffusion and collapse seem inevitable. Similar reasons have been invoked in order to reject an equivalent model (IH2) for N G C 6153 considered by Liu et al. (2000). Model Y Sl (Table 3.3) contains high density inclusions with enhanced helium and heavy element abundances compared to the remaining material which occupies most of the nebular volume. With this model we investigate the possibility whether 3.3. Results and discussion 147 high-density inclusions of the type envisaged by Viegas & Clegg (1994) could be of practical value regarding the abundance discrepancy problem. As we already saw, for NGC 5882 simple temperature fluctuation models fail to provide a consistent picture for several important lines, far-IR CELs and metal ORLs. We therefore start the optimization with a high Ne in the second component, while we also reduce the initial abundance contrast between the two nebular media. Apart from the strongest 3d-4f 0 II transition, at A4089.3, for this model we also include in the fitting the multiplet V 1 A4661.6 ( J = 3/2-3/2) line, whose intensity £ls it turns out is reproduced particularly well (this choice of line was guided by its strength and unblended nature, making it an attractive alternative to other strong, but blended, V 1 lines such as the AA4649, 4650 transitions). The best fit for this model (YSl) is obtained when the second component occupies a fraction of only 1 . 7 7 xlO“^ of the nebula’s volume and has an elec­ tron density of 3 6 1 0 0 0 cm"^, an electron temperature of 7 4 2 3 K, He/H = 0 .1 2

(i.e. slightly helium-rich) and 0 /H = 1 3 .4 x 1 0 " " ^ . The remaining material has

Nq = 1 6 5 6 cm“ ^. Te = 8 9 1 3 K, He/H = 0 .1 0 and 0 /H = 7 .3 9 x 1 0 “ “^. This model gives a much better fit to the far-IR lines than the temperature fiuctuation (YO) model, as well as a considerably better fit to the optical [O III] lines and permitted metal ORLs too. The high critical density [Ne III] 15.5-/im line behaves properly, being mostly emitted from the rare medium, while suppression of it within the dense inclusions keeps its overall flux at acceptable levels. Overall, this model simulates well the effects of Viegas & Clegg-type density inhomogeneities. The auroral [O III] A4363 is predominantly emitted by the dense inclusions, while the opposite is true for the nebular AA4959, 5007 lines which are collissionaly suppressed there (it is worth pointing out that this is already occurring even though the electron density within the condensations is only about half the critical density of the [O III] nebular lines; Ncnt = 6.9 x 10^cm“^). As a result, the nebular to auroral [O III] electron temperature obtained from the model for each component is actually lower than that derived from the observed 3.3. Results and discussion 148 line ratio for the whole nebula. The enhanced total nebular oxygen abundance resulting from the lower temperature is 7.41 x 10~'^, i.e. 1.5 times s o la r b u t 51 per cent higher than the value deduced by the standard CEL analysis and 26 per cent lower than that of the ORL analysis. The total nebular He/H abundance is 0.096. It is quite pleasing that the model succeeds in fitting both the Baimer jump and the high Baimer lines, without encountering the difficulties of the models of NGC6153 by Liu et al. (2000). Obviously the smaller difference between the electron temperatures derived by the observed B J/H ll and [0 III] ratios helps in this regard. As this solution dictates, the dense inclusions contribute 55 per cent of the high Baimer line fluxes and 52 per cent of the Baimer jump (even though the total volume occupied by the inclusions is small, the high density enhances recombination rates sufficiently for them to emit strongly in the Hi lines). It is thus probable that the picture provided by Fig. 3.4 fails to highlight a density dis­ tribution like the one predicted by model YSl. The NGC5882 ionized inclusions in this case are not of a density as high as that of the model IH3 considered by Liu et al. for NGC6153, where was 2.15 x 10^cm~^; they are only 0.17 times as dense as the NGC6153 dense inclusions—albeit 517 times more voluminous— and may have escaped ‘detection’ in the graph taking into account their relative contribution to the emission measure. Furthermore, the smoothing effect of the long-slit observations throughout the PN’s volume and their relatively low spatial resolution, along with the small filling factor of the inclusions, probably assists in this interpretation as well. Amongst the electron densities derived straight from the observations it seems that the one given by the [O III] 88/rm/52/rm ratio (A^ = 1170 cm“^) is the most unbiased and closest to the value predicted by the model for component 1 (cf. Tables 3.1 and 3.3). An attempt to explore the density inhomogeneities scenario further is taken up in model YE2. A high density for the second component is still pursued

^ Adopting the Allende Prieto at al. (2001) solar oxygen abundance of 4.90 x 10 otherwise it is only 10 per cent higher than the Grevesse & Sauvai (1998) solar value. 3.3. Results and discussion 149 here, but this time we favour a more chemically homogeneous nebula and start the optimization with an oxygen abundance close to that deduced observationally from the O II ORLs. The solution is thus defined by one 0/H ratio, a common He/H fraction, plus the corresponding Tg and Ne of the two components. The best fit turns out to be only slightly worse than that of the previous model, but is better than the first three models. The resulting nebula exhibits a temperature contrast of only 730 K between the rarefied component and the embedded dense inclusions, which is much more agreeable compared to that of model YM (AT = 8707 K). The inclusions have now a slightly higher electron density of 4.15 x 10^cm“^, compared to YSl, whereas the ambient medium has very much the same density as in model YSl. The model reproduces very well the [O III] nebular and auroral line intensities, the far-IR CELs and the Baimer discontinuity, while it is best in fitting the [Ne ill] 15.5//m line among the solutions reached thus far. The total nebular oxygen abundance is 9.25 x 10"'^, i.e. 1.9 times solar (with respect to the solar oxygen abundance of Allende Prieto at al. 2001) and quite importantly is only 8 per cent lower than that derived based on observations of O II ORLs. We have therefore demonstrated that for a nebula which displays mild ORL/ CEL abundance discrepancies, a satisfactory fit to the various plasma and abundance diagnostics can be achieved, by adopting a composite nebular model that has high-density, small filling-factor inclusions embedded in a lower density compo­ nent. Without resorting to temperature fiuctuations, the model increases the total oxygen abundance via a reduced electron temperature for the low density, large filling factor medium. The high-density knots are only slightly more helium-rich than the rest of the nebula, in contrast to the more severe H-deficiency invoked for the hypothesized inclusions of the extreme PN NGC 6153 (cf. the composite models of Liu et al. 2000). Still the postulated inclusions have evaded detection by the high-order Baimer lines density diagnostic, but this may be due to the limited spatial resolution of our observations and the very small filling factor of high-density condensations. 3.3. Results and discussion 150

It remains to be seen, perhaps through self-consistent photoionization mod­ elling of the nebula, whether NGC 5882 is truly oxygen-rich, compared to the planetary nebula average and the sun. 3.3. Results and discussion 151

0.5 C 111+ O II 0.4

0.3

OO uO Z 0.2

.15 [SII] He I He I He I He II

4000 4020 4040 4060 4080 4100 4120 4140 4160 4180 4200 4220 4240

0.4

o :

• ■ .\v .1/• ■ Hy [O HI] He I 0.1 .1 1. 11...... Il 4240 4260 4280 4300 4320 4340 4360 4380 4400 4420 4440 4460 4480

0.4 C HI+O II

[ArlV]N IH+O II [Fe HI] [ArlV]N

4480 4500 4520 4540 4560 4580 4620 4640 4660 4680 4700 4720

0.4

[O HI]

[Ar IV] ghost

4720 4740 4760 4780 4800 4820 4840 4860 4880 4900 4920 4940 4960 X(A)

Figure 3.3: Log-iutensil,y spectrum of NGC 5882 from 4000 to 4960 Â featuring the prominent recombination lines from C, N, O and Ne ions; it was obtained by uniformly scanning the entire nebular surface using a narrow long-slit. The dotted lines show the adopted continuum level. 'Fhe steej) rise of the continuum after 4880 A is an instrumental artifact. The intensity is in units such that F(H/3) = 100. s. 3. Results and discussion 152

0.75 NGC 5882 HIO

0.50 éH 14

— 0.25

U 0.00 (cm) H17

-0.25

H23 -0.50 1.0 1.1 1.2 1.3 1.4 Log n

Figure 3.4: Reddening-corrected intensities [in units where F(H/3) = 100] of high-order H I Baimer lines (/?- —> 2, for n = 10, 11,....24), versus the principal quantum number, n, of their upper levels. II 14 at 3721.94 À is blended with the [S ill] A3721.63 line. The different curves plot the predicted Balmer decrements for electron densities from 100 to 10® cm“^, assuming an electron temperature of 7800 K, as derived from the BJ/H 11 ratio. The data can be fitted with a uniform electron density of 4430 cm"^. C h a p t e r 4

Elemental Abundances in Galactic

and Magellanic Cloud H II Regions

In this Chapter I present an abundance analysis of two Galactic and three Magel­ lanic cloud H II regions involving both collisionally excited lines and recombination lines of C II, N II and O II. A comparison of the two types of abundances is made. Complications in the emission line analysis arising from the presence of bright, scattered stellar continuum within the nebulae are pointed-out.

4.1 Observations

The observational dataset consists of long-slit spectra obtained during runs at the European Southern Observatory (ESQ) using the 1.52-m telescope and the 3.5-m New Technology Telescope (NTT). Additional long-slit spectroscopy for one target was performed at the 3.9-m Anglo-Australian Telescope (AAT). The journal of observations is presented in Table 4.1.

The Galactic H II regions M 17 and NGC 3576 were observed at ESQ with the 1.52-m telescope equipped with a B&C spectrograph. The detector was a Loral 2048 X 2048, 15^m x Ibfim CCD in July 1996 and a Ford 2048 x 2048, 15/im x

153 4-1. Observations 154

Table 4.1: Journal of observations.

H II region Date A-range FWHM PA RA DEC Exp. (UT) (A) (A) (deg) (2000) (2000) (sec)

ESC 1.52-m

M 17 0 7 /7 /9 6 3995-4978 1.5 -21 18 20 40.0 -16 09 29 3 x1200

M 17 1 3 /7 /9 6 3535-7400 4.5 -21 ” 60, 300, 600

NGC 3576 1 1 /2 /9 7 3995-4978 1.5 -50 11 11 56.9 -61 17 25 6 x1800

AAT 3.9-m

NGC 3576 08/02/95 3509-3908 1 -50 11 12 00.5 -61 18 24 300, 1800

NGC 3576 ” 3908-4305 1 -50 ”” 2x1800

NGC 3576 ” 3635-7360 8.5 -50 ” 120, 300, 600

NTT 3.5-m

30 Doradus 15/12/95 3635-4145 2 76 05 38 45.6 -69 05 24 3x1200

30 Doradus ” 4060-4520 2 76 ” 3 x1200

30 Doradus ” 4515-4975 2 76 ” ” 300, 4x1200

30 Doradus ” 6507-7828 3 .8 76 ” 3x1200

30 Doradus ” 3800-8400 11 76 ” 60, 300, 600

LMC N llB 16/12/95 3635-4145 2 -57 04 55 58.8 -66 25 12 600

LMC N llB ” 4060-4520 2 -57 ” 2x1800

LMC N llB ” 4515-4975 2 -57 ” 2x1800

LMC N llB " 6507-7828 3 .8 -57 ” ” 600

LMC N llB ” 3800-8400 11 -57 ” ” 600

SMC N66 16/12/95 3635-4145 2 -57 00 58 55.2 -72 12 32 600

SMC N66 ” 4060-4520 2 -57 ” 2x1800

SMC N66 4515-4975 2 -57 ” 2x1800

SMC N66 ” 6507-7828 3 .8 -57 600

SMC N66 ” 3800-8400 11 -57 ” 300, 600 4-1- Observations 155

15/im CCD in February 1997. A 2 arcsec wide, 3.5 arcmin long slit was employed. The CCDs were binned by a factor of two along the slit direction, in order to reduce the read-out noise. The spatial sampling was 1.63 arcsec per pixel projected on the sky. Two wavelength regions of M 17 were observed in July 1996: a 2400 lines mm“^ holographic grating was used in first order to cover the 3995-4978 Â range at a spectral resolution of 1.5 Â (FWHM); a second grating in first order, along with a WC345 order sorting filter, was used to cover the 3535-7400 Â range, at a resolution of 4.5 Â. The shortest integration time was chosen so as to ensure that strong emission lines like Ho; and the [O III] AA4959, 5007 nebular lines would not be saturated. NGC 3576 was observed in February 1997, in the 3995-4978 Â range only, at a resolution of 1.5 Â. Additional spectra of NGC 3576 were taken at the AAT with the RGO spectrograph and a TEK 1000 x 1000, 24/am x 24/im CCD. A 1200 lines mm~^ grating was used in two settings to cover the 3509-3908, 3908-4305 Â ranges, at a resolution of 1 Â, while another grating with 250 lines mm“ ^ was used to cover the 3655-7360 Â range at a resolution of 8.5 Â. The CCD was again binned by a factor of two along the slit direction, yielding a plate scale of 1.54 arcsec per pixel.

The Magellanic Cloud H II regions 30 Doradus, NllB (in the LMC) and N66 (in the SMC) were observed with the NTT 3.5-m telescope in December 1995. The ESQ Multi Mode Instrument (EMMI) was used in the following modes: red imaging and low dispersion grism spectroscopy (RILD), blue medium dispersion spectroscopy (BLMD) and dichroic medium dispersion spectroscopy (DIMD). The detector was a TEK 1024 x 1024, 24pm x 24pm CCD (no. 31), used while in BLMD observing mode and a TEK 2048 x 2048, 24pm x 24pm CCD (no. 36), while in RILD mode. Both cameras are in use when observing in DIMD mode. In this case, a dichroic prism is inserted into the beam path so that light is di­ rected to the blue and red grating units in synchronization, allowing simultaneous exposures to be obtained in the blue and red part of the optical spectrum. In all exposures, both CCDs were binned by a factor of two in both directions. The 4^.2. Data reduction 156

spatial sampling was thus 0.74 and 0.54 arcsec per pixel projected on the sky, for CCDs no. 31 and no. 36, respectively. Five wavelength regions were observed with two different gratings (#3, #7) and a grism unit (#3) at spectral resolutions of approximately 2 A (AA3635-4145, AA4060-4520, AA4515-4975), 3.8 A (AA6507- 7828), and 11A FWHM (AA3800-8400), respectively. An OG530 filter was used when observing in DIMD mode. The slits used were 5.6 arcmin long and 1, 1.5 arcsec wide. The relevant exposure times, position angles and target coordinates are listed in Table 4.1.

4.2 Data reduction

The two-dimensional spectra were reduced with the MIDAS software package, fol­ lowing standard procedures. They were bias-subtracted, fiat-fielded via division by normalized flat held frames, cosmic-rays cleaned, and then wavelength cali­ brated using exposures of He-Ar, Th-Ar and Cu-Ar calibration lamps. During the 1995 and 1997 runs, twilight sky flat-helds were also obtained, in order to correct the small variations in illumination along the slit. The ESO 1.52-m spectra were reduced to absolute intensity units using wide-slit (8 arcsec) observations of the HST standard stars FeigellO, (the nucleus of the planetary nebula) NGC 7293 (Walsh 1993), and the CTIO standards LTT 4364 and LTT6248 (Hamuy et al. 1994). All NTT spectra were hux-calibrated using wide-slit (5 arcsec) observa­ tions of FeigellO. The AAT spectra were hux-calibrated using observations of the standard star LTT 3218. In all cases, hux-calibration was done using the IRAF software package. A complete list of observed emission lines and their huxes for all nebulae can be found in the Appendix. All huxes are in a scale where F(H/3) = 100, with the dereddened hux given by,

7(A) = ioc(7f/3)y(A)

The amount of interstellar extinction is given by c(H/?) which is the logarithmic 4.3. Nebular analysis 157 difference between observed and dereddened H/? fluxes, while / (A) is the adopted extinction curve in each case (see below) normalized such that /(H/?) = 0. A ratio of total extinction Ay, to colour-excess E{B — V), oî R = 3.1 was assumed. The line fluxes were retrieved using line profile fitting techniques, in a similar manner to that described in Chapter 2, employing multiple Gaussian fitting for blended features.

4.3 Nebular analysis

4.3.1 Reddening correction

In the case of the galactic nebulae M 17 and NGC 3576, I used the galactic red­ dening law of Howarth (1983) to derive the interstellar extinction, by comparing the observed Balmer Ho;/H/3, H7/H/5 and H^/H/3 decrements to their Case B theoretical values from Storey & Hummer (1995). For the M 17 and NGC3576 sightlines, c(H/?) = 1.84 and 1.25 was found respectively. For the Cloud nebulae, the contribution from galactic foreground reddening was derived from the reddening maps of Burstein & Heiles (1982), using the ex­ tinction law of Howarth (1983) in all cases. For 30 Doradus c(gal) = 0.087 was found. The remaining extinction due to the 30 Doradus immediate environment was found to be c(30Dor) = 0.32, from the resulting Balmer decrements using the LMC extinction law of Howarth (1983). For LMC N llB and SMC N66, no further correction was attempted after the foreground reddening was removed, since the resulting Balmer line ratios were in good agreement with their Case B theoret­ ical values. I derived c(gal) = 0.073 and 0.087 for LMC NllB and SMC N66, respectively.

4.3.2 Electron temperatures and densities

The nebular electron temperatures and densities were derived from several CEL diagnostic ratios by solving the equations of statistical equilibrium using the multi­ 4-3. Nebular analysis 158

level (> 5) atomic model EQUIB and are presented in Table 4.2. The atomic data sets used for this purpose, as well as for the derivation of abundances are the same as those used in our extensive PN study (Chapter 2, Table 2.3). The procedure to derive temperatures and densities is as follows: a representative initial Tg of

9000 K was assumed in order to derive N e{C l III) and Ne(Ar iv); the mean density from these diagnostics was then used to derive Te([0 III]) and I iterated once to get the final values. In a similar manner, Tg([N llj) was derived in conjunction with Ne{0 ll) and Nq{S ll). I derived Balmer jump temperatures from the ratio of the continuum Balmer discontinuity at 3646 Â to H 11 A3770 for M 17 and NGC 3576. These are presented in Table 4.2, along with the implied mean nebular temperature. To, and temperature fluctuation parameter, (Peimbert 1967; cf. Chapter 1, Eqs. 1.16 and 1.17). As is the case for the PN diagnostics presented earlier in this study, the electron temperatures deduced from the [N II] nebular to auroral line ratio are higher than those of the corresponding [O III] line ratio—except from LMC NllB, where the values derived are probably consistent within the errors; the differences exceed 2000 K for 30 Doradus and SMC N66. Using photoionization models of metal-rich

H II regions, Stasihska (1980) has argued that the electron temperature increases outward as a function of radius, probably due to a combination of hardening of the radiation field with increasing optical depth, plus stronger cooling from fine- structure lines of [O III] in the inner parts of the nebulae where dominates; thus the temperature in zones where singly ionized species exist is predicted to be higher than Te([0 ill]). However, as we saw in Chapter 2 the contribution of recombination to the excitation of the [N ll] A5754 line, coupled with the potential presence of high- density inclusions in nebulae, may lead deceptively high temperatures to be derived from the [N ll] nebular to auroral line ratio, as well as from the corresponding [O ll] line ratio. Following the procedure laid out in Chapter 2, I estimated corrections to the [N ll] temperatures of M 17, NGC 3576 and 30 Doradus, making Jf.,3. Nebular analysis 159

Table 4.2:: Plasm a Diagnostics.

Diagnostic ratio M17 NGC 3576 30 Doradus LMC NllB SMC N66

%RK)

[0 III] (A4959+A5007)/A4363 8200 8850 10100 9400 12400 [N ii] (A6548+A6584)/A5754 9100 9000 12275 9250 14825 7900“ 8500“ 11600“ ** [S n] A4068/(A6717+A6730) 8850 7900 7500 6950 11050 B J /H ll 7700 8070 *** To 7840 8300 *** e 0.011 0.017 * **

Ne (cm “ ^)

[A r IV] A4740/A4711 1500 1700 1800 * < 100 [Cl III] A5537/A5517 1050 2700 480 1700 3700 [S II] A6731/A6716 600 1350 390 80 60 [0 ii] A3739/A3726 * 1300 370 110 50

“ [N ll] temperatures after correction for recombination excitation contributions.

use of the derived ORL fractions presented in Table 4.5. The revised [N ll] temperatures for these three nebulae respectively, are 7900 K, 8500 K and

11 600 K; improved agreement with [O III] temperatures is found in all cases. Given the inherent uncertainties I have neglected Te([N ll])’s in the abundance analysis that follows and adopted the [0 III] temperatures instead. With respect to the derived electron densities, it is found that in four out of five nebulae the densities deduced from the [S ll] A6731/A6716 ratio are in good agreement with those derived from [O ll] A3729/A3727 (the spectral resolution was not adequate to allow the determination of the [0 ll] doublet ratio in M 17), but are lower than the ones given by the [Ar iv] and [Cl III] diagnostics. This can be seen especially in LMC NllB and SMC N66 where the [O ll] and [S ll] doublet ratios approach the low-density limit, but certainly imply an order of magnitude less than from the doubly ionized species. This behaviour is consistent with the presence of strong density variations in the nebulae, so that the diagnostic line ratios with higher critical densities yield higher derived nebular electron densities (see Rubin 1989, Liu et al. 2001 and Chapter 2). 4- 4- Ionic and total elemental abundances from CELs 160

4.4 Ionic and total elemental abundances from CELs

The ionic abundances deduced from optical CELs are presented in Table 4.4. The electron temperature derived from the [O III] nebular to auroral line ratio has been adopted for all ionic species and for all five nebulae, according to the discussion detailed in the previous section. Regarding the choice of electron densities, we have adopted the mean values deduced from the [O ll] and [S ll] doublet ratios in order to derive abundances of singly ionized species, while the means of the [Ar iv ] and [Cl III] values were used to derive abundances of doubly and triply ionized species. Total abundances from CELs have been derived adopting the ICF scheme of KB94, apart from Cl which is not discussed by those authors. For that element and in the cases of M 17, 30 Doradus and SMC N66, the prescription of Liu et al. (2001) was used according to which Cl/H = (S/S^+) x CF'^/H'*’, based on the similarities of the ionization potentials of Cl ionic stages to those of the S ionic stages. On the other hand, for NGC 3576 and LMC NllB, whose listed S abundances are lower limits, we adopted Cl/H = (Ar/Ar^"*") x CF"^/H"^. The total oxygen abundance is the sum of singly and doubly ionized oxygen. No significant amounts of are expected, since no He II lines are seen.

4.5 Ionic abundances from ORLs

Since ionic abundances derived from optical recombination lines, relative to H+, are almost completely insensitive to temperature and density variations, the stan­ dard [O III] temperature and mean [Cl III], [Ar iv ] densities for each nebula, were adopted for the calculations.

4.5.1 C2+/H+ and N^+/H+

The case-insensitive 3d-4f C II A4267 line has been detected and used to derive the C^’^'/H'^ abundance ratios presented in Table 4.4; an upper limit only has Jf.,5. Ionic abundances from ORLs 161

Table 4.3: Ionic and total Helium and O, N, C, Ne, Ar, S and Cl abundances.“

M17 NGC 3576 30 Doradus LMC NllB SMC N66

4471 He+/H+ 0.0955(1) 0.0949(1) 0.0874(1) 0.0856(1) 0.0791(1)

5876 He+/H+ 0.0959(4) 0.0926(4) 0.0910(3) 0.1022(4) 0.1006(4)

6678 He+/H+ . 0.0909(1) 0.0879(1) 0.0926(1) 0.0906(1) 0.0756(1)

Avg. He+/H+ 0.0950 0.0922 0.0906 0.0975 0.0928

4686 He^+/H+ 0.0000 0.0000 0.0000 0.0000 0.0000

H e/H 0.0950 0.0922 0.0906 0.0975 0.0928

3727 0+/H+ 9.28e-5 l.lOe-4 4.13e-5 1.12e-4 3.35e-5

7330 1.18e-4: 2.06e-4: 5.77e-5: 3.79e-4: 8.77e-6:

4931 0^+/H+ 2.23e-4: 3.31e-4: * 2.36e-4: 6.30e-5:

4959 2.66e-4 2.21e-4 1.76e-4 1.47e-4 9.41e-5

ic /( 0 ) 1.00 1.00 1.00 1.00 1.00

0 /H 3.59e-4 3.310-4 2.17e-4 2.59e-4 1.28e-4

6584 N + /H + 8.14e-6 1.17e-5 1.48e-6 3.66e-6 5.37e-7

3.87 3.01 5.25 2.31 &82

N/H 3.15e-5 3.52e-5 7.77e-6 8.45e-6 2.05e-6

3868 N e^+/H + 7.17e-5 3.46e-5 3.71e-5 2.40e-5 1.72e-5

zc/(Ne) 1.39 1.50 1.23 1.76 1.36

Ne/H 9.97e-5 5.19e-5 4.56e-5 4.22e-5 2.34e-5

7135 Ar^+/H+ 1.48e-6 1.70e-6 l.lle-6 1.12e-6 4.71e-7

4740 Ar^+/H+ 2.72e-8 2.15e-8 3.39e-8 2.27e-8 3.16e-8

ïc/(Ar) 1.35 1.50 1.24 1.76 1.35

Ar/H 2.03e-6 2.58e-6 1.42e-6 2.01e-6 6.79e-7

4069 S + /H + 5.79e-7: 7.50e-7: 2.53e-7: 8.82e-7: 2.60e-7:

6725 3.60e-7 6.20e-7 2.80e-7 7.32e-7 2.59e-7

6312 S^+/H+ 8.05e-6 * 4.28e-6 * 1.66e-6

tc/(S ) 1.19 * 1.29 * 1.19

S/H l.OOe-5 6.57e-7 5j& ^ 6 4.98e-6 2.28e-6

5517 Cl^+/H+ 1.07e-7 9.95e-8 5.14e-8 7.69e-8 2.63e-8

ic/(C l) 1.24 1.52 1.37 1.79 1.37

Cl/H 1.33e-7 1.51e-7 7.04e-8 1.38e-7 3.60e-8

“ Numbers followed by have not been used at any point in the analysis. 4.5. Ionic abundances from ORLs 162

Table 4.4; Ionic carbon abundances from optical recombination lines.

M 17 NGC 3576 30 Doradus LMC NllB SMC N66

/(A4267) 0.482 0.312 0.0919 0.196 < 0.0419 10^ X C 2+/H + 4.35 2.87 0.882 1.82 < 0 .4 3 3

been estimated for SMC N66. The doublet from multiplet V 4 (A3920) has also been detected from NGC 3576, but not used for abundance determination, since it is directly connected to the C"*" 2p^F° ground term and is therefore potentially affected by optical depth effects. Regarding the ORL abundance ratios, these have been derived for M 17 and NGC 3576 and are presented in Table 4.5; lines from the 3s-3p multi­ plets V 3 and V 5 have been detected and Case B recombination has been assumed. Under Case A the abundance ratios deduced from V 5 lines would be several times higher, while V 3 results would be higher by only about 20 per cent. Longer ex­ posure times would be required to detect weaker lines from the 3d-4f group. N II

ORLs have not been unambiguously detected from any of the Cloud H II regions; the strongest predicted multiplet V 3 (A5630) is probably present in 30 Doradus, but the resolution of our spectra is too low (FWHM 11 Â) at this wavelength to allow us to be conclusive. In that nebula’s spectrum the strongest N II A4630 multi­ plet V 5 line is marginally detected; an upper limit N^'^/H+ fraction of 5.48 x 10~^ was derived from it. From Table 4.5 we adopt fractions of 3.53 x 10"^ and 2.70 x 10^^ for M 17 and NGC 3576, respectively.

4.5.2 Q2+/H+

This is the first time that O II ORLs arising from Magellanic Cloud H II regions have been recorded. We have detected transitions from 3s-3p, as well as from 3p-3d and 3d-4f configurations. In Table 4.6 we present a comparison between the observed and predicted intensities of O II fine structure lines relative to the 4-5. Ionic abundances from ORLs 163

Table 4.5: Recombination line abundances. N2+ N2+ N2 + Ao Mult. Jobs lobs lobs H+ H+ H+ (Â) (10-4) (10-4) (10-4)

MIT NGC 3576 30 Doradus V3 3s^P°-3p^D 5666.63 V3 .05699 4.35 .03581 2.80 * * 5676.02 V3 .01702 2.93 .01589 2.72 ** 5679.56 V3 .05444 2.23 .06669 2.75 * * V5 3s^P°-3p^P 4630.54 V5 .06016 4.99 .03086 2.56 .00658 .548 V20 3p^D-3d^D° 4803.29 V20 .03039 4.52 * * * * S um .219 3.53 .149 2.70 * .548

strongest expected line within each multiplet. As with our extensive PN recombi­ nation line survey, it is assumed that L^-coupling holds for the 3s-3p transitions, while intermediate coupling is assumed for those between 3p-3d and 3d-4f states. Table 4.6 is important both for checking whether observation agrees with theory and in excluding some lines from further consideration when blending or misiden- tihcation are suspected.

Table 4.6: Comparison of the observed and predicted relative intensities of H II region O II lines.

A o ( Â ) Mult Termi-Term^ gl-gu Ipred lobs lobs/lpred

M 17

4638.86 VI 3s 4P-3p 4D* 2-4 0.21 0.77[.12] 3.7[0.6]

4641.81 VI 3s 4P-3p 4D* 4-6 0.53 1.11[.16] 2.1[0.3]

4649.13 VI 3s 4P-3p 4D* 6-8 1.00 1.00[.10] 1.0[0.1]

4650.84 VI 3s 4P-3p 4D* 2-2 0.21 0.63[.ll] 3.0[0.5]

4661.63 VI 3s 4P-3p 4D* 4-4 0.27 0.76[.ll] 2.8[0.4]

4676.24 VI 3s 4P-3p 4D* 6-6 0.22 0.46[.10] 2.1[0.5]

4317.14 V2 3s 4P-3p 4P* 2-4 0.44 1.04[.32] 2.6[0.8]

4319.63 V2 3s 4P-3p 4P* 4-6 0.43 0.31[.23] 0.7[0.5]

4345.56 V2 3s 4P-3p 4P* 4-2 0.40 0.63[.20] 1.5[0.5] 4.5. Ionic abundances from ORLs 164

Ao(Â) Mult Term/-Terniu gl-g u ïpred Io6s loba/Ipred

4349.43 V2 3s 4P-3p 4P* 6-6 1.00 1.00[.10] 1.0[0.1]

NGC 3576

4638.86 VI 3s 4P-3p 4D* 2-4 0.21 0.55[.ll] 2.6[0.5]

4641.81 VI 3s 4P-3p 4D* 4-6 0.53 0.95[.16] 1.8[0.3] 4649.13 VI 3s 4P-3p 4D* 6-8 1.00 1.00[.10] 1.0[0.1]

4650.84 VI 3s 4P-3p 4D* 2-2 0.21 0.60[.ll] 2.9[0.5]

4661.63 VI 3s 4P-3p 4D* 4-4 0.27 0.99[.16] 3.7[0.6]

4676.24 VI 3s 4P-3p 4D* 6-6 0.22 0.60[.16] 2.7[0.7]

(3 p —3d)

4069.89 VIO 3p 4D*-3d 4P 4-6 0.74 1.40[.20] 1.9[0.3]

4072.16 VIO 3p 4D*-3d 4P 6-8 0.69 1.20[.82] 1.7[1.2]

4075.86 VIO 3p 4D*-3d 4P 8-10 1.00 1.00[.10] 1.0[0.1]

4085.11 VIO 3p 4D*-3d 4P 6-6 0.13 0.17[.10] 1.3[0.8]

4087.15 T/48C 3d 4P-4f 03* 4-6 0.27 1.30[.48] 4.8[1.8] 4089.29 V48a 3d 4P-4f 05* 10-12 1.00 1.00[.10] 1.0[0.1] 4275.55 V67a 3d 4D-4f P4* 8-10 1.20 4.18[1.1] 3.5[0.9] 4357.25 V63a 3d 4D-4f D3* 6-8 0.05 2.36[.85] 47.[17.] 4609.44 T/92a 3d 2D-4f P4* 6-8 0.14 1.69[.64] 12.[5.0]

30 Doradus

4638.86 VI 3s 4P-3p 4D* 2-4 0.21 1.00[.15] 4.8[0.7]

4641.81 VI 3s 4P-3p 4D* 4-6 0.53 1.33[.18] 2.5[0.3]

4649.13 VI 3s 4P-3p 4D* 6-8 1.00 1.00[.10] 1.0[0.1]

4650.84 VI 3s 4P-3p 4D* 2-2 0.21 1.00[.19] 4.8[0.9]

4661.63 VI 3s 4P-3p 4D* 4-4 0.27 0.88[.12] 3.3[0.4]

4673.73 VI 3s 4P-3p 4D* 4-2 0.04 0.19[.03] 4.8[0.8]

4676.24 VI 3s 4P-3p 4D* 6-6 0.22 0.35[.05] 1.6[0.2]

4317.14 V2 3s 4P-3p 4P* 2-4 1.00 1.00[.10] 1.0[0.1] 4319.63 V2 3s 4P-3p 4P* 4-6 0.98 1.06[.12] 1.1[0.1]

4069.89 VIO 3p 4D *-3d 4P 4-6 1.00 1.00[.10| 1.0[0.1] 4-5. Ionic abundances from ORLs 165

A o ( A ) Mult Termj-Termu gl-gu Ipred lob s lo b s /I p r e d

4072.16 VIO 3 p 4D *-3d 4F 6-8 0.93 0.51[.09] 0.6[0.1]

4 0 7 8 . 8 4 VIO 3 p 4D *-3d 4F 4 -4 0.14 0.24[.05] 1.7[0.4]

4085.11 VIO 3 p 4D *-3d 4F 6-6 0.17 0.19[.04] 1.1[0.2]

4083.90 V48b 3 d 4 F -4 f G 4 * 6-8 0.29 0.51[.31] 1.8(1.1]

4087.15 V48c 3d 4F-4f G3* 4-6 0.27 0.88[.23] 3.3[0.9]

4 0 8 9 . 2 9 V48a 3d 4F-4f G5* 10-12 1.00 1.00[.13] 1.0[0.1]

4275.55 V67a 3d 4D-41 F4* 8-10 1.33 3.73[.71] 2.8[0.5]

4 2 8 8 . 8 2 V53c 3d 4P-4f Dl* 2-4 0.83 11.9[2.6] 14.[3.1]

4303.83 V53a 3d 4P-4f D 3 * 6-8 0.50 2.47[.44] 4.9[0.9]

4313.44 V78a 3d 2F-4f F4* 8-10 0.61 2.06[.68] 3.4(1.1]

4315.69 T /6 3 C 3d 4D-4f Dl* 6-4 0.11 1.82[.67] 17.(6.1]

4609.44 V92a 3d 2D-4f F4* 6-8 0.57 1.00[.24] 2.0(0.4]

LMG N llB

4 6 3 8 . 8 6 VI 3s 4P-3p 4D* 2-4 0.21 1.54[.15] 7.3 (0.7]

4641.81 VI 3s 4P-3p 4D* 4-6 0.53 0.70[.18] 1.3 (0.3]

4649.13 VI 3s 4P-3p 4D* 6-8 1.00 1.00[.10] 1.0 (0.1]

4650.84 VI 3s 4P-3p 4D* 2-2 0.21 0.83[.19] 4.0 (0.9]

4661.63 VI 3s 4P-3p 4D* 4-4 0.27 0.84[.12] 3.1 (0.4]

4673.73 VI 3s 4P-3p 4D* 4-2 0.04 0.41 [.03] 10. (1.0]

4676.24 VI 3s 4P-3p 4D* 6-6 0.22 0.36[.05] 1.6 (0.2]

SMC N66

4 6 3 8 . 8 6 VI 3s 4P-3p 4D* 2-4 0.21 0.57[.15] 2.7(0.7]

4649.13 VI 3s 4P-3p 4D* 6-8 1.00 1.00[.10] 1.0(0.1]

4650.84 VI 3s 4P-3p 4D* 2-2 0.21 0.94[.19] 4.5(0.9]

4661.63 VI 3s 4P-3p 4D* 4-4 0.27 1.47[.32] 5.4(1.2]

4.5.2.1 Relative intensities of fine-structure V 1 lines

In contradiction to our results from a comparison of O II relative line intensities from PN (Chapter 2, Table 2.12), there is clear evidence that several transitions are stronger than expected in all five H II regions studied. The effect is especially 4-5. Ionic abundances from ORLs 166 pronounced amongst the fine structure lines of multiplet V 1; those of multiplet V 10, however, are in better agreement with theory even though their formal mea­ surement errors are somewhat larger due to partial blending with the [S ll] AA4068, 4076 doublet. A similar situation has been reported by Esteban et al. (1999) from their echelle observations of the Lagoon Nebula (M8). Esteban et al. (1998) how­ ever, found quite good agreement between O II multiplet V 1 lines from the Orion Nebula, a result supported by our examination of the relative intensities of the same transitions using the recent Orion line atlas by Baldwin et al. (2000), which is of better spectral resolution than the data used by the former authors. Our observations show that in all 5 H ll regions the A4638.86 (J = 1/2-3/2), A4650.84 ( J = 1/2-1/2) transitions are enhanced by the same amount within each nebula: a factor ranging from ~ 2.8-4.8 across the five objects under study; in Orion this fac­ tor is only ~ 1.4 for both lines, but less for the other V 1 transitions. The A4661.63 (J = 3/2-3/2) line follows closely in displaying abnormal strength compared to theory. One possible explanation for this behaviour involves the breakdown of thermal equilibrium among the fine structure levels of the ^Po, 1,2 ground term of O^"*". Our analysis of 3s-3p transitions employs term-averaged effective recombination coefficients from Storey (1994). The corresponding effective recombination coef­ ficient for a particular fine structure line, J —>■ J', is proportional to the statistical weight {çl = 2Jl T 1) of the level under LTE conditions; it thus follows that the line strength is also proportional to gi^. In cases however where LTE does not hold and the level populations are far from equilibrium, individual line strengths are perturbed and lines will display abnormal relative intensities. This could come about if the electron density of the nebula is lower than the critical densities of the ^Pq, 1,2 levels [~ (2-3) x 10^ cm“^]. The mean derived from various diagnostics in the objects under study is in all cases lower than 2000 cm“^, while in the three Cloud Hll regions it is lower than 1000 cm"^ (Table 4.2); contrary to Orion where

Esteban et al. find Nq ~ 5000 cm~^). Thus, it seems that within multiplet V 1 4-5. Ionic abundances from ORLs 167

the strongest expected line A4649.13 (J = 5/2-T/2), is weakened at the expense

of transitions between levels of lower Jl . This argument seems to be supported

by our previous analysis of O II ORLs originating from planetary nebulae. For example, the observed relative intensities of V 1 lines from the dense PN IC4191 and NGC 5315 are in perfect agreement with theory (cf. Chapter 2, Table 2.12); the derived mean of these objects is 10 700 and 14100 cm"^ respectively. The opposite is true in the case of PN NGC 3132 (Ng — 600 cm~^), where the V 1 lines

display abnormal ratios just like they do for our H II regions. If the above inter­ pretation is correct however, the total observed intensity of the multiplet should not be affected and can be used to derive a reliable ORL abundance. A similar effect may be affecting the relative intensities of lines from the 3d-4f group, weakening the A4089 transition. A major point to be deduced from this analysis is that the observed relative

intensities of O II ORLs point towards an origin for these lines in very low density

gas, similar to that emitting the [0 III] CELs.

4.5.2.2 Continuum observations and scattered light

An issue that had to be addressed in the course of this analysis is the potential influence of dust-scattered stellar light on the intensities of weak nebular emission

features like the O II ORLs of interest. Unlike PN where the size of the emitting

region results in relatively small dust columns, the situation is different in H II regions whose big volumes contain much larger dust columns. The dust effectively scatters light from the nebula’s illuminating stars, which then makes up a major fraction of the observed continuum at optical and UV wavelengths. It is thus possible that emission or absorption features in the spectra of exciting stars in

an H II region may contaminate the nebular emission spectrum. Peimbert et al.

(1993) dealt with this problem in an analysis of the Orion and M l7 H II regions. They employed previously published medium resolution (5-7 Â) spectroscopic data in order to derive ORL abundances; their resolution however was not 4-5. Ionic abundances from ORLs 168

adequate to allow for the effect of blending of O II ORLs with other lines e.g.

[Fe ll], N I I I and C I I I , so they resorted to older photographic line intensities in order to estimate the necessary corrections. Our CCD spectrophotometry is of much better resolution (FWHM 1, 1.5, 2Â for the coverage of O II ORLs) and our line intensities were consistently corrected for blends with other nebular emission lines. In order to investigate the potential effect of dust-scattered stellar light on the nebular line fluxes, we have measured the continuum emission at A4089 and

A4650, which coincide with the strongest 3d-4f O II A4089.3 transition and O II multiplet V 1 ORLs respectively. In Table 4.7 the continuum intensity, /^(A), and the scattered light contribution, /^(A)d, to the observed continuum are presented. The observed continuum is mainly due to two components: the atomic continua and the dust-scattered light, so that,

r(A) = r(A)a-kr(A)d, where /^(A)a represents the sum of the Hi and He I atomic continua calculated using emissivities by Storey & Hummer (1995) and Brown & Mathews (1970), respectively; the derived Tg, iVg and He"^/H"*" for each object were also taken into account. We find that for M 17 dust-scattered light accounts for 78 and 69 per cent of the observed continuum at A4089 and A4650, respectively; for NGC 3576 these values are 87 and 82 per cent; for 30 Doradus the corresponding values are 74 and 67per cent; for LMC N llB they are 85 and 83 per cent; finally, for SMC N66 they are 86 and 85 per cent, respectively (sky-subtraction has not been done; contribution assumed negligible for these dark-of-Moon conditions^).

In Table 4.8 we present estimated equivalent widths, £'Wabs(^e6), for stellar lines in absorption (in mÂ) in the observed continuum given by,

EWabs(Tie6) = EWabs(sW/or)r(A)d/r(A), (4.1)

^Typical dark-of-Moon sky brightness at B (4400 A) is 22”^.8 per arcsec^ corresponding to

4.73 X 10~^® ergs“^ cm“^ arcsec"^ when B = 0 is 6.24 x 10“® ergs“^ cm“^ A“^; thus for e.g.

30 Doradus the sky contribution to the observed continuum at A4650 is ~ 0.082 per cent only. 4-5. Ionic abundances from ORLs 169

Table 4.7: Continuum emission and scattered light, expressed as log/‘^(A)//(H/3) when / ‘^(A)//(H/3) is in units of

M17 NGC3576 30 Doradus LMC NllB SMC N66 r(A) r(A)d r(A) r(A)a r(A) r(A)a r(A) r(A)a r(A) r(A)a

A4089 -2.402 -2.508 -2.218 -2.279 -2.433 -2.564 -2.063 -2.135 -2.103 -2.168

A4650 -2.645 -2.806 -2.435 -2.522 -2.624 -2.797 -2.278 -2.357 -2.214 -2.283

with the continuum intensity values being those of Table 4.7. This analysis pertains to the stellar content of 30 Doradus in the following way: according to Walborn & Blades 1997 (WB97), the visually brightest stars in the 30 Doradus association are B-type supergiants, along with the Of-type star R 139 (= Parker 952). Therefore, we will assume that their spectra dominate the ob­ served dust-scattered light. We have singled out seven stars which are the bright­ est in the central ~1 arcmin^ of the cluster (excluding the compact core R 136): R137 = P548 (B0.7-1.5I, V = 12.14); R138 (AOla, V = 11.87); R139 = P952 (WNL 4- Of, y = 11.94); R140 = P877, 880 (WN -f WC, V = 12.22, 12.79); R141 = P1253 (BN0.5Ia, V = 12.57); R142 = P987 (B0.5-0.7I, R = 11.91); and P 767 (0 3 If*, V = 12.87). We have high resolution spectrograms (at 0.9 À/pix) of R 139 (Fig. 4.1) and R 140 (Fig. 4.2), since they fell on our slit; the remaining five stars are positioned on either side of it. It is assumed that the spectra of R137 and R142 are similar to that of Parker 3157 (BCl la, V = 12.47; Parker et al. 1992) a supergiant belonging to the LMC LH 10 (N il) association, whose spectrogram (at 0.9 A/pix) was extracted from our LMC N llB frames. High resolution spectra of R 141 and P 767 can be found in WB97; we do not have a spectrum of R 138, but we can safely assume that it is featureless at the wavelengths of interest judging from its spectral type.

The O II spectrum reaches a sharp maximum in absorption in stars of early- B spectral type (e.g. Walborn & Fitzpatrick 1990), thus the very existence of 4.-5. Ionic abundances from ORLs 170

Table 4.8: Estimated equivalent widths (in units of mÂ).

30 Doradus LMC N llB

F W e x n E W a h s { n e b ) E W e m E W a h s { n e b )

A4072 185 107 * *

A 4 0 8 9 57 87 151 98

A4640 637 37 2 8 6 167

A4650 540 185 264 2 9 8

A4661 238 40 121 78

B-type supergiants in nebulae complicates the analysis of nebular O II ORLs. It is therefore recommended that fields of view relatively clear of such stars are chosen for future weak-emission line analyses of Hll regions. The stellar absorption and/or emission features that are of relevance to this analysis are: Si IV -f O II

A4089, N iii-b O II AA4638, 4640, C lil + O ii A4650 and O ii A4661. Values of EWa\^s{ste.llar) were estimated at several wavelengths of interest on our high-resolution stellar spectrograms mentioned above and subsequently averaged, weighted according to the brightness ratios of the seven illuminating stars in the blue wavelength region, i.e., 5.6 : 6.3 : 4.3 : 6.8 : 3.2 : 9.8 : 1.0 (for R 137 : .. : R142 ; P 767). We have assumed a total-to-selective extinction ratio of R b = 5 (Hill et al. 1993) and adopted apparent magnitudes and (B-V), E{B-V) values from Parker (1993). We were then able to compute £'Wabs(^e6) values using Eq. 4.1. Finally, in order to correct the observations we added these to the observed EWem values. Moving on to LMC NllB, we find the early-B supergiant Parker 3157 (BCl la,

V = 12.47; Parker et al. 1992) with its extremely rich O II absorption line spec­ trum directly in the field of view of the spectrograph’s slit (Fig. 4.3). The scattered light from this star is responsible for a significant drop in the intensity of the ob­ served continuum, in the region of AA4638, 4640, AA4649, 4650 and A4661 lines of multiplet VI. In contrast to 30 Doradus however, the dust-scattered light in ^.5. Ionic abundances from ORLs 171

N llB does not contain any obvious contribution from Of- and WR-type luminous stars equivalent to e.g. R 139 or R 140 (no Wolf-Rayet stars are listed in the census of N llB by Parker et al.), whose emission in the A4650 region partly compensates absorption effects due to early-B supergiants and results in a fairly level continuum at those wavelengths for the former nebula. Therefore, along with the BCl la-type supergiant P 3157, we will assume that the observed continuum in N llB in dom­ inated by the scattered light of six other stars, which are visually the brightest amongst those studied by Parker et al.; they are the following: P 3209 (03III(f*), 12.66); P3252 (B2II, V = 11.35); P3271 (Bill, V = 12.99); P3070 (06V , V = 12.75); P3223 (08.5IV, V = 12.95); and P 3120 (0 5 .5 V((f*)), V = 12.80). High resolution rectified spectrograms for these stars can be found in Parker et al. and were used to estimate mean EWahsi^tellar) values. More accurate mea­ surements were possible for P 3157 for which we have a digital spectrogram. A cross-examination between that digital spectrogram and the ones in Parker et ah, which include P 3157, shows agreement, pointing towards a reliable estimate of the EWs. In Table 4.9 we present the measured intensities and the derived abundances from O II ORLs for all five Hll regions, before any correction for scattered light contamination. Along with values from individual transitions, abundances de­ rived from total multiplet intensities are also listed and are discussed in more detail below for each object accordingly. In the context of that discussion revised abundance ratios are presented, incorporating the effects of the scattered stellar continua on the emission line intensities wherever possible. 4-5. Ionic abundances from ORLs 172

Table 4.9: Recombination line abundances in H II regions.

0 ^+ Ao(Â) Mult lobs H+ (1 0 -') M 17

4638.86 VI .1056 10.26

4641.81 VI .1524 5.87

4649.13 VI .1375 2.78

4650.84 V I .08713 8.46

4661.63 VI .1043 7.93

4676.24 V I .06298 5.70

VI 3s^P-3p^D° 0.650 5.41

4317.14 V2 .1115: 14.82:

4319.63 V2 .03304 4.07

4345.56 V2 .06700 8.55

4349.43 V2 .1070 5.67

V2 3s^P-3p^P° 0.207 5.95

A dopted 5.68

NGC 3576

4638.86 V I .05086 4.94

4641.81 VI .08796 3.39

4649.13 VI .09293 1.88

4650.84 VI .05530 5.37

4661.63 VI .09155 6.96

4676.24 VI .05604 5.07

V 1 3s^P-3p^D° 0.435 3.62

4069.62 VIO .1409: 5.45:

4072.16 VIO .1207 5.01

4075.86 VIO .1010 2.90

4085.11 VIO .1700 3.78

V 10 3p^D°-3d^F 0.239 3.77

4087.15 V48c .04147 13.49

4089.29 V48a .03201 2.82 4275.55 V67a .1339 9.85 j^.5. Ionic abundances from ORLs 173

-^o(Â) Mult lobs (10-4) H+ 3d -4f 0.207 7.40

Adopted 3.70

30 Doradus

4638.86 V I .06404 6.11

4641.81 VI .08537 3.23

4649.13 VI .06416 1.28

4650.84 VI .06413 6.12

4661.63 VI .05652 4.23

4673.73 VI .0 1202 5.80

4676.24 VI .02235 1.99

VI 3s4p_3p4D° 0.369 2.97

4317.14 V2 .07054 6.96

4319.63 V2 .08763 10.91

V 2 3s^P-3p4p° 0.158: 8.42:

4414.90 V5 .03316; 6.27: 4416.97 V5 .04359: 14.8:

V5 3s2p-3p2D° 0.033: 6.27:

4069.89 VIO .1328 5.14

4072.16 VIO .06828 2.84

4078.84 VIO .03243 8.86

4085.11 VIO .02576 5.74

VIO 3p^D°-3d4p 0.259 4.47

4132.80 V19 .05082 9.29

4153.30 V19 .06091 7.80

4156.53 V19 .02318: 18.6:

V 19 3p4p°_3d4p 0.112 8.41

4110.78 V 20 .03218: 13.4:

4906.83 V 28 .03694: 14.8:

4083.90 V48b .01081 3.41 4087.15 T/48c .01852 6.17

4089.29 V48a .02103 1.90

4275.55 V67a .07849: 5.33: 4-5. Ionic abundances from ORLs 174

o^+ Ao(Â) Mult lofas 1 0 H+ ( -") 4288.82 V53c .1091: 11.85:

4303.83 V53a .05188: 9.97:

4313.44 V78a .03833: 28.5:

4315.69 V63c .05334: 43.2:

4609.44 V92a .02094 3.35

3d-4f 0.071 3.03

A dopted 3.49

LMC N llB

4638.86 VI .1170 11.24

4641.81 VI .05303 2.0 2

4649.13 VI .07599 1.52

4650.84 VI .06309 6.06

4661.63 VI .06403 4.81

4673.73 VI .03078 14.93

4676.24 VI .02711 2.43

4696.35 VI .03461 24.07

V 1 3s^P-3p^D° 0.466 3.14

4349.43 V2 .08712 4.60

V2 3s^P-3p'^P° 0.087 4.60

4072.16 VIO .1027 4.27

V 10 3p^D°-3d^F 0.103 4.27

4153.30 V19 .09819 12.54

V19 3p^P°-3d'^P 0.098 12.54

4089.29 V48a .1305 11.41

3d -4 f 0.131 11.41

A dopted 7.19

SMC N 66

4638.49 VI .03110 2.91

4649.13 VI .05449 1.06

4650.84 VI .05126 4.78

4661.63 VI .08000 5.84

V 1 3s^P-3p^D° 0.217 2.15 4.5. Ionic abundances from ORLs 175

Ao(Â) Mult U s (10~^)

A dopted 2.1 5

M 17: the two triplet 3s-3p multiplets V 1 and V 2 are detected, yielding similar results. The mean abundance ratio is 5.68 x 10“'^. Peimbert et al. (1993) have derived a very similar result of 4.85 x 10"^ using multiplets V 1 and V 5, af­ ter correcting the line intensities for underlying absorption due to dust-scattered stellar light. We have not taken into account such an effect for this nebula, since Peimbert et al.’s corrections amount to only 14 per cent for multiplet V 1, the one affected the most. A comparison with the forbidden-line abundance from Table 4.4 yields an ORL/CEL abundance discrepancy ratio of 2.1.

NGC 3576: multiplets V 1 and V 10 yield consistent results; lines of the 3d-4f group display abnormal intensity ratios. It is likely that underlying absorption is affecting the A4089 transition. We do not attempt any corrections and adopt as the ORL the mean from V1 and V10 results, i.e. 3.70 x 10“'^. The ORL/CEL discrepancy ratio in this case is 1.7.

30 D oradus: we have detected lines from the 3p-3d multiplets V 10, V 19, V 20, as well as from the 3s-3p multiplets V 1, V 2, V 5 and several transitions from the 3d- 4f group. Abundance values derived from V 2 and V 19 agree very well with each other, but are about a factor of 2.8 higher than those derived by multiplet V 1 and 88 per cent higher than those of multiplet V 10. The relative intensities amongst the multiplet V 19 components are in good agreement with theory (Table 4.6), apart from A4156.63 (J = 5/2-3/2) which is a factor of 2 stronger than expected relative to A4132.80 (J = 1/2-3/2), probably as a result of blending;^ the V 1 lines

^This is a ratio of lines originating from the same upper level and thus invariably fixed by the ratio of their transition probabihties; the fact that it is observed to differ between several nebulae (e.g. NGC 3242, 3918, 5882, 6153 and 30 Doradus) by a factor of more than 4 suggests that blending with an unknown line affects the A4156 line. J^.5. Ionic abundances from ORLs 176

were discussed in the previous subsection. This abundance discrepancy is quite striking, since from observations of 10 PNe it has been found that on average multiplet V 2 yields abundance values that are 24 per cent lower than those of V I (cf. Chapter 2); furthermore, abundances from multiplets V 10 and V 19 generally agree within 40 per cent. Case dependence does not seem to be an issue since multiplets V 1 and V 10 are almost insensitive to optical depth effects, while multiplets V 2 and V 19, which are assumed here to be under Case B, would yield even higher abundances under Case A. A single line from V 20 is detected, A4110.78 (J = 3/2-1/2), which yields a similar abundance value to those from the V 2 and V 19 lines; the remaining V 20 lines coincide with He I A4120.84 and cannot be reliably used to derive an inde­ pendent abundance estimate. Of the two detected multiplet V 5 lines, A4416.97 ( J — 1/2-3/2) is most probably blended with [Fe II] A4416.27 and is excluded from the analysis, while A4414.90 (J = 3/2-5/2) yields an abundance ratio rather consistent with that from the multiplet V 10 lines. Multiplet V 2 is fed from high-lying terms via A4119 (V 20) and less via V 19 transitions. Multiplet V 19 arises from the Od'^P term which can be reached via permitted transitions from the 0+ 2p^ '^5'° ground state—e.g. the 2p^ ‘^5'°-3d^P A430 line—themselves excited by resonance fluorescence, either by starlight or by another emission line. Along with resonance scattering, the ex­ cited O II 3d level will decay emitting cascade line photons via the multiplets ^5°-4p A4924 (V28), ^po_4p (V19) and 4po_4p ^3907 (Vll), with mul­ tiplet strength ratios of about 26.9 :18.5 :1.0. The strongest line of V 28, A4924.53

( J = 3/2-5/2), coincides in wavelength with [Fe III] A4924.50 and is blended with the much stronger He I A4921.93 line at the resolution of our observations. The second strongest component, A4906.83 (J = 3/2-3/2), is detected in our spec­ tra, yielding an even higher abundance ratio than the V 19 lines, by 70 per cent. There is therefore marginal evidence—based on results from V 28 and V 19—that the population of the 3d term is enhanced by resonance fluorescence from the 4-5. Ionic abundances from ORLs 177

ground state; this might explain the overabundances derived from these lines com­ pared to multiplet V 1 results; V 1 itself is on a cascade route not affected by line fluorescence excitation. However, the single A4110.78 (V 20) line which also pro­ duces an overabundance yields consistent results with V 19, while its upper term 3d^D cannot be reached via permitted transitions from the ground state. Grandi (1976), from a model analysis of the permitted emission line spectrum of Orion, surmised that starlight excitation of the A430 line contributes only 20 per cent as much as recombination to the A4153 (V19) line, while he did not discuss the potentially worse affected AA4906, 4924 (V 28) lines; perhaps in 30 Doradus that contribution is larger. Due to their possible contamination by fluorescence, mul­ tiplet V 19 and V 28 lines will be excluded from our abundance analysis. Regarding the detected 3d-4f transitions, several of them are evidently blended with [Fe ll] and/or Fe II lines and were omitted from further consideration. The

3d-4f O II ORLs are insensitive to optical depth effects and their strengths have been consistently proven to be in excellent agreement with theoretical predictions under an intermediate coupling scheme for a signiflcant number of PN (cf. LSBC, Liu et al. 2000, and this study Chapter 2). From this growing body of work it has emerged that on average abundance ratios derived from the best detected 3s-3p transitions (those of multiplet V 1) are lower than from 3d-4f lines by about 30per cent (c.f. LSBC for a possible explanation of this behaviour). Transitions of the 3d-4f group are close to hydrogenic in nature and the one amongst them with the highest total angular momentum quantum number (A4089, upper state J = 11/2) is not affected by a change of coupling scheme (Storey 1994). Furthermore, their upper terms cannot be reached by permitted resonance lines and are thus unaffected by fluorescence effects. Overall, the 3d-4f lines are the best indicators of 0^'^/H’*' ORL abundances. In 30 Doradus the strongest expected 3d-4f A4089 line is abnormally weak compared to other 3d-4f lines (Table 4.6); as a consequence, the abundance ratio derived from this line alone is lower than that from the V 1 lines, by 36per cent. This is in stark contrast with the standard 4-5. Ionic abundances from ORLs 178 observed behaviour of these two sets of lines as outlined above. An explanation for this result may involve the effect of dust-scattered stellar light on the nebular line intensities. Based on observations of the continuum and scattered light described previously, we estimated corrections to the observed emission line intensities; these amount to 58per cent (A4072); a factor of 2.5 (A4089); 6per cent (AA4638, 4640); 34per cent (AA4649, 4650); and 17per cent (A4661), respectively. The resulting values from the various multiplets display now the more regular pattern already established from the PN analysis; multiplet V 1 results are now 25 per cent lower than those from the co-added 3d- 4f transitions and 20 per cent lower than those from the A4072 V 10 line. The final corrected line intensities and abundance ratios are presented in Table 4.10. The ORL abundance before any corrections is 3.49 x 10“ '^, while that incorporating the corrections becomes 4.67 x 10"^. In 30 Doradus, the revealed ORL/CEL abundance discrepancy ratio before corrections for underlying absorption is 2.0; that after the correction is 2.7.

LMC N llB : in this nebula too the emission line spectrum is contaminated by the scattered continuum of illuminating stars. Regarding the nebular O II ORLs, lines of multiplet V 1 are affected the most. According to our calculations presented in Table 4.8, the following corrections to the observed intensities have been estimated; 65 per cent (A4089); 58 per cent (AA4638, 4640); a factor of 2.1 (AA4649, 4650); and 64per cent (A4661). The corrected line intensities and resulting 0^+/H+ abundance ratios for this nebula are listed in Table 4.10. The ORL abundance before any correction is 7.19 x 10“ ^, while that after is 12.0 x 10"^. Therefore, in this nebula the ORL/CEL discrepancy ratio before the correc­ tions for absorption is 4.9, while that after rises to 8.2.

SMC N66: only multiplet V 1 lines are reliably detected; no corrections for underlying absorption have been estimated. Co-adding the intensities of A4638, 4.5. Ionic abundances from ORLs 179

Table 4.10: Corrected O II intensities and resulting ORL abundances. q 2+ Ao Mult. Icoi Lor H+ H+ (Â) (10-*) ( 1 0 -')

30 Doradus LMC NI IB

4072.16 VIO .1077 4.47 **

4638-40 VI .1601 4.34 .2389 6.51

4649-50 VI .1723 2.84 .2964 4.91

4661 VI .06610 4.94 .1050 7.89

V 1 3s^P-3p^D° 0.399 3.59 0.640 5.26

4083.90 V48b .01081 3.41 * *

4087.15 V48c .01852 6.17 * *

4089.29 V48a .05310 4.79 .2154 18.8

4275.55 V53c .07849 5.34 **

4609.44 V92c .02094 3.35 * *

3d -4f 0.182 4.76 0.215 18.8

A dopted 4.67 12.0

AA4649, 4650 and A4661 transitions, we find an ORL abundance ratio of

2.15 X 10"4. The derived ORL/CEL abundance discrepancy ratio is 2.3.

In Table 4.10 corrected O II line intensities [in units of H/? = 100] are presented, along with revised ORL 0^+/H+ abundances for 30 Doradus and LMC N llB only.

4.5.2.3 Uncertainty

It is likely that our estimated corrections to the observed ORL intensities and the final ORL abundance results of Table 4.10 in the cases of 30 Doradus and LMC NllB, are upper limits only. Clearly, even though the dust-scattered com­ ponent to the observed continuum is found to be very high (~ 80 per cent typi­ cally) and the various absorption/emission EWs were meticulously measured, not all stars that contribute to the continuum were traced. For instance, the field 4-6. Discussion 180

in 30 Doradus contains the luminous core R 136 whose effect was neglected. This means that our estimated underlying stellar absorption should probably be diluted by a greater factor (the denominator in Eq. 4.1) than what was assumed here.

4.6 Discussion

4.6.1 Abundance discrepancies and elemental ratios

It is clear from the above analysis that ORL versus CEL abundance discrepancies are present^ not only in planetary nebulae, but in Hll regions too. In the past, Esteban et al. (1998) found ORL/CEL discrepancy factors of 1.6 and 2.2, respectively, for O^^ and C^"*" in the Orion nebula [adopting C^‘^(A1908)/H+ CEL abundances from Walter et al. 1992]. Unpublished long-slit spectra from our group obtained at the AAT and ESC 1.52-m telescopes, show a factor of only 1.3 discrepancy for O^"*" in that nebula, but imply an ORL/CEL C^”*" discrepancy of 2.4, when compared with Walter et al.’s lU E results. Discrepancies have also been reported for M 17, where Peimbert, Storey & Torres-Peimbert (1993) found an ORL/CEL factor of ~ 1.7 for the O^^ ion. The

C III] A1908 line in that nebula has not been detected by the lUE^ so no compari­ son with the ORL abundance value has been possible. Esteban et al. (1999) also reported a factor of 1.6 abundance discrepancy for the total, recombination-line oxygen abundance versus the forbidden-line value in the Hll region M8. A stan­ dard abundance study of the Gum 38a complex, which contains NGC 3576, has been made by Girardi et al. (1997), but no detection of C II A4267 or any other heavy element optical recombination line was reported. NGC 3576 was never ob­ served by the lU E either, in order to allow a comparison with the ORL C^'*'/H+ abundance ratio reported herein. The present analysis shows that the ORL/CEL O^’*' discrepancy factors for our Hli regions are: 2.1, 1.7, 2.0-2.7, 4.9-S.2 and 2.3, for M 17, NGC 3576, 30 Do­ radus, LMC NllB and SMC N66 respectively. Even if we adopt the conservative 4-6. Discussion 181

ORL abundances for 30 Doradus and NllB (neglecting any underlying absorp­ tion), the mean discrepancy factor is 2.6. NllB especially, displays the largest discrepancy documented so far for an HII region (= 4.9), which is in the neighbor­ hood of the mean value (= 5.1) found for a sample of eighteen PN (Chapter 2). How is one to explain such an extreme discrepancy for an HII nebula? Before commenting on this issue let us have a look on the derived elemental abundance ratios for this sample of nebulae presented in Table 4.11 and compared with the solar values. The and values are from abundances derived from pure ORL ratios and should be almost completely unaffected by temperature variations along the line of sight. They should give a good measure of the C /0 and N /0 ratio—especially of the latter since the ionization potentials of nitrogen ionic stages are very similar to those of oxygen stages. From Table 4.11 we see that the agreement for the C^'''/0^'*‘ ratio amongst the three galactic nebulae is excellent and is only ~ 7 per cent higher than the solar value. We should stress here the fact that we adopt the recent solar oxygen abundance of Allende Prieto et al. (2001; O = 8.69 expressed as log O /H when H = 12) throughout this study. Esteban et al. (1998) using the former oxygen value of 8.87, found that the C /0 ratio in Orion is higher than the solar value (by ~ 17 per cent) and interpreted this as a result of Galactic chemical evolution; the revised solar O abundance implies that their C /0 ratio would actually be 22 per cent less than solar (and would not reveal any C /0 enhancement). Clearly, our galactic C /0 results agree much better with the solar ratio, pointing towards only a mild C /0 galactic increase in the last 4.5 Gyr. Our N^^/O^’*' ORL results for M 17 and NGC 3576 are higher than the respec­ tive N^/O ’*' CEL ratios by almost a factor of 7, for both nebulae. Disagreement is found between the same ORL and CEL ratio in 30 Doradus too. On the other hand the Orion N /0 ratios are in agreement. Our Orion N^'^/O^'*’ value may be favourably compared with the Rubin et al. (1998) N /0 value of 0.16, with the Rubin et al. (1991a) far-IR N^‘^/0^"'' ratio of ~0.2 and with the sun. 4-6. Discussion 182

Table 4.11: Elemental abundance ratios.

Sun“ Orion^ M17 NCC3576 30 Doradus LMC NllB SMC N 66

C 2 + /Q 2 + (ORLs) 0.72*’ 0.77 0.77 0.78 0.25 0.25 < 0 .2 0

N 2 + / o 2 + (ORLs) 0.19^ 0.15 0.62 0.71 <0.16 *■ *

N + /0 + (CELs) * 0.15 0.088 0.106 0.036 0.033 0.016

N e/O (CELs) 0.25 0.18 0.28 0.16 0.21 0.16 0.18

S / 0 (CELs) 0.033 0.028 0.027 * 0.027 0.019; 0.018

A r /0 (CELs) .0074 .011 .0057 .0078 .0065 .0078 .0053

“ Solar elemental ratios adopting O — 8.69 (Allende Prieto et al. 2001) and C, N, Ne, S, Ar

abundances from Grevesse et al. (1996);

These refer to solar C/O and N/O ratios respectively;

From Esteban et al. (1998), apart from the and N^'^'/O^"'" ratios which are from

unpublished data.

With regard to M 17 and NGC 3576, either the ORL values point towards N /0 ratios that are much higher than the mean galactic Hll region ratio (= 0.074, Du- four 1984), or there is the possibility that N II ORLs from these two nebulae yield unreliable abundances because they are affected by fluorescence effects. Grandi

(1976) showed that in Orion the N II multiplets V3, V5 and V 30 are excited by fluorescence via the He I A508.6 resonance line. In this study abun­ dances for M 17 and NGC 3576 have been derived from multiplets V 3 and V 5 (Table 4.5). It is therefore possible that fluorescence contributes to the excita­ tion of these lines. Unfortunately, the N II 3d-4f high-lying transitions that are generally not biased by this effect have not been detected from either of these two nebulae; their measurement would offer a means of checking the fluorescence hypothesis. For the two LMC nebulae the ORL values are in excellent agreement, being only 11 per cent smaller than the mean G/0 ratio in LMC Hll regions (= 0.28). Their N+/0''' CEL ratios are in excellent agreement with the LMC 4-6. Discussion 183

mean N/0 ratio of 0.035. Finally, regarding SMC N66 the derived ORL ratio is in very good agreement with the SMC mean C/O value of 0.15, while the N"^/0"^ ratio is only 46 per cent of the mean N /0 value of 0.035. The comparisons

are with respect to the LMC and SMC H II regions abundances of Kurt &: Dufour (1998).

4.6.2 Causes of the discrepancies

For M 17 and NGC 3576 the temperature fluctuation parameter, derived from

a comparison between the BJ and [O III] forbidden-line temperatures is 0.011 and 0.017, respectively (Table 4.2). On the other hand the ORL/CEL 0^+ discrepancy factors for these two nebulae (2.1 and 1.7 respectively) imply values of 0.038, 1.e. considerably larger. For the three Magellanic Cloud nebulae the implied from the abundance discrepancies is greater reaching the uncomfortable value of ~ 0.1 in the case of LMC N llB, which is very similar to the one needed to reconcile ORL and CEL O^'*' abundances for the rather extreme planetary nebula NGC 7009 (LSBC). Typical values predicted by photoionization models of chemically and density homogeneous nebulae yield t^~0.01 (Garnett 1992, Gruenwald & Viegas 1995). As in the case of PN, it is unlikely that simple temperature fluctuations are to blame for the discrepancies observed here, especially for objects like LMC N llB. High density clumps with > 10^ cm“^ embedded in a medium of lower density could lead to high ORL C, N, O abundances via an effect of pseudo­ temperature fluctuations, whereby the observationally derived [O III] temperature is significantly overestimated from collisional quenching of the nebular AA4959, 5007 lines (Viegas & Clegg 1992). HST images of the Orion nebula show numer­ ous clumpy microstructures surrounded by more or less uniform nebulosity. Walsh & Rosa (1998) reported HST observations of a partially ionized globule and a fil­ ament in the Orion nebula that display high densities of up to 5 x 10^ cm“^. They do not however disclose their density-sensitive diagnostic. They also comment that in the core of the nebula the filaments contribute half the [N ll] emission, but 4.6. Discussion 184 only 10 per cent of the Baimer line flux. Such high-density clumps would have a significant effect on the derived [O III] temperature and the fact that, according to the authors, they are not emitting strongly in HI Baimer lines would make their detection difficult via other means as we will discuss soon after. The existence of high-density, partially ionized condensations with ~ 10^ cm"^ has also been postulated by Bautista, Pradhan & Osterbrock (1994) in order to explain abnormalities in the emission spectrum of [Fe ll] in the Orion nebula. Independent observations by Esteban et al. (1998) and Baldwin et al. (2000) however, showed that [Fe ll] lines are formed in lower density gas along with [O l] lines. Finally, as Verner et al. (2000) discuss, the formation of [Fe ll] emission lines is affected by radiative pumping fluorescence processes in ways that render them unsuitable as straightforward density diagnostics in Hll regions. In Fig. 4.4 we plot the reddening-corrected intensities of high-order Baimer lines from our high resolution (1 Â) spectrum of NGC 3576 and compare them with theoretical predictions for various nebular electron densities. Since such lines are sensitive indicators of ionized high-density regions (assuming those have a normal hydrogen content), it is interesting to see that our data do not show any evidence for significant amounts of high-density material. In fact, the spectrum can be fitted with a uniform electron density of 2990 cm~^, rather consistent with that derived from the optical [Cl III] density diagnostic (Table 4.2). It is likely however that existing clumps are only partially ionized and so they do not emit significant amounts of Baimer line radiation just like Walsh & Rosa (1998) report for the Orion nebula dense inclusions. Our long-slit spectroscopy could easily miss to pick up such emission, especially since the ambient emission from material of low density ionized gas would be overwhelmingly higher. Unfortunately, we are not in a position to compute a detailed empirical model, incorporating high-density, small filling factor inclusions for NGC 3576. We lack emission lines that could serve as probes of high-density plasma (other than the

H 10-24 lines); there are no C III] A1908, [Ne III] 15.5-/im line observations. Both J^.6. Discussion 185 these lines arise from high critical density upper levels and would be invaluable in such a modelling attempt. Likewise, we have no high-resolution observations of the region around the Baimer discontinuity for any of the Magellanic Cloud nebulae. An examination of the high-order Baimer lines in LMC NllB, as well a measurement of the Baimer jump temperature of that nebula, would be of interest. Another explanation for the uniform, low density derived by the Baimer lines could involve the existence of high density, but hydrogen-poor material, enhanced in heavy elements so that its ORL emission rates would be high and lead to the observed abundance discrepancies. Rosa & Mathis (1987) report the existence of a hydrogen-poor (He/H = 0.14) region in the outskirts of the 30 Doradus nebula, enhanced in most metals, apart from nitrogen. They hypothesize that it might contain the chemically processed ejecta from a massive, late-type Wolf-Rayet star. It would be interesting to obtain high resolution spectra of that region and search for enhanced heavy element ORL emission and indications of high electron density. Our long-slit spectra sample a slice through the central 1 arcmin^ of 30 Doradus, which is known to contain a number of WR stars (Moffat et al. 1987, Parker 1993). Could it be that knots of heavy element enhanced material have been ejected from such stars and are the cause of the high ORL abundances? It is well known from kinematical studies (e.g. Meaburn 1984) that 30 Doradus contains pressure-driven bubbles, shells and sheets of ionized gas, originating in the intense winds from WR, Of stars and supernova remnants. In fact two such potential culprits have been directly encountered by our long slit, i.e. R 139 (Of -f- WNL) and R 140 (WN -f WC). On the other hand, LMC NllB does not contain any known Wolf-Rayet stars (Parker et al. 1992) and its abundance discrepancy factor is more than twice that of 30 Doradus. In conclusion, the current data set does not settle the issue for the existence of C II, N II, 0 II ORL emitting, high density ionized condensations in our sam­ ple of nebulae, even though they could offer a viable solution to the abundance discrepancy problem. Our analysis rather points towards a low density (cf. Sec- 4-6. Discussion 186

tion 4.5.2.1) and also very cold (so that it does not emit CELs) gas origin for these lines. From our optical density diagnostics, however, we have deduced the exis­

tence of strong density variations in these H II regions, especially within the Cloud nebulae. Furthermore, we can not firmly exclude the possibility that real tem­ perature fiuctuations cause part of the discrepancy. A future abundance analysis employing far-IR line fluxes would be invaluable in deriving and N^+/H+ fractions from IR CELs in order to compare them with the corresponding ORL abundances; this would demonstrate to what extent real temperature fluctuations exist in these nebulae. If, however, it turns out that abundances from IR CELs are in agreement with those from optical CELs, then one way towards a better understanding of the ORL/CEL problem may be to postulate the existence of cold, very rarefied, metal-rich gas in these H II regions. ^.6. Discussion 187

1.2 R139 = Parker 952 ■ 1.1

1.0

0 . 9 =i

0.8

41004 1 5 0 4200 4 2 5 0 4300 4 3 5 0 4 4 0 0 4 4 5 0 1.2 R139 = Parker 952 - 1.1

1.0 He II (WR)

0 . 9

0.8

4 0 0 4 5 5 045 4600 47004650 4 7 5 0 4 8 0 0 4 8 5 0 4 9 0 0 Wavelength (A)

Figure 4.1: Rectified linear-intensity, blue-violet spectrogram of the 06.5 Ib(f) 4- WNL supergiant

R 139 (= Parker 952). The identified features are: H6 A4102, H7 A4340, H/3 A4861, He I AA4121,

4144, 4387, 4471; He II AA4200, 4541, 4686 (broad WR emission + narrow absorption); N III A4379, (AA4638, 4640+4641 emission); Si iv AA4089, 4116; Mg II A4481; the diffuse interstellar bands at AA4430, 4502 and the unidentified emission lines at AA4485, 4503. See text for more details. 4.6. Discussion 188

R140

0 . 9 He II 0.8

0 . 7 40004200 43004100 4400 4500 4600 48004700 4900 Wavelength (A)

Figure 4.2: Rectified log-intensity, blue-violet spectrogram of the WN 6 (h) component of the spectroscopic binary R 140 - it may have a black hole companion (Wang 1995). The identified features are: H5 A4102, H 7 A4340, H/3 A4861, He I AA4026; He II AA4026, 4200, 4541, 4686; N III AA4638 + 4640 + 4641; N iv A4058; N v A4603, 4619; see text for more details.

1.1

O O

1.0

0 . 9

0.8 Parker 3157

4 5 5 0 4600 4650 4700 4 7 5 0 4 8 0 0 Wavelength (A)

Figure 4.3: Rectified, blue spectrogram of the blue supergiant Parker 3157 (BC 1 la) with a full complement of O ll V 1 lines; see text for more details. 4-6. Discussion 189

0.75 :rio NGC 3576

0.50 f H14

0.25

(cm ) . oWD 0 .0 0 h -3

-0.25

-0.50

1.0 1.1 1.2 1.3 1.4 Log n

Figure 4.4: Observed log-intensities (in units of H/5 = 100) of high-order Baimer lines (n ^ 2, n = 10, 24) as a function of the principal quantum number n; H 14 at 3721.94 Â is blended with the [S h i] A3721.63 line. The various curves show respectively the predicted Balmer decrements for electron densities from TVe = 10^ to 10^ cm“^. A constant temperature of 8070K, derived from the nebular continuum Balmer discontinuity, has been assumed in all cases. Ch a pter 5

Conclusions

In Chapter 2 ionic and total elemental abundances, relative to hydrogen, were derived for twelve Galactic and three Magellanic Cloud PN. The C, N and O abundances were derived from both collisionally excited lines and optical recombi­ nation lines. For ten Galactic and three Cloud nebulae lU E low resolution spectra were analyzed, along with optical spectra, while for eight Galactic nebulae ionic abundances were also derived from ISO fine-structure lines. Good agreement was found amongst abundances derived from the collisionally excited UV, optical and IR lines, after accounting for modest density variations (on average a contrast of a factor of 6) within the nebular volumes. In order, however, to reconcile the q2+/h+ ^1653 ionic abundance values with those derived from the optical AA4959, 5007 lines, much larger density contrasts may need to be called for, necessitating the existence of high density {Nq = 10^-10® cm~^), ionized inclusions. In all cases C, N, 0 abundances from ORLs were found to be consistently higher than the corresponding CEL values. For the first time ORL/CEL abundance discrepancies were documented for Magellanic Cloud PN. A further two extreme nebulae - NGC 2022 and LMC N66 - were revealed to have discrepancies greater than a factor of 10 for the 0^"*" ionic abundance. It was shown that the ORL/CEL discrepancies differ amongst PN spanning a range of more than 2 decades, from

190 191

~ 2 to > 20 for Even though quite similar elemental ratios were derived from pure ORL and pure CEL ratios, suggesting that any embedded dense inclusions originated from the same nuclear-processed, asymptotic giant branch material as the ambient gas, there is evidence that for some PN the C/O and N /0 ratios from ORLs are higher than those from CELs. This result, if real, may be indicative of enhanced nucleosynthetic effects in the postulated inclusions.

The high-excitation C II lines that populate the upper level of the C II A4267 3d-4f transition, were detected from a number of PN (as well as positively iden­ tified in the extensive line-list of the Orion nebula presented by Baldwin et al. 2000). Their relative intensities were found to be in good agreement with the predictions of recombination theory, thereby disproving claims that the A4267 line is augmented via means other than recombination; on the contrary the C II A4267 line is found to be a reliable abundance diagnostic. It was shown that the abundance enhancements of doubly ionized C, N, O are positively correlated with the difference (AT) between the Balmer recombination continuum and [O III] forbidden-line temperatures, suggesting that uncertainties in the adopted temperature caused by real or induced fluctuations, are partly to blame for the discrepancies. However, the relative uniformity of the overabundance patterns and lack of consistent correlation with CEL excitation energies, point away from Peimbert-type, simple temperature fluctuations as the cause of the problem. I confirm recent results by Garnett & Dinerstein (2000), showing that the ORL/CEL discrepancies are i) anticorrelated with the intrinsic nebular surface brightness and ii) positively correlated with PN absolute radii, indicating that young, bright nebulae display less abundance discrepancies than older, more ex­ tended ones. The results, however, further show that very similar correlations exist between AT and the nebular surface radii and brightnesses, suggesting that more evolved objects are more likely to have temperature discrepancies than less evolved ones. These findings strongly indicate an association of the ORL/CEL 192 abundance discrepancy problem with the evolution of planetary nebulae. Chapter 3 introduces empirical composite models for the planetary nebula NGC 5882, which shows mild ORL/CEL discrepancy factors of ~2. A model is favoured which fits well the various plasma and abundance indicators, by incor­ porating high-density, metal-enhanced inclusions of a small filling factor into a normal component of much lower density. The results show that high density condensations of the type suggested by Viegas & Clegg (1992) induce pseudo­ temperature variations, thus leading to an observationally derived low oxygen abundance from CELs, while the model’s oxygen value is substantially higher. The introduction, however, of a new solar oxygen abundance (Allende Prieto et al. 2001) led to the removal of the solar versus mean PN oxygen abundance dis­ crepancy; therefore it remains to be seen, whether PN heavy element abundances are truly higher than the mean nebular values derived from CELs. This will probably require detailed photoionization modelling and further observations. In Chapter 4 an ORL abundance analysis was presented for two Galactic and three Magellanic Cloud Hll regions, along with a standard CEL analysis. The abundance discrepancy factors are found, for the first time, to be in the range of 2 to 5 thus placing these objects in the planetary nebula ORL/CEL discrep­ ancy regime. For the first time ORL/CEL discrepancies were documented for extragalactic Hll regions els well. The nebula LMC N llB is found to display a large discrepancy factor ~ 5 for the O^'*' ion. Claims that H-poor nebular material originating from Wolf-Rayet stars is found in 30 Doradus should be investigated with respect to the ORL/CEL abundance discrepancy issue. Quite importantly, deviations are found in the observed relative intensities of O II ORLs from all nebulae that might be indicative of a low density gas origin for these lines. This could mean that the O II ORLs in this sample of H II regions are emitted from very rarefied metal-rich gas. Accurate, temperature insensitive C^’^/O^'*' ionic ratios are derived from ORLs for these objects; mean values are 0.77 for the galactic nebulae, 0.25 for the LMC nebulae and less than 0.20 for the SMC N66 nebula. 193

Future studies should aim at establishing whether clumpy regions in nebulae are really the sources of enhanced heavy element ORL emission. This can be pursued through high spatial resolution, spectroscopic observations of peculiar features, such as knots, filaments etc. These features could be first identified from high resolution imaging. A p p e n d i x A

Emission line fluxes: the catalogue

In this appendix the complete catalogue of measured emission line fluxes for the fifteen PN and five Hll regions studied in the context of this thesis is presented. The various c(H/3) logarithmic extinction coefficients used to deredden the ob­ served emission line fluxes are also given here, before each linelist. The tables contain from left to right the following: measured line wavelengths, wavelength specific extinction law, dereddened line intensities in units such that /(H/3) = 100, ionic identification (with ‘?’ for unidentified features), multiplet number, lower term of the transition, upper term of the transition, statistical weight of the lower level, statistical weight of the upper level. The observed fluxes relative to H/? can be retrieved as:

F{X) = I{\) io"(H/5)/(A),

The series of papers by Hyung et al. (1994, 1995, 1996, 1999) were used for line identification and guidance.

194 A.I. Galactic Planetary Nebulae 195

A .l Galactic Planetary Nebulae c(H/3)(Bal) = 0.375, c(H/3)(rad) - 0.545, c(H/3)^:G40 = ^ q7

Table A.l; NGC 2022 linelist

-^obs (•'^) f(A) /(A) Ion Ao Mult Lower term Upper term gi g2

1401 0.257 1.242e + 02 O IV] 1401 1486 0.257 9.010e+ 01 N IV] 1486 1549 0.257 8.139e + 02 C IV 1549 1574 0.257 2.127e + 01 [Ne V] 1574 1601 0.257 1.483e + 01 [Ne IV] 1601 1640 0.257 7.896e + 02 He II 1640 1663 0.257 1.774e + 01 O III] 1663 1908 0.257 5.255e + 02 C III] 1908 0.257 * [S III] 3721.63 F2 3p2 3? 3p2 IS 3 1 0.257 * H 14 3721.94 H14 2p+ 2P* 14d+ 2D 8 * 0.257 1.455e + 01 [O II] 3726.03 FI 2p3 4S* 2p3 2D* 4 4 0.256 * [O II] 3728.82 FI 2p3 4S* 2p3 2D* 4 6 0.255 * H 13 3734.37 H13 2p+ 2P* 13d+ 2D 8 * 3750.89 0.253 4.713e + 00 H 12 3750.15 H12 2p+ 2P* 12d+ 2D 8 * 3757.98 0.251 3.487e + 00 ? 3757.98 * * 3760.61 0.251 3.265e + 00 O III 3759.87 V2 3s 3P* 3p 3D 5 7 3771.37 0.249 5.135e + 00 H 11 3770.63 Hll 2p+ 2P* lld+ 2D 8 * 3797.88 0.244 6.687e + 00 H 10 3797.90 HIO 2p+ 2P* 10d+ 2D 8 * 3815.06 0.241 9.780e - 01 ? 3815.06 * * 3835.37 0.237 8.809e + 00 H 9 3835.39 H9 2p+ 2P* 9d+ 2D 8 * 3869.16 0.230 5.561e + 01 [Ne III] 3868.75 FI 2p4 3P 2p4 ID 5 5 3889.46 0.226 1.323e + 01 He I 3888.65 V2 2s 3S 3p 3P* 3 9 3922.34 0.220 7 .5 7 1 e -0 1 He II 3923.48 4.15 41+ 2F* 15g+ 2G 32 * 3967.88 0.211 3.187e + 01 [Ne HI] 3967.46 FI 2p4 3P 2p4 ID 3 5 4025.99 0.198 1.593e + 00 He I 4026.21 V18 2p 3P* 5d 3D 9 15 4060.39 0.191 1.848e - 01 [F IV] 4060.23 * * 4068.47 0.189 4.357e - 01 C HI 4067.87 ** 4069.57 0.189 5.707e - 01 C III 4068.97 * * 4070.90 0.189 7.448e - 01 C HI 4070.30 ** 4072.76 0.188 3.164e - 01 O II 4072.16 VIO 3p 4D* 3d 4F 6 8 4076.46 0.187 2.949e - 01 O II 4075.86 VIO 3p 4D* 3d 4F 8 10 4090.02 0.184 2.028e - 01 0 II 4089.29 V48a 3d 4F 41 G5* 10 12 4098.09 0.183 1.225e + 00 N III 4097.33 VI 3s 2S 3p 2P* 2 4 A.l. Galactic Planetary Nebulae 196

-^obs (Â) f(A) /(A) Ion Ao Mult Lower term Upper term gi g2 4102.21 0.182 2.802e + 01 H 6 4101.74 H6 2p+ 2P* 6d+ 2D 8 72 4119.75 0.178 8.347e - 02 O II 4119.22 V20 3p 4P* 3d 4D 6 8 4122.71 0.177 2.099e - 01 0 II 4121.46 V19 3p 4P* 3d 4P 2 2 4143.97 0.172 2.021e - 01 He I 4143.76 V53 2p IP* 6d ID 3 5 4163.78 0.168 1.466e - 01 ? 4163.78 * * 4185.98 0.163 9.405e - 02 0 II 4185.45 V36 3p’ 2F* 3d’ 2G 6 8 4187.43 0.162 3.323e - 01 C III 4186.90 V18 4f IF* 5g IG 7 9 4200.42 0.160 2.062e + 00 He II 4199.83 4.11 4f+ 2F* l lg + 2G 32 * 4228.17 0.153 5.899e - 01 N II 4227.74 V33 3p ID 4s IP* 5 3 4267.78 0.144 8.200e - 01 C II 4267.15 V6 3d 2D 4f 2F* 10 14 4276.62 0.142 1.855e - 01 0 II 4276.75 V67b 3d 4D 4f F3* 6 8 4341.00 0.127 5.111e + 01 H 5 4340.47 H5 2p+ 2P* 5d+ 2D 8 50 4350.13 0.125 1.167e - 01 O II 4349.43 V2 3s 4P 3p 4P* 6 6 4363.78 0.121 1.436e + 01 [0 III] 4363.21 F2 2p2 ID 2p2 IS 5 1 4367.46 0.120 2.779e - 01 O II 4366.89 V2 3s 4P 3p 4P* 6 4 4370.43 0.120 1.453e - 01 Ne II 4369.86 V56 3d 4F 4f 0j3i* 4 6 4377.24 0.118 1.927e - 01 ? 4377.24 ** 4379.98 0.118 4.722e - 01 N III 4379.11 V18 4f 2F* 5g 2G 14 18 4388.50 0.115 8.337e - 02 He I 4387.93 V51 2p IP* 5d ID 3 5 4409.86 0.110 5.075e - 02 Ne II 4409.30 V55e 3d 4F 4f 2j5i* 8 10 4415.46 0.109 1.247e - 01 O II 4414.90 V5 3s 2P 3p 2D* 4 6 4435.13 0.104 6.446e - 02 N III 4434.14 V2 3p 2P* 3d 2D 2 4 4449.43 0.101 7.612e - 02 O II 4448.19 V35 3p’ 2F* 3d’ 2F 8 8 4453.61 0.100 2.839e - 01 ? 4453.61 * * 4459.36 0.098 1.627e - 01 ? 4459.31 * * 4472.10 0.095 6.876e - 01 He I 4471.50 V14 2p 3P* 4d 3D 9 15 4476.25 0.094 9.357e - 02 O II 4477.90 V88 3d 2P 4f G3* 4 6 4511.52 0.086 1.594e - 01 N III 4510.91 V3 3s’ 4P* 3p’ 4D 4 6 4517.74 0.084 7.375e - 02 N III 4518.15 V3 3s’ 4P* 3p’ 4D 2 2 4535.41 0.080 6.979e - 02 N III 4534.58 V3 3s’ 4P* 3p’ 4D 6 6 4542.22 0.078 3.798e + 00 He II 4541.59 4.9 4f+ 2F* 9g+ 2G 32 * 4571.37 0.071 4.499e - 02 Mg I] 4571.10 3s2 IS 3s3p 3P* 1 3 4607.24 0.062 1.124e - 01 N II 4607.16 V5 3s 3P* 3p 3P 1 3 4610.33 0.062 5.776e - 02 O II 4609.44 V92a 3d 2D 4f F4* 6 8 4626.09 0.058 4.334e - 02 [Ar V] 4624.54 3s2 ID 3s2 IS 5 1 4634.78 0.056 5.609e - 01 N III 4634.14 V2 3p 2P* 3d 2D 2 4 4639.50 0.055 1.051e - 01 O II 4638.86 VI 3s 4P 3p 4D* 2 4 4641.28 0.054 9.296e - 01 N III 4640.64 V2 3p 2P* 3d 2D 4 6 A.l. Galactic Planetary Nebulae 197

'^obs (Â) f(A) 7(A) Ion Ao Mult Lower term Upper term gi g2 4642.45 0.054 3.659e - 01 O II 4641.81 VI 3s 4P 3p 4D* 4 6 4648.11 0.053 5.686e - 01 C III 4647.42 VI 3s 3S 3p 3P* 3 5 4649.82 0.052 3.674e - 01 O II 4649.13 VI 3s 4P 3p 4D* 6 8 4650.94 0.052 2.820e - 01 C III 4650.25 VI 3s 3S 3p 3P* 3 3 4651.53 0.052 8.923e - 02 O II 4650.84 VI 3s 4P 3p 4D* 2 2 4652.16 0.052 1.232e - 01 C III 4651.47 VI 3s 3S 3p 3P* 3 1 4659.05 0.050 1.030e + 00 C IV 4658.10 F3 3d6 5D 3d6 3F2 9 9 4662.20 0.049 1.093e - 01 O II 4661.63 VI 3s 4P 3p 4D* 4 4 4677.18 0.046 7.715e - 02 O II 4676.24 VI 3s 4P 3p 4D* 6 6 4681.24 0.044 2.206e - 01 ? 4681.24 * * 4686.41 0.043 1.128e -f 02 He II 4685.68 3.4 3d + 2D 41-1- 2F* 18 32 4712.08 0.037 1.435e-|-01 [Ar IV] 4711.37 FI 3p3 4S* 3p3 2D* 4 6 4715.20 0.036 1.073e -j- 00 He I 4713.17 V12 2P 3p* 4S 3s 9 3 4724.92 0.034 1.070e 4- 00 [Ne IV] 4724.15 FI 2p3 2D* 2p3 2P* 4 4 4726.39 0.033 8.650e - 01 [Ne IV] 4725.62 FI 2p3 2D* 2p3 2P* 4 2 4740.91 0.030 1.246e-H01 [Ar IV] 4740.17 FI 3p3 4S* 3p3 2D* 4 4 4862.13 0.000 l.OOOe -f 02 H 4 4861.33 H4 2p-b 2P* 4d+ 2D 8 32 4922.81 -0.015 2.257e - 01 He I 4921.93 V48 2p IP* 4d ID 3 5 4932.04 -0.017 1.335e - 01 [O HI] 4931.80 FI 2p2 3P 2p2 ID 1 5 4959.16 -0.024 2.465e + 02 [O HI] 4958.91 FI 2p2 3P 2p2 ID 3 5 5007.09 -0.036 7.351e 4- 02 [O HI] 5006.84 FI 2p2 3P 2p2 ID 5 5 5147.34 -0.070 3.591e - 01 7 5147.34 * * 5177.13 -0.077 1.306e - 01 7 5177.13 ** 5345.87 -0.118 1.348e - 01 ? 5345.87 * * 5411.95 -0.134 8.675e -f 00 He II 5411.52 4.7 4f+ 2F* 7g+ 2G 32 98 5518.34 -0.154 4.294e - 01 [Cl III] 5517.66 FI 2p3 4S* 2p3 2D* 4 6 5538.28 -0.158 3.581e - 01 [Cl HI] 5537.60 FI 2p3 4S* 2p3 2D* 4 4 5592.99 -0.167 2.387e - 01 O HI 5592.37 V5 3s IP* 3p IP 3 3 5801.82 -0.202 3.930e - 01 C IV 5801.51 VI 3s 2S 3p 2P* 2 4 5812.45 -0.204 2.378e - 01 C IV 5812.14 VI 3s 2S 3p 2P* 2 2 5868.54 -0.213 3.278e - 01 7 5868.54 * * 5876.42 -0.215 1.943e4-00 He I 5875.66 V l l 2p 3P* 3d 3D 9 15 5930.89 -0.223 1.070e - 01 He II 5931.78 ** 5953.59 -0.227 1.483e - 01 He II 5952.39 * * 5977.29 -0.231 1.405e - 01 He II 5977.03 * * 6006.09 -0.236 1.762e - 01 ? 6006.09 * * 6037.45 -0.241 1.875e - 01 He II 6036.70 * * 6074.81 -0.247 2 .3 4 5 e - 01 He II 6074.10 5.20 5g4- 2G 20h+ 2H* 50 * A.I. Galactic Planetary Nebulae 198

-^obs (^) f(A) /(A) Ion Ao Mult Lower term Upper term gi g2 6102.86 -0.251 6.585e - 01 [K IV] 6101.83 FI 3p4 3P 3d4 ID 5 5 6119.23 -0.254 2.235e - 01 He II 6118.20 5.19 5g+ 2G 19h-b 2H* 50 * 6172.20 -0.262 2.511e - 01 He II 6170.60 5.18 5g+ 2G 18h+ 2H 5 * 6234.40 -0.272 3.146e - 01 He II 6233.80 5.17 5g+ 2G 17h+ 2H 5 * 6313.08 -0.283 1.880e-h00 [S III] 6312.10 F3 2p2 ID 2p2 IS 5 1 6407.30 -0.297 4.917e - 01 He II 6406.30 5.15 5g+ 2G 15h+ 2H 5 * 6436.10 -0.302 1.806e4- 00 [Ar V] 6435.10 3s2 3P 3s2 ID 3 5 6528.47 -0.315 7.251e - 01 He II 6527.11 5.14 5g-b 2G 14h+ 2H 5 * 6549.22 -0.318 1.060e4-00 [N II] 6548.10 FI 2p2 3P 2p2 ID 3 5 6563.89 -0.320 3.008e -f 02 H 3 6562.77 H3 2p+ 2P* 3d-b 2D 8 18 6584.62 -0.323 2.285e -b 00 [N II] 6583.50 FI 2p2 3P 2p2 ID 5 5 6679.86 -0.336 4.761e - 01 He I 6678.16 V46 2p IP* 3d ID 3 5 6684.56 -0.337 7.274e - 01 He II 6683.20 ** 6717.81 -0.342 6.557e - 01 [S II] 6716.44 F2 2p3 4S* 2p3 2D* 4 6 6732.19 -0.344 7.316e - 01 [S II] 6730.82 F2 2p3 4S* 2p3 2D* 4 4 6796.37 -0.352 1.937e - 01 [K IV] 6795.00 FI 3p4 3P 3p4 ID 3 5 6828.16 -0.357 7.166e - 01 ? 6828.16 * * 6862.86 -0.361 2.132e - 01 ? 6862.86 ** 6892.24 -0.365 7.764e - 01 He II 6890.88 * * 7007.02 -0.380 3.647e -b 00 [ArV] 7005.67 3s2 3P 3s2 ID 5 5 7066.61 -0.387 4.597e - 01 He I 7065.25 VIO 2p 3P* 3s 3S 9 3 7136.82 -0.396 9.5136 “b 00 [Ar III] 7135.80 FI 3p4 3P 3p4 ID 5 5 7171.65 -0.401 5.114e - 01 [Ar IV] 7170.62 F2 3p3 2D* 3p3 2P* 4 4 7178.53 -0.401 1.083e-b00 He II 7177.50 5.11 5g-b 2G llh -b 2H 5 * 7238.48 -0.409 5.599e - 01 [Ar IV] 7237.26 F2 3p3 2D* 3p3 2P* 6 4 7263.99 -0.412 4.181e - 01 [Ar IV] 7262.76 F2 3p3 2D* 3p3 2P* 4 2 7320.46 -0.418 3.476e - 01 [O II] 7318.92 F2 2p3 2D* 2p3 2P* 6 2 7331.22 -0.420 3.598e - 01 [O II] 7329.67 F2 2p3 2D* 2p3 2P* 4 2

c(H^)(Bal) = 0.456, c(H/3)(rad) = 0.516 , c(H/3)^iG40 = 0.746

Table A.2: NGC2440 linelist.

'^obs (^) f(A) /(A) Ion Ao Mult Lower term Upper term gi g2

1240 0.260 1.048e -b 02 N V 1240 1401 0.260 5.290e -b 01 O IV] 1401 1484 0.260 1.856e-b02 N IV] 1484 1549 0.260 6.589e -b 02 C IV 1549 1574 0.260 1.2216-b 01 [Ne V] 1574 1640 0.260 5.1436-b 02 He II 1640 A.I. Galactic Planetary Nebulae 199

-^obs (Â) f(A) /(A) Ion Ao Mult Lower term Upper term gi g2 1663 0.260 3.920e + 01 O III] 1663 1750 0.260 1.396e + 02 N III] 1750 1908 0.260 9.235e + 02 C III] 1908 4004.83 0.203 5.879e - 02 N III 4003.58 V17 4d 2D 5f 2F* 6 8 4010.20 0.202 1.676e - 01 He I 4009.26 V55 2p IP* 7d ID 3 5 4026.89 0.198 2.372e + 00 He I 4026.21 V18 2p 3P* 5d 3D 9 15 4042.29 0.195 6.184e - 02 N II 4041.31 ** 4061.54 0.191 7.959e - 02 [FIV] 4060.00 * * 4069.36 0.189 2.438e + 00 [S II] 4068.60 FI 2p3 4S* 2p3 2P* 4 4 4070.52 0.189 5.789e - 01 O II 4069.62 VIO 3p 4D* 3d 4F 2 4 0.189 * O II 4069.89 VIO 3p 4D* 3d 4F 4 6 4072.92 0.188 1.512e - 01 0 II 4072.16 VIO 3p 4D* 3d 4F 6 8 0.187 * 0 II 4075.86 VIO 3p 4D* 3d 4F 8 10 4077.11 0.187 8.504e - 01 [S II] 4076.35 FI 2p3 4S* 2p3 2P* 2 4 0.183 * 0 II 4097.26 V48b 3d 4F 4f G4* 8 10 4098.23 0.183 2.841e + 00 N III 4097.33 VI 3s 2S 3p 2P* 2 4 4102.57 0.182 2.762e + 01 H 6 4101.74 H6 2p+ 2P* 6d+ 2D 8 72 4121.46 0.178 1.719e - 01 O II 4119.22 V20 3p 4P* 3d 4D 6 8 0.178 * 0 II 4120.28 V20 3p 4P* 3d 4D 6 6 0.177 * O II 4120.54 V20 3p 4P* 3d 4D 6 4 4123.08 0.177 1.113e - 01 He I 4120.84 V16 2p 3P* 5s 3S 9 3 4123.70 0.177 3.595e - 02 O II 4121.46 V19 3p 4P* 3d 4P 2 2 4129.82 0.176 6.911e - 02 O II 4129.32 V19 3p 4P* 3d 4P 4 2 4144.57 0.172 2.943e - 01 He I 4143.76 V53 2p IP* 6d ID 3 5 4164.05 0.168 1.305e - 01 [K V] 4163.12 ** 0.163 * 0 II 4185.45 V36 3p’ 2F* 3d’ 2G 6 8 4187.72 0.162 2.218e - 01 C HI 4186.90 V18 4f IF* 5g IG 7 9 4196.65 0.160 6.347e - 02 N HI 4195.76 V6 3s’ 2P* 3p’ 2D 2 4 4200.80 0.160 1.582e + 00 He II 4199.83 4.11 41+ 2F* l lg + 2G 32 * 0.160 * N HI 4200.10 V6 3s’ 2P* 3p’ 2D 4 6 4228.22 0.153 1.857e - 01 N II 4227.74 V33 3p ID 4s IP* 5 3 0.151 * N II 4236.91 V48a 3d 3D* 4f l\3i 3 5 4237.63 0.151 4.951e - 02 N II 4237.05 V48b 3d 3D* 4f lj4i 5 7 4242.36 0.150 5.579e - 02 N II 4241.24 V48a 3d 3D* 4f 5 5 0.150 * N II 4241.78 V48a 3d 3D* 4f l\3i 5 7 4268.02 0.144 4.026e - 01 C II 4267.15 V6 3d 2D 4f 2F* 10 14 4341.36 0.127 4.875e + 01 H 5 4340.47 H5 2p+ 2P* 5d+ 2D 8 50 0.125 * 0 II 4345.56 V2 3s 4P 3p 4P* 4 2 A.I. Galactic Planetary Nebulae 200

^obs (Â) f(A) 7(A) Ion Ao Mult Lower term Upper term gi g2 4364.13 0.121 2.517e + 01 [O III] 4363.21 F2 2p2 ID 2p2 IS 5 1 4380.10 0.118 3.239e - 01 N III 4379.11 V18 4f 2F* 5 g 2 G 14 18 0.115 3.577e - 01 He I 4387.93 V51 2p IP* 5d ID 3 5 4415.42 0.109 4.889e - 02 O II 4414.90 V5 3s 2P 3p 2D* 4 6 4417.49 0.109 3.935e - 02 O II 4416.97 V5 3s 2P 3p 2D* 2 4 4438.37 0.103 4.725e - 02 He I 4437.55 V50 2p IP* 5s IS 3 1 4453.94 0.100 8.317e - 02 O II 4452.37 V5 3s 2P 3p 2D* 4 4 4459.44 0.098 3.844e - 02 7 *** 4472.45 0.095 3.295e + 00 He I 4471.50 V14 2p 3P* 4d 3D 9 15 4499.49 0.089 4.765e - 02 ? * **

4511.91 0.086 2.615e - 01 N HI 4510.91 V3 3s’ 4P* 3p’ 4D 4 6 4515.86 0.085 8.010e - 02 N HI 4514.86 V3 3s’ 4P* 3p’ 4D 6 8 4519.15 0.084 1.040e - 01 N HI 4518.15 V3 3s’ 4P* 3p’ 4D 2 2 4524.55 0.083 8.185e - 02 N HI 4523.58 V3 3s’ 4P* 3p’ 4D 4 4 4531.38 0.081 4.152e - 02 N II 4530.41 V58b 3d IF* 4f 2\bi 7 9 4535.55 0.080 8.911e - 02 N HI 4534.58 V3 3s’ 4P* 3p’ 4D 6 6 4542.56 0.078 2.501e + 00 He II 4541.59 4.9 4f+ 2F* 9g+ 2G 32 * 4563.57 0.073 1.123e - 01 Mg I] 4562.60 3s2 IS 3s3p 3P* 1 5 4572.06 0.071 2.224e - 01 Mg I] 4571.10 3s2 IS 3s3p 3P* 1 3 4590.21 0.066 1.127e - 01 O II 4590.97 V15 3s’ 2D 3p’ 2F* 6 8 4604.75 0.063 1.815e - 01 ? 4604.75 * * 4607.44 0.062 4.750e - 01 7 4607.44 * * 4620.02 0.059 2.542e - 02 ? 4620.02 ** 4621.31 0.059 5.543e - 02 N II 4621.39 * * 4632.99 0.056 2.998e - 01 ? 4632.99 * * 4635.17 0.056 2.199e + 00 N HI 4634.14 V2 3p 2P* 3d 2D 2 4 4639.89 0.055 2.841e - 01 O II 4638.86 VI 3s 4P 3p 4D* 2 4 4641.67 0.054 4.012e + 00 N HI 4640.64 V2 3p 2P* 3d 2D 4 6 4642.84 0.054 6.5246 - 01 O II 4641.81 V I 3s 4P 3p 4D* 4 6 4648.35 0.053 2.263e -01 C HI 4647.42 V I 3s 3S 3p 3P* 3 5 4650.06 0.052 1.3136-01 O II 4649.13 VI 3s 4P 3p 4D* 6 8 4651.18 0.052 1.024e - 01 C HI 4650.25 VI 3s 3S 3p 3P* 3 3 4651.77 0.052 5.2876 - 02 O II 4650.84 VI 3s 4P 3p 4D* 2 2 4652.40 0.052 4.8536 - 02 C HI 4651.47 V I 3s 3S 3p 3P* 3 1 4659.26 0.050 7.2176 - 01 C IV 4658.10 F3 3d6 5D 3d6 3F2 9 9 4662.55 0.049 3.5886 — 02 O II 4661.63 V I 3s 4P 3p 4D* 4 4 4666.59 0.048 3.0536 - 02 7 4666.59 * * 0.046 * O II 4673.73 V I 3s 4P 3p 4D* 4 2 A.I. Galactic Planetary Nebulae 201

Aobs(Â) %A) /(A) Ion Ao Mult Lower term Upper term gi g2 4677.19 0.046 7.364e - 02 O II 4676.24 VI 3s 4P 3p 4D* 6 6 4686.76 0.043 7.234e -H 01 He II 4685.68 3.4 3d-|- 2D 4f-H 2F* 18 32 0.041 * O II 4696.35 VI 3s 4P 3p 4D* 6 4 4702.71 0.039 4.190e - 02 [Fe III] 4701.62 F3 3d6 5D 3d6 3F2 7 7 4708.43 0.038 1.167e - 01 ? 4708.43 * * 4712.45 0.037 8.399e -H 00 [Ar IV] 4711.37 FI 3p3 4S* 3p3 2D* 4 6 4714.25 0.036 1.279e-H00 He I 4713.17 vl2 2P 3p* 4S 3s 9 3 4725.26 0.034 8.991e - 01 [Ne IV] 4724.15 FI 2p3 2D* 2p3 2P* 4 4 4726.73 0.033 7.422e - 01 [Ne IV] 4725.62 FI 2p3 2D* 2p3 2P* 4 2 4741.34 0.030 9.429e -t- 00 [Ar IV] 4740.17 FI 3p3 4S* 3p3 2D* 4 4 4803.70 0.015 9.179e - 02 N II 4802.29 V20 3p 3D 3d 3D* 7 7 4853.99 0.002 8.569e - 02 ? 4853.99 * * 4862.50 0.000 1.037e 4- 02 H 4 4861.33 H4 2p+ 2P* 4d+ 2D 8 32 4867.67 -0.002 2.476e - 01 7 4867.67 * * 4882.17 -0.005 5.168e - 02 [Fe HI] 4881.11 * * 4907.95 -0.011 2.711e - 02 O II 4906.83 V28 3p 4S* 3d 4P 4 4 4923.20 -0.015 8.208e - 01 He I 4921.93 V48 2p IP* 4d ID 3 5 -0.016 * O II 4924.53 V28 3p 4S* 3d 4P 4 6 4932.11 -0.017 3.543e - 01 [O HI] 4931.80 FI 2p2 3P 2p2 ID 1 5 4935.80 -0.018 3.051e - 02 ? 4934.63 ** 4939.50 -0.019 2.446e - 02 7 *** 4945.84 -0.021 2.864e - 01 ? 4945.84 * * 4960.13 -0.024 3.748e-h02 [O HI] 4958.91 FI 2p2 3P 2p2 ID 3 5

c(H/3)(Bal) = 0.200, c(H/3)(rad) = 0.331, c(H/3)^i64o ^ g.643

Table A.3: NGC3132 linelist.

Aobs {A) f(A) /(A) Ion Ao Mult Lower term Upper term gi g2

1549 0.260 7.400e + 00 C IV 1549 1640 0.260 2.590e -f 01 He II 1640 1663 0.260 3.000e -t- 00 O HI] 1663 1751 0.260 4.600e + 00 N HI] 1751 1908 0.260 4.110e4-01 C HI] 1908 0.259 * H 15 3711.97 ** 0.257 * H 14 3721.94 * * 0.257 * [S HI] 3721.63 * * 0.257 5.516e -|- 02 [O II] 3726.03 ** 0.256 * [0 II] 3728.82 * * 0.255 * H 13 3734.37 * * A.I. Galactic Planetary Nebulae 202

Aobs (-^) f(A) /(A) Ion Ao Mult Lower term Upper term gi g2 0.253 * H 12 3750.15 H12 2p+ 2P* 12d+ 2D 8 * 3797.32 0.251 * O III 3759.87 V2 3s 3P* 3p 3D 5 7 3797.32 0.244 5.114e + 00 H 10 3797.90 HIO 2p+ 2P* 10d+ 2D 8 * 3867.99 0.230 1.182e + 02 [Ne III] 3868.75 FI 2p4 3P 2p4 ID 5 5 3888.03 0.226 2.476e + 01 He I 3888.65 V2 2s 3S 3p 3P* 3 9 3967.30 0.211 5.069e + 01 [Ne III] 3967.46 FI 2p4 3P 2p4 ID 3 5 4009.21 0.202 3.077e - 01 He I 4009.26 V55 2p IP* 7d ID 3 5 4025.89 0.198 2.917e + 00 He I 4026.21 V18 2p 3P* 5d 3D 9 15 4068.20 0.189 6.288e + 00 [S II] 4068.60 FI 2p3 4S* 2p3 2P* 4 4 4069.36 0.189 4.086e - 01 O II 4069.62 VIO 3p 4D* 3d 4F 2 4 0.189 * O II 4069.89 VIO 3p 4D* 3d 4F 4 6 4072.07 0.188 1.484e - 01 O II 4072.16 VIO 3p 4D* 3d 4F 6 8 4075.46 0.187 3.592e - 01 O II 4075.86 VIO 3p 4D* 3d 4F 8 10 4075.95 0.187 1.912e + 00 [S II] 4076.35 FI 2p3 4S* 2p3 2P* 2 4 4084.95 0.185 4.328e - 02 O II 4085.11 VIO 3p 4D* 3d 4F 6 6 0.183 * 0 II 4097.25 V20 3p 4P* 3d 4D 2 4 0.183 * O II 4097.26 V48b 3d 4F 4f G4* 8 10 4096.95 0.183 3.624e - 01 N HI 4097.33 VI 3s 2S 3p 2P* 2 4 4101.33 0.182 2.580e + 01 H 6 4101.74 H6 2p+ 2P* 6d+ 2D 8 72 4118.80 0.178 2.417e - 02 O II 4119.22 V20 3p 4P* 3d 4D 6 8 0.178 * O II 4120.28 V20 3p 4P* 3d 4D 6 6 0.177 * O II 4120.54 V20 3p 4P* 3d 4D 6 4 4120.42 0.177 2.714e - 01 He I 4120.84 V16 2p 3P* 5s 3S 9 3 0.177 * 0 II 4121.46 V19 3p 4P* 3d 4P 2 2 4143.38 0.172 3.784e - 01 He I 4143.76 V53 2p IP* 6d ID 3 5 4153.25 0.170 7.865e - 02 O II 4153.30 V19 3p 4P* 3d 4P 4 6 4168.53 0.167 7.632e - 02 O II 4169.22 V19 3p 4P* 3d 4P 6 6 4266.77 0.144 6.971e - 01 C II 4267.15 V6 3d 2D 4f 2F* 10 14 4276.63 0.142 5.8396 — 02 O II 4275.55 V67a 3d 4D 4f F4* 8 10 4303.36 0.135 1.0136-01 O II 4303.61 V65a 3d 4D 4f G5* 8 10 4340.05 0.127 4.7126 + 01 H 5 4340.47 H5 2p+ 2P* 5d+ 2D 8 50 4345.32 0.125 7.922e - 02 O II 4345.56 V2 3s 4P 3p 4P* 4 2 4349.19 0.125 5.7346 - 02 0 II 4349.43 V2 3s 4P 3p 4P* 6 6 4362.82 0.121 4.3016 + 00 [0 HI] 4363.21 F2 2p2 ID 2p2 IS 5 1 4365.91 0.120 4.3396 - 02 O II 4366.89 V2 3s 4P 3p 4P* 6 4 4371.84 0.119 3.2156 - 02 O II 4371.62 V76b 3d 2F 4f G4* 8 10 4378.20 0.118 5.7126 - 02 7 4378.20 * * 4387.59 0.115 7.6796 - 01 He I 4387.93 V51 2p IP* 5d ID 3 5 A.I. Galactic Planetary Nebulae 203

'^obs (Â) f(A) /(A) Ion Ao Mult Lower term Upper term gi g2 4392.68 0.115 4.090e - 02 Ne II 4391.99 V55e 3d 4F 4f 2j5i* 10 12 4436.95 0.103 6.577e - 02 He I 4437.55 V50 2p IP* 5s IS 3 1

4438.30 0.103 4.512e - 02 ? 4438.30 * * 4471.10 0.095 6.144e -f 00 He I 4471.50 V14 2p 3P* 4d 3D 9 15 4491.12 0.091 5.2906 — 02 O II 4491.23 V86a 3d 2P 4f D3* 4 6 4541.26 0.078 1.177e - 01 He II 4541.59 4.9 4f+ 2F* 9g+ 2G 32 * 4562.25 0.073 3.446e - 01 Mg I] 4562.60 3s2 IS 3s3p 3P* 1 5 4570.75 0.071 3.163e - 01 Mg I] 4571.10 3s2 IS 3s3p 3P* 1 3 4599.08 0.064 5.130e - 02 ? 4599.08 * * 4621.36 0.059 3.301e - 02 N II 4621.39 * * 4630.61 0.057 4.107e - 02 N II 4630.54 V5 3s 3P* 3p 3P 5 5 * 4633.11 0.056 6.146e - 02 ? 4633.11 * 4712.90 0.056 6.172e - 01 N III 4634.14 V2 3p 2P* 3d 2D 2 4 4638.39 0.055 1.043e - 01 O II 4638.86 VI 3s 4P 3p 4D* 2 4 4640.17 0.054 2.437e - 01 N HI 4640.64 V2 3p 2P* 3d 2D 4 6 4641.36 0.054 2.940e - 01 O II 4641.81 VI 3s 4P 3p 4D* 4 6 4648.91 0.052 1.682e - 01 O II 4649.13 VI 3s 4P 3p 4D* 6 8 4650.62 0.052 1.227e - 01 O II 4650.84 VI 3s 4P 3p 4D* 2 2 4658.09 0.050 2.038e - 01 [Fe HI] 4658.10 F3 3d6 5D 3d6 3F2 9 9 4661.34 0.049 l.llOe - 01 0 II 4661.63 VI 3s 4P 3p 4D* 4 4 4673.30 0.046 6.069e - 02 O II 4673.73 VI 3s 4P 3p 4D* 4 2 4675.81 0.046 6.740e - 02 O II 4676.24 VI 3s 4P 3p 4D* 6 6 4685.40 0.043 4.018e -t- 00 He II 4685.68 3.4 3d-t- 2D 41+ 2F* 18 32 4695.78 0.041 2 .8 8 1 e - 02 O II 4696.35 VI * * 4701.45 0.039 5.009e - 02 [Fe HI] 4701.62 F3 3d6 5D 3d6 3F2 7 7 4711.10 0.037 4.573e - 01 [Ar IV] 4711.37 FI 3p3 4S* 3p3 2D* 4 6 0.036 * He I 4713.17 v l2 2P 3p* 4S 3s 9 3 4739.99 0.030 3.438e - 01 [Ar IV] 4740.17 FI 3p3 4S* 3p3 2D* 4 4 * 4755.26 0.026 3.744e - 02 [Fe HI] 4754.83 * 4769.16 0.023 3.389e - 02 [Fe III] 4769.40 3F * * 4777.53 0.021 3.537e - 02 [Fe HI] 4777.88 V20 ** 0.000 l.OOOe -f 02 H4 4861.33 H4 2p-k 2P* 4d+ 2D 8 32 ** 4866.06 -0.001 3.588e - 02 7 * 4881.00 -0.005 8.072e - 02 [Fe III] 4881.11 F2 ** 4893.79 -0.008 3.532e - 02 [Fe VII] 4893.90 F2 * * 4910.55 -0.012 4.432e - 02 7 *** 4921.78 -0.015 1.554e-h 00 He I 4921.93 V48 2p IP* 4d ID 3 5 4924.38 -0.016 7.075e - 02 O II 4924.53 V28 3p 4S* 3d 4P 4 6 A.I. Galactic Planetary Nebulae 204

•^obs(Â) f(A) /(A) Ion Ao Mult Lower term Upper term gi g2 4930.70 -0.017 8.918e - 02 [O III] 4931.80 FI 2p2 3P 2p2 ID 1 5 4932.31 -0.017 4.970e - 02 ? 4932.31 * * 4958.07 -0.024 2.679e -f 02 [O III] 4958.91 FI 2p2 3P 2p2 ID 3 5 5005.99 -0.036 7.899e + 02 [O III] 5006.84 FI 2p2 3P 2p2 ID 5 5 5190.28 -0.081 6.825e - 01 [Ar III] 5191.82 * * 5198.30 -0.083 1.048e -f 01 [NI] 5199.84 FI ** 5282.08 -0.103 3.532e - 01 ? 5282.08 * * 5410.88 -0.134 3.656e — 01 He II 5411.52 4.7 4f-b 2F* 7g+ 2 0 32 98 5517.08 -0.154 1.112e -b 00 [Cl III] 5517.66 FI 2p3 4S* 2p3 2D* 4 6 5537.01 -0.158 9.053e - 01 [Cl III] 5537.60 FI 2p3 4S* 2p3 2D* 4 4 5576.60 -0.164 6.834e - 01 [0 1] 5577.34 F3 2p4 ID 2p4 IS 5 1 5754.04 -0.195 8.155e-b00 [N II] 5754.60 F3 2p2 ID 2p2 IS 5 1 5875.09 -0.215 1.632e-b01 He I 5875.66 V ll 2p 3P* 3d 3D 9 15 6300.10 -0.282 3.705e 4- 01 [O I] 6300.34 FI 2p4 3P 2p4 ID 5 5 6311.86 -0.283 2.404e -b 00 [S III] 6312.10 F3 2p2 ID 2p2 IS 5 1 6363.54 -0.291 1.289e-b01 [O I] 6363.78 ** 6547.80 -0.318 2.055e + 02 [N II] 6548.10 FI 2p2 3P 2p2 ID 3 5 6562.47 -0.320 2.869e + 02 H 3 6562.77 H3 2p-b 2P* 3d-b 2D 8 18 6583.20 -0.323 6.420e -b 02 [N II] 6583.50 FI 2p2 3P 2p2 ID 5 5 6678.03 -0.336 4.459e -b 00 He I 6678.16 V46 2p IP* 3d ID 3 5 6716.31 -0.342 5.4886 -b 01 [S II] 6716.44 F2 2p3 4S* 2p3 2D* 4 6 6730.69 -0.344 5.389e -b 01 [S II] 6730.82 F2 2p3 4S* 2p3 2D* 4 4 7064.61 -0.387 3.822e -b 00 He I 7065.25 VIO 2p 3P* 3s 3S 9 3 7135.15 -0.396 2.534e -b 01 [Ar III] 7135.80 FI 3p4 3P 3p4 ID 5 5 7280.47 -0.414 8.474e - 01 He I 7281.35 V45 2p IP* 3s IS 3 1 7318.65 -0.418 6.686e -b 00 [O II] 7318.92 F2 2p3 2D* 2p3 2P* 6 2 7329.40 -0.420 5.650e -b 00 [0 II] 7329.67 F2 2p3 2D* 2p3 2P* 4 2

c(H/3)(Bal) = 0.100, c(H/3)(rad) = 0.198, c(H/3)^i640 = Q.450

Table A.4; NGC3242 linelist.

■^obs (^ ) f(A) /(A) Ion Aq Mult Lower term Upper term gi g2

1640 0.260 1.723e -b 02 He II 1640 1663 0.260 8.5006 -b 00 O III] 1663 1751 0.260 5.5006 -b 00 N HI] 1751 1908 0.260 1.4976 + 02 C HI] 1908 3726.73 0.257 1.4466 -b 01 [0 II] 3726.03 * * 0.256 * [0 II] 3728.82 * * 3750.21 0.253 3.5516-b 00 H 12 3750.15 H12 2p+ 2P* 12d-b 2D 8 * A.I. Galactic Planetary Nebulae 205

-^obs (Â) f(A) /(A) Ion Ao Mult Lower term Upper term gi g2 3757.30 0.251 1.517e + 00 O III 3759.87 V2 * * 3759.93 0.251 2.504e + 00 O III 3759.87 V2 3s 3P* 3p 3D 5 7 3770.69 0.249 4.288e + 00 H 11 3770.63 Hll 2p+ 2P* l ld + 2D 8 * 3797.35 0.244 5.467e + 00 H 10 3797.90 HIO 2p+ 2P* 10d+ 2D 8 * 3819.00 0.240 1.099e + 00 He I 3819.62 V22 2p 3P* 6d 3D 9 15 3834.91 0.237 7.434e + 00 H 9 3835.39 H9 2p+ 2P* 9d+ 2D 8 * 3868.37 0.230 1.039e + 02 [Ne III] 3868.75 FI 2p4 3P 2p4 ID 5 5

3 8 8 8 .6 7 0.226 1.980e + 01 He I 3888.65 V2 2s 38 3p 3P* 3 9 3923.65 0.220 2.134e - 01 He II 3923.48 4.15 ** 3967.07 0.211 4.515e + 01 [Ne III] 3967.46 FI 2p4 3P 2p4 ID 3 5 3998.55 0.204 2.033e - 02 N III 3998.63 V17 4d 2D 5f 2F* 4 6 4003.60 0.203 2.377e - 02 N III 4003.58 V17 4d 2D 5f 2F* 6 8 4009.20 0.202 1.731e - 01 He I 4009.26 V55 2p IP* 7d ID 3 5 4012.03 0.201 1.597e - 02 ? * * * 4026.02 0.198 2.268e + 00 He I 4026.21 V18 2p 3P* 5d 3D 9 15 4041.28 0.195 1.274e - 02 N II 4041.31 * * 4060.55 0.191 6.482e - 03 [FIV] 4060.00 * * 4067.67 0.189 1.877e - 01 C III 4067.94 4f 3F* 5g 3G * * 4068.40 0.189 2.400e - 02 [S II] 4068.60 FI 2p3 48* 2p3 2P* 4 4 4068.77 0.189 2.459e - 01 C III 4068.91 ** 4069.56 0.189 1.828e - 01 O II 4069.62 VIO 3p 4D* 3d 4F 2 4 0.189 * O II 4069.89 VIO 3p 4D* 3d 4F 4 6 4070.10 0.189 3.210e - 01 C III 4070.26 ** 0.188 * O II 4071.23 V48a 3d 4F 4f G5* 8 10 4071.96 0.188 2.153e - 01 O II 4072.16 VIO 3p 4D* 3d 4F 6 8 4075.66 0.187 1.664e - 01 O II 4075.86 VIO 3p 4D* 3d 4F 8 10 4076.15 0.187 7.892e - 03 [S II] 4076.35 FI 2p3 48* 2p3 2P* 2 4 4078.58 0.187 2 .1 2 6 e -0 2 O II 4078.84 VIO 3p 4D* 3d 4F 4 4 4080.84 0.186 1.394e - 02 O III 4081.10 DAS ** 4083.59 0.186 2.108e - 02 O II 4083.90 V48b 3d 4F 4f G4* 6 8 4084.79 0.185 3.509e - 02 O II 4085.11 VIO 3p 4D* 3d 4F 6 6 4086.84 0.185 2.043e - 02 O II 4087.15 V48c 3d 4F 4f G3* 4 6 4088.97 0.184 1.113e - 01 O II 4089.29 V48a 3d 4F 4f G5* 10 12 4093.04 0.184 2.402e - 02 O II 4092.93 VIO 3p 4D* 3d 4F 8 8 0.183 * O II 4097.25 V20 3p 4P* 3d 4D 2 4 0.183 * O II 4097.26 V48b 3d 4F 4f G4* 8 10 4097.11 0.183 1.670e + 00 N III 4097.33 VI 3s 28 3p 2P* 2 4

0.182 * O II 4098.24 V46a 3d 4F 4f D3* 4 6 A.I. Galactic Planetary Nebulae 206

•^obs (Â) f(A) /(A) Ion Ao Mult Lower term Upper term gi g2 4101.52 0.182 2.661e + 01 H 6 4101.74 H6 2p+ 2P* 6d+ 2D 8 72 4101.52 0.182 2.596e + 01 H 6 4101.74 H6 2p+ 2P* 6d+ 2D 8 72 0.180 * O II 4110.78 V20 3p 4P* 3d 4D 4 2 4115.52 0.178 1.976e - 02 ? 4116.14 ** 4118.98 0.178 6.772e - 02 O II 4119.22 V20 3p 4P* 3d 4D 6 8 0.178 * 0 II 4120.28 V20 3p 4P* 3d 4D 6 6 0.177 * O II 4120.54 V20 3p 4P* 3d 4D 6 4 4120.60 0.177 1.9586 — 01 He I 4120.84 V16 2p 3P* 5s 3S 9 3 4121.22 0.177 2.833e - 02 O II 4121.46 V19 3p 4P* 3d 4P 2 2 4128.49 0.176 5.166e — 02 0 II 4129.32 V19 3p 4P* 3d 4P 4 2 4131.97 0.175 2.104e - 02 O II 4132.80 V19 3p 4P* 3d 4P 2 4 4143.51 0.172 3.167e - 01 He I 4143.76 V53 2p IP* 6d ID 3 5 4152.93 0.170 3.037e - 02 O II 4153.30 V19 3p 4P* 3d 4P 4 6 4156.16 0.169 3.727e - 02 0 II 4156.53 V19 3p 4P* 3d 4P 6 4 4162.75 0.168 1.9396 - 02 [K V] 4163.12 * * 4168.85 0.167 3.9726 - 02 O II 4169.22 V19 3p 4P* 3d 4P 6 6 4185.25 0.163 3.596e - 02 O II 4185.45 V36 3p’ 2F* 3d’ 2G 6 8 4186.70 0.162 2.0436 - 01 C HI 4186.90 V18 4f IF* 5g IG 7 9 4189.59 0.162 2.2806 - 02 O II 4189.79 V36 3p’ 2F* 3d’ 2G 8 10 4195.60 0.160 3.2716 - 02 N HI 4195.76 V6 3s’ 2P* 3p’ 2D 2 4 4199.66 0.160 4.8136 - 01 He II 4199.83 4.11 4f+ 2F* l lg + 2G 32 * 0.160 * N HI 4200.10 V6 3s’ 2P* 3p’ 2D 4 6 4215.64 0.156 1.1146 - 02 ? 4215.64 * * 0.155 * Ne II 4219.37 V52a 3d 4D 4f 2\4i* 8 8 4219.32 0.155 1.822e - 02 Ne II 4219.74 V52a 3d 4D 4f 2j4i* 8 10 4227.32 0.153 1.0726 - 01 N II 4227.74 V33 3p ID 4s IP* 5 3 0.151 * N II 4236.91 V48a 3d 3D* 4f M3i 3 5 4237.50 0.151 1.7286 - 02 N II 4237.05 V48b 3d 3D* 4f lj4^ 5 7 4253.67 0.147 1.6776 - 02 O II 4254.00 V lO l 3d’ 2G 4P H5* 18 22 4266.96 0.144 6.2026 - 01 C II 4267.15 V6 3d 2D 4f 2F* 10 14 4275.31 0.142 3.1036 - 02 0 II 4275.55 V67a 3d 4D 4f F4* 8 10 0.142 * 0 II 4275.99 V67b 3d 4D 4f F3* 4 6 4276.16 0.142 2.7576 - 02 O II 4276.28 V67b 3d 4D 4f F3* 6 6 0.142 * 0 II 4276.75 V67b 3d 4D 4f F3* 6 8 4277.41 0.142 1.7156 - 02 O II 4277.43 V67c 3d 4D 4f F2* 2 4 0.141 * 0 II 4277.89 V67b 3d 4D 4f F3* 8 8 4282.78 0.140 8.6546 - 03 O II 4282.96 V67c 3d 4D 4f F2* 4 6 4283.55 0.140 9.5336 - 03 O II 4283.73 V67c 3d 4D 4f F2* 4 4 A.I. Galactic Planetary Nebulae 207

-^obs f(A) 7(A) Ion Ao Mult Lower term Upper term gl g2 4285.51 0.140 1.199e - 02 O II 4285.69 V78b 3d 2F 4f F3* 6 8 4288.41 0.139 5.451e - 03 O II 4288.82 V53c 3d 4P 4f Dl* 2 4 4291.21 0.138 1.615e - 02 O II 4291.25 V55 3d 4P 4f G3* 6 8 4292.17 0.138 8.559e - 03 O II 4292.21 V78c 3d 2F 4f F2* 6 6 4294.74 0.137 3.144e - 02 O II 4294.78 V53b 3d 4P 4f D2* 4 6 4303.64 0.135 5.824e - 02 O II 4303.61 V65a 3d 4D 4f G5* 8 10 4307.11 0.135 8.666e - 03 O II 4307.23 V53b 3d 4P 4f D2* 2 4 4313.14 0.133 8.5676 — 03 0 II 4313.44 V78a 3d 2F 4f F4* 8 10 4317.02 0.132 4.539e - 02 O II 4317.14 V2 3s 4P 3p 4P* 2 4 0.132 * O II 4317.70 V53a ** 4319.23 0.132 2.270e - 02 O II 4319.63 V2 3s 4P 3p 4P* 4 6 4325.54 0.130 3.107e - 02 O II 4325.76 V2 3s 4P 3p 4P* 2 2 4340.25 0.127 4.749e + 01 H 5 4340.47 H5 2p+ 2P* 5d+ 2D 8 50 4344.95 0.125 8.813e - 02 O II 4345.56 V2 3s 4P 3p 4P* 4 2 4349.30 0.125 7.1236 - 02 O II 4349.43 V2 3s 4P 3p 4P* 6 6 4353.91 0.124 9.7816 - 03 O II 4353.59 V76c 3d 2F 4f G3* 6 8 4357.57 0.123 1.8156 - 02 O II 4357.25 V63a * * 4362.99 0.121 1.3616 + 01 [O III] 4363.21 F2 2p2 ID 2p2 IS 5 1 4362.65 0.121 1.3576 + 01 [O III] 4363.21 F2 2p2 ID 2p2 IS 5 1 4366.76 0.120 9.0926 - 02 O II 4366.89 V2 3s 4P 3p 4P* 6 4 4369.73 0.120 3.0566 - 02 Ne II 4369.86 V56 3d 4F 4f Oj3i* 4 6 0.119 * O II 4371.62 V76b 3d 2F 4f G4* 8 10 4376.56 0.118 2.6066 - 02 ? 4376.56 * * 4378.98 0.118 1.5836 - 01 N III 4379.11 V18 4f 2F* 5g 2G 14 18 4385.83 0.116 2.7846 - 02 O II 4386.01 * * 4387.75 0.115 4.4266 - 01 He I 4387.93 V51 2p IP* 5d ID 3 5 4391.81 0.115 1.9526 - 02 Ne II 4391.99 V55e 3d 4F 4f 2i5i* 10 12 4409.03 0.110 1.8616 - 02 Ne II 4409.30 V55e 3d 4F 4f 2j5i* 8 10 4414.63 0.109 3.4866 - 02 O II 4414.90 V5 3s 2P 3p 2D* 4 6 4416.70 0.109 3.3116 - 02 O II 4416.97 V5 3s 2P 3p 2D* 2 4 4428.32 0.106 8.674e - 03 Ne II 4428.64 V60c 3d 2F 4f 1|3^* 6 8 4430.74 0.105 9.3546 - 03 Ne II 4430.94 V61a 3d 2D 4f 2\4i* 6 8 4437.35 0.103 5.2196 - 02 He I 4437.55 V50 2p IP* 5s IS 3 1 4442.12 0.103 4.8006 - 03 N II 4442.02 V55a 3d 3P* 4f 2\3i 3 5 4448.11 0.101 1.2146 - 02 O II 4448.19 V35 3p' 2F* 3d’ 2F 8 8 4453.50 0.100 1.7886 - 02 O II 4452.37 V5 3s 2P 3p 2D* 4 4 4458.18 0.099 1.826e - 02 Ne II 4457.24 V61d 3d 2D 4f 2\2i* 4 4 4465.79 0.097 2.5976 - 02 O II 4466.42 V86b 3d 2P 4f D2* 4 6 A.I. Galactic Planetary Nebulae 208

■^obs ( Â ) f(A) /(A) Ion Ao Mult Lower term Upper term gl g2 4471.29 0.095 4.124e + 00 He I 4471.50 V14 2p 3P* 4d 3D 9 15 4477.70 0.094 1.448e — 02 O II 4477.90 V88 3d 2P 4f G3* 4 6 4481.01 0.093 2.229e - 02 Mg II 4481.21 V4 3d 2D 4f 2F* 10 14 4488.05 0.091 1.004e - 02 O II 4488.20 V104 3d’ 2P 4P D2* 4 6 4489.34 0.091 5.510e —03 O II 4489.49 V86b ** 4491.08 0.091 2.351e - 02 O II 4491.23 V86a 3d 2P 4f D3* 4 6 4510.76 0.086 8.144e - 02 N III 4510.91 V3 3s’ 4P* 3p’ 4D 4 6 4514.71 0.085 2.228e - 02 N III 4514.86 V3 3s’ 4P* 3p’ 4D 6 8 4518.00 0.084 2.086e -02 N III 4518.15 V3 3s’ 4P* 3p’ 4D 2 2 4523.39 0.083 1.798e - 02 N III 4523.58 V3 3s’ 4P* 3p’ 4D 4 4 4534.39 0.080 2.639e - 02 N III 4534.58 V3 3s’ 4P* 3p’ 4D 6 6 4541.40 0.078 8.555e —01 He II 4541.59 4.9 4f+ 2F* 9g+ 2G 32 * 4544.61 0.078 3.962e - 02 N III 4544.80 ** 4562.43 0.073 5.068e - 02 Mg I] 4562.60 3s2 IS 3s3p 3P* 1 5 4570.93 0.071 8.480e - 02 Mg I] 4571.10 3s2 IS 3s3p 3P* 1 3 4590.85 0.066 2.878e - 02 O II 4590.97 V15 3s’ 2D 3p’ 2F* 6 8 4596.06 0.065 2.156e - 02 O II 4596.18 V15 3s’ 2D 3p’ 2F* 4 6 4601.96 0.064 4.988e - 03 N II 4601.48 V5 3s 3P* 3p 3P 3 5 4609.42 0.062 2.276e - 02 O II 4609.44 V92a ** 4610.18 0.062 2.339e - 02 O II 4610.20 V92c 3d 2D 4f F2* 4 6 4620.00 0.059 9.805e - 03 N II 4620.04 ** 4630.69 0.057 1.787e - 02 N II 4630.54 V5 3s 3P* 3p 3P 5 5 4633.96 0.056 8.341e - 01 N III 4634.14 V2 3p 2P* 3d 2D 2 4 4638.68 0.055 1.692e - 01 O II 4638.86 VI 3s 4P 3p 4D* 2 4 4640.46 0.054 1.575e + 00 N III 4640.64 V2 3p 2P* 3d 2D 4 6 4641.63 0.054 3.820e - 01 O II 4641.81 VI 3s 4P 3p 4D* 4 6 4647.27 0.053 3.279e - 01 C HI 4647.42 VI 3s 3S 3p 3P* 3 5 4648.98 0.052 2.231e - 01 O II 4649.13 VI 3s 4P 3p 4D* 6 8 4650.10 0.052 1.967e - 01 C III 4650.25 V I 3s 3S 3p 3P* 3 3 4650.69 0.052 9.785e - 02 O II 4650.84 VI 3s 4P 3p 4D* 2 2 4651.32 0.052 6.557e —02 C HI 4651.47 VI 3s 3S 3p 3P* 3 1 4658.01 0.050 7.735e - 02 C IV 4658.10 F3 3d6 5D 3d6 3F2 9 9 4661.54 0.049 8.697e -02 O II 4661.63 VI 3s 4P 3p 4D* 4 4 4669.14 0.047 1.193e - 02 O II 4669.27 V89b 3d 2D 4f D2* 4 6 4673.60 0.046 3.040e - 02 0 II 4673.73 VI 3s 4P 3p 4D* 4 2 4676.11 0.046 5.739e - 02 0 II 4676.24 VI 3s 4P 3p 4D* 6 6 4685.59 0.043 2.575e + 01 He II 4685.68 3.4 3d+ 2D 4f+ 2F* 18 32 0.041 * O II 4696.35 VI 3s 4P 3p 4D* 6 4 A.I. Galactic Planetary Nebulae 209

-^obs (Â) f(A) /(A) Ion Ao Mult Lower term Upper term gl g2 4699.30 0.040 5.854e — 03 O II 4699.22 V25 3p 2D* 3d 2F 4 6 4699.21 0.039 6.766e — 03 [Fe III] 4701.62 F3 3d6 5D 3d6 3F2 7 7 4711.25 0.037 4.893e -f 00 [Ar IV] 4711.37 FI 3p3 4S* 3p3 2D* 4 6 4713.05 0.036 5.675e - 01 He I 4713.17 vl2 2P 3p* 4S 3s 9 3 4724.05 0.034 4.930e - 02 [Ne IV] 4724.15 FI 2p3 2D* 2p3 2P* 4 4 4725.52 0.033 4.070e - 02 [Ne IV] 4725.62 FI 2p3 2D* 2p3 2P* 4 2 4740.12 0.030 4.529e + 00 [Ar IV] 4740.17 FI 3p3 4S* 3p3 2D* 4 4 4789.74 0.018 1.033e - 02 N II 4788.13 V20 * * 4797.67 0.016 1.648e - 02 ? 4797.67 * * 4802.52 0.015 1.913e - 02 C II 4802.23 41 2F* 8g2G 14 18 4861.28 0.000 l.OOOe 4-02 H 4 4861.33 H4 2p-b 2P* 4d+ 2D 8 32 4906.89 -0.011 1.861e - 02 O II 4906.83 V28 3p 4S* 3d 4P 4 4 4921.91 -0.015 1.026e -f 00 He I 4921.93 V48 2p IP* 4d ID 3 5 -0.016 * O II 4924.53 V28 3p 4S* 3d 4P 4 6 4931.28 -0.017 1.767e - 01 [O III] 4931.80 FI 2p2 3P 2p2 ID 1 5 4934.63 -0.018 1.186e - 02 ? 4934.63 * * 4958.37 -0.024 4.274e + 02 [O HI] 4958.91 FI 2p2 3P 2p2 ID 3 5 5006.30 -0.036 1.289e-b03 [O III] 5006.84 FI 2p2 3P 2p2 ID 5 5 5130.26 -0.066 5.232e - 02 7 5130.26 * * 5190.80 -0.081 1.215e - 01 [Ar III] 5191.82 * * 5303.80 -0.108 3.615e - 02 ? * ** 5344.65 -0.118 1.184e - 01 ? 5344.65 * * 5411.06 -0.134 1.914e -b 00 He II 5411.52 4.7 4f-t- 2F* 7g+ 2G 32 98 5517.36 -0.154 3.178e - 01 [Cl III] 5517.66 FI 2p3 4S* 2p3 2D* 4 6 5537.29 -0.158 2.779e - 01 [Cl HI] 5537.60 FI 2p3 4S* 2p3 2D* 4 4 5576.67 -0.164 3.565e - 02 [O I] 5577.34 F3 2p4 ID 2p4 IS 5 1 5591.52 -0.167 5.072e - 02 O HI 5592.37 V5 3s IP* 3p IP 3 3 5754.10 -0.195 6.346e - 02 [N II] 5754.60 F3 2p2 ID 2p2 IS 5 1 5801.00 -0.202 1.037e - 01 C IV 5801.51 VI 3s 2S 3p 2P* 2 4 5811.63 -0.204 8.744e - 02 C IV 5812.14 VI 3s 2S 3p 2P* 2 2 5875.26 -0.215 1.163e-b01 He I 5875.66 Vll 2p 3P* 3d 3D 9 15 5913.18 -0.221 1.953e - 02 He II 5913.26 5.26 5g+ 2G 26h+ 2H* 50 * 5931.76 -0.224 1.098e - 02 He II 5931.84 5.25 5g-t- 2G 25h4- 2H* 50 * 5952.86 -0.227 1.730e - 02 He II 5952.94 5.24 5g-b 2G 24h-f 2H* 50 * 5975.14 -0.231 1.903e - 02 7 5975.14 * * 6004.90 -0.236 4.596e - 02 He II 6004.73 5.22 5g-t- 2G 22h+ 2H* 50 * 6036.87 -0.241 4.883e - 02 He II 6036.70 5.21 5g+ 2G 21h+ 2H* 50 * 6074.27 -0.247 4.739e - 02 He II 6074.10 5.20 5g+ 2G 20h+ 2H* 50 * A.I. Galactic Planetary Nebulae 210

-^o b s ( Â ) f(A) /(A) Ion Ao Mult Lower term Upper term gl g2 6101.57 -0.251 2.718e - 01 [KIV] 6101.83 FI 3p4 3P 3d4 ID 5 5 6118.37 -0.254 5.987e - 02 He II 6118.20 5.19 5g-t- 2G 19h-H 2H* 50 * 6152.59 -0.259 1.864e - 02 C II 6151.43 V16.04 4d 2D 6f 2F* 10 14 6171.70 -0.262 4.093e - 02 He II 6170.60 5.18 5g-t- 2G 18h-|- 2H* 50 * 6234.91 -0.272 7.091e - 02 He II 6233.80 5.17 5g+ 2G 17h4- 2H* 50 * 6299.99 -0.282 4.674e - 02 [OI] 6300.34 FI 2p4 3P 2p4 ID 5 5 6311.93 -0.283 6.318e - 01 [S HI] 6312.10 F3 2p2 ID 2p2 IS 5 1 6407.45 -0.297 1.107e - 01 He II 6406.30 5.15 5g+ 2G 15h+ 2H* 50 * 6435.37 -0.302 4.134e - 02 7 6435.37 ** 6461.46 -0.306 6.329e - 02 C II 6461.95 4f 2F* ^ :2 G 14 18 6527.18 -0.315 1.457e - 01 He II 6527.11 5.14 5g+ 2G 14h-t- 2H* 50 * 6548.19 -0.318 9.852e - 01 [N II] 6548.10 FI 2p2 3P 2p2 ID 3 5 6562.86 -0.320 2.887e -f 02 H 3 6562.77 H3 2p+ 2P* 3d+ 2D 8 18 6583.59 -0.323 2.631e + 00 [N II] 6583.50 FI 2p2 3P 2p2 ID 5 5 6678.41 -0.336 3.190e-h00 He I 6678.16 V46 2p IP* 3d ID 3 5 6716.69 -0.342 2.780e - 01 [S II] 6716.44 F2 2p3 4S* 2p3 2D* 4 6 6731.07 -0.344 3.922e - 01 [S II] 6730.82 F2 2p3 4S* 2p3 2D* 4 4 6780.76 -0.350 2.365e - 02 7 6780.76 * * 6794.47 -0.352 6.201e - 02 [KIV] 6795.00 * * 6862.64 -0.361 2.553e - 01 7 6862.64 * * 6891.16 -0.365 1.535e - 01 He II 6890.88 7 * * 7006.03 -0.380 8.357e - 02 [ArV] 7005.74 * * 7065.03 -0.387 2.959e -f 00 He I 7065.25 VIO 2p 3P* 3s 3S 9 3 7135.58 -0.396 6.999e -H 00 [Ar HI] 7135.80 FI 3p4 3P 3p4 ID 5 5 7170.23 -0.401 1.073e - 01 [Ar IV] 7170.62 F2 3p3 2D* 3p3 2P* 4 4 7177.11 -0.401 1.753e - 01 He II 7177.50 5.11 5g-t- 2G l lh + 2H* 50 * 7236.27 -0.408 2.634e - 01 C II 7236.42 V3 3p 2P* 3d 2D 4 6 7262.61 -0.412 1 .2 4 6 e - 01 [Ar IV] 7262.76 F2 3p3 2D* 3p3 2P* 4 2 7281.20 -0.414 4.233e- 01 He I 7281.35 V45 2p IP* 3s IS 3 1 7318.85 -0.418 5.045e - 01 [O II] 7318.92 F2 2p3 2D* 2p3 2P* 6 2 7329.59 -0.420 4.005e - 01 [O II] 7329.67 F2 2p3 2D* 2p3 2P* 4 2

c(H/3)(Bal) = 0.437, c(H/3)(rad) = 0.398, c(H/3)^i640 = 0.489

Table A.5: NGC3918 linelist.

-^obs (-^) f(A) /(A) Ion Ao Mult Lower term Upper term gi g2

1240 0.257 4.140e 4- 01 N V 1240 1401 0.257 4.890e -f 01 O IV] 1401 1486 0.257 4.640e -f 01 N IV] 1486 1549 0.257 4,5766 4"" 02 C IV 1549 A.I. Galactic Planetary Nebulae 211

-^obs {N) f(A) 7(A) Ion Ao Mult Lower term Upper term gl g2 1601 0.257 5.900e + 00 [Ne IV] 1601 1640 0.257 2.836e + 02 He II 1640 1663 0.257 3.150e + 00 O III] 1663 1751 0.257 2.670e + 01 N III] 1751 1908 0.257 4.925e + 02 C III] 1908 2326 0.257 3.124e + 01 C II] 2326 3726.46 0.257 9.452e + 01 [Oil] 3726.03 FI 2p3 4S* 2p3 2D* 4 4 0.256 * [O II] 3728.82 FI 2p3 4S* 2p3 2D* 4 6 3749.69 0.253 3.700e + 00 H 12 3750.15 H12 2p+ 2P* 12d+ 2D 8 * 3759.41 0.251 5.110e + 00 O III 3759.87 V2 3s 3P* 3p 3D 5 7 3770.07 0.249 4.072e + 00 H 11 3770.63 H ll 2p+ 2P* lld+ 2D 8 * 3797.23 0.244 5.624e + 00 H 10 3797.90 HIO 2p+ 2P* 10d+ 2D 8 * 3819.14 0.240 7.001e - 01 He I 3819.62 V22 2p 3P* 6d 3D 9 15 3834.62 0.237 7.730e + 00 H 9 3835.39 H9 2p+ 2P* 9d+ 2D 8 * 3867.98 0.230 1.355e + 02 [Ne III] 3868.75 FI 2p4 3P 2p4 ID 5 5 3888.09 0.226 1.982e + 01 He I 3888.65 V2 2s 3S 3p 3P* 3 9 3922.02 0.220 3.118e - 01 He II 3923.48 4.15 4f+ 2F* 15g+ 2G 32 * 3967.20 0.211 5.759c 4” 01 [Ne HI] 3967.46 FI 2p4 3P 2p4 ID 3 5 4008.79 0.202 1.348e - 01 He I 4009.26 V55 2p IP* 7d ID 3 5 4025.64 0.198 2.270e + 00 He I 4026.21 V18 2p 3P* 5d 3D 9 15 4040.83 0.195 1.338e - 02 N II 4041.31 V39b 3d 3F* 4f 2 < 5 > 9 11 4059.65 0.191 3.755e - 02 O II 4060.60 2s2 2F 2s2 2 [4] 8 10 4067.99 0.189 1.976C + 00 [S II] 4068.60 FI 2p3 4S* 2p3 2P* 4 4 4069.17 0.189 5.330e - 01 O II 4069.62 * * 4071.55 0.188 2.378e - 01 O II 4072.16 VIO 3p 4D* 3d 4F 6 8 4075.25 0.187 6.576e - 02 O II 4075.86 VIO ** 4075.74 0.187 7.170c - 01 [S II] 4076.35 FI 2p3 4S* 2p3 2P* 2 4 4083.46 0.186 2.771c - 02 O II 4083.90 V48b 3d 4F 4f G4* 6 8 4084.67 0.185 1.939c - 02 O II 4085.11 VIO 3p 4D* 3d 4F 6 6 4087.12 0.185 2.390c - 02 0 II 4087.15 V48c 3d 4F 4f G3* 4 6 4088.85 0.184 5.603c — 02 O II 4089.29 V48a 3d 4F 4f G5* 10 12 4096.84 0.183 2.163c+ 00 N III 4097.33 VI 3s 2S 3p 2P* 2 4 4101.20 0.182 2.798c+ 01 H 6 4101.74 H6 2p+ 2P* 6d+ 2D 8 72 4118.62 0.178 6.098c - 02 0 II 4119.22 V20 3p 4P* 3d 4D 6 8 0.178 * O II 4120.28 V20 3p 4P* 3d 4D 6 6 4120.24 0.177 2.055c - 01 He I 4120.84 V16 2p 3P* 5s 3S 9 3 4121.91 0.177 9.197c - 02 O II 4121.46 V19 3p 4P* 3d 4P 2 2 4128.18 0.176 2.745c - 02 ? 4128.18 * * A.I. Galactic Planetary Nebulae 212

-^obs (À) f(A) /(A) Ion Ao Mult Lower term Upper term gl g2 0.176 * O II 4129.32 V19 3p 4P* 3d 4P 4 2 4131.76 0.175 1.800e - 02 O II 4132.80 V19 3p 4P* 3d 4P 2 4 4143.18 0.172 3.330e - 01 He I 4143.76 V53 2p IP* 6d ID 3 5 4152.63 0.170 2.600e - 02 O II 4153.30 V19 3p 4P* 3d 4P 4 6 4155.84 0.169 3.701e - 02 O II 4156.53 V19 3p 4P* 3d 4P 6 4 4168.17 0.167 3.098e - 02 O II 4169.22 V19 3p 4P* 3d 4P 6 6 4184.89 0.163 3.747e - 02 O II 4185.45 V36 3p’ 2F* 3d’ 2G 6 8 4186.34 0.162 1.519e - 01 C III 4186.90 V18 4f IF* 5g IG 7 9 4189.23 0.162 3.848e — 02 O II 4189.79 V36 3p’ 2F* 3d’ 2G 8 10 4195.17 0.160 1.998e - 02 N III 4195.76 V6 3s’ 2P* 3p’ 2D 2 4 4199.32 0.160 8.101e - 01 He II 4199.83 4.11 4f+ 2F* l lg + 2G 32 * 4219.20 0.155 7.947e - 03 Ne II 4219.74 V52a 3d 4D 4f 2<4>* 8 10 4226.95 0.153 2.566e —01 N II 4227.74 V33 3p ID 4s IP* 5 3 4237.30 0.151 1.630e - 02 N II 4237.05 V48b 3d 3D* 4f 1<4> 5 7 4241.13 0.150 1.628e - 02 N II 4241.78 V48a 3d 3D* 4f 1< 3> 5 7 4253.53 0.147 1.275e - 02 O II 4254.00 VlOl 3d’ 2G 4f’ H5* 18 22 4266.62 0.144 4.971e - 01 C II 4267.15 V6 3d 2D 4f 2F* 10 14 4275.21 0.142 3.0346 —02 O II 4275.55 V67a 3d 4D 4f F4* 8 10 0.142 * O II 4275.99 V67b 3d 4D 4f F3* 4 6 4276.19 0.142 1.372e - 02 O II 4276.75 V67b 3d 4D 4f F3* * * 4276.87 0.142 9.295e - 03 O II 4277.43 V67c 3d 4D 4f F2* * * 4291.09 0.138 2.437e - 02 O II 4291.25 V55 3d 4P 4f G3* 6 8 4294.62 0.137 2.538e - 02 O II 4294.78 V53b 3d 4P 4f D2* 4 6 4303.35 0.135 2.292e - 02 O II 4303.83 V53a 3d 4P 4f D3* * * 4316.71 0.132 4.170e - 02 O II 4317.14 V2 3s 4P 3p 4P* 2 4 4319.20 0.132 1.701e - 02 O II 4319.63 V2 3s 4P 3p 4P* 4 6 4325.33 0.130 1.414e - 02 O II 4325.76 V2 3s 4P 3p 4P* 2 2 4339.89 0.127 4.929e + 01 H 5 4340.47 H5 2p+ 2P* 5d+ 2D 8 50 0.125 * 0 II 4345.56 V2 3s 4P 3p 4P* 4 2 4349.00 0.125 6.383e - 02 O II 4349.43 V2 3s 4P 3p 4P* 6 6 4362.64 0.121 2.1556 + 01 [O HI] 4363.21 F2 2p2 ID 2p2 IS 5 1 4376.53 0.118 3.1536 - 02 ? 4376.53 * * 4378.68 0.118 1.4636 - 01 N HI 4379.11 V18 4f 2F* 5g 2G 14 18 4387.39 0.115 4.2346 - 01 He I 4387.93 V51 2p IP* 5d ID 3 5 4391.45 0.115 2.1546 - 02 Ne II 4391.99 V55e 3d 4F 4f 2<5>* 10 12 4408.70 0.110 2.4926 - 02 Ne II 4409.30 V55e 3d 4F 4f 2<5>* 8 10 4414.30 0.109 5.1786 - 02 O II 4414.90 V5 3s 2P 3p 2D* 4 6 4416.37 0.109 2.9006 - 02 O II 4416.97 V5 3s 2P 3p 2D* 2 4 A.I. Galactic Planetary Nebulae 213

-^obs {N) f(A) 7(A) Ion Ao Mult Lower term Upper term gl g2 4428.11 0.106 9.465e — 03 Ne II 4428.64 V60c 3d 2F 4f 1[3]* 6 8 4434.28 0.105 3.000e - 02 N II 4433.48 * * 4437.02 0.103 5.083e - 02 He I 4437.55 V50 2p IP* 5s IS 3 1 4448.19 0.101 1.760e - 02 O II 4448.19 V35 3p’ 2F* 3d’ 2F 8 8 4452.37 0.100 6.525e - 02 O II 4452.37 V5 3s 2P 3p 2D* 4 4 4458.10 0.099 3.756e - 02 Ne II 4457.05 * * 4465.13 0.097 2.425e - 02 O II 4466.42 * * 4470.93 0.095 3.885e + 00 He I 4471.50 V14 2p 3P* 4d 3D 9 15 4480.62 0.093 1.208e - 02 Mg II 4481.21 V4 3d 2D 4f 2F* 10 14 4490.55 0.091 1.534e - 02 C II 4491.07 4f 2F* 9g 2G 14 18 0.091 * O II 4491.23 V86a 3d 2P 4f D3* 4 6 4510.36 0.086 8.656e - 02 N HI 4510.91 V3 3s’ 4P* 3p’ 4D 4 6 4514.31 0.085 2.429e - 02 N HI 4514.86 V3 3s’ 4P* 3p’ 4D 6 8 4517.60 0.084 2.351e - 02 N III 4518.15 V3 3s’ 4P* 3p’ 4D 2 2 4523.03 0.083 1.848e - 02 N HI 4523.58 V3 3s’ 4P* 3p’ 4D 4 4 4529.86 0.081 1.476e - 02 N II 4530.41 V58b 3d IF* 4f 2<5> 7 9 4534.03 0.080 3.241e - 01 N HI 4534.58 V3 3s’ 4P* 3p’ 4D 6 6 4541.03 0.078 1.515e + 00 He II 4541.59 4.9 4f+ 2F* 9g+ 2G 32 * 4544.24 0.078 2.952e - 02 N HI 4544.80 * * 4552.24 0.076 1.403e - 02 N II 4552.53 V58a 3d IF* 4f 2<4> 7 9 4562.01 0.073 1.335e - 01 Mg I] 4562.60 3s2 IS 3s3p 3P* 1 5 4570.51 0.071 4.749e - 01 Mg I] 4571.10 3s2 IS 3s3p 3P* 1 3 4590.54 0.066 4.080e - 02 O II 4590.97 V15 3s’ 2D 3p’ 2F* 6 8 4595.75 0.065 2.754e - 02 O II 4596.18 V15 3s’ 2D 3p’ 2F* 4 6 4606.04 0.062 6.603e - 02 N II 4607.16 V5 3s 3P* 3p 3P 1 3 0.062 * O II 4609.44 V92a 3d 2D 4f F4* 6 8 4609.09 0.062 3.778e - 02 O II 4610.20 V92c 3d 2D 4f F2* 4 6 4619.68 0.059 1.698e - 02 ? 4619.68 * * 4624.75 0.058 4.665e - 02 [Ar V] 4624.54 3s2 ID 3s2 IS 5 1 4629.97 0.057 7.283e - 02 N II 4630.54 V5 3s 3P* 3p 3P 5 5 4631.71 0.056 2.167e - 01 7 4631.71 * * 4633.58 0.056 1.532e + 00 N HI 4634.14 V2 3p 2P* 3d 2D 2 4 4638.30 0.055 2.200e - 01 O II 4638.86 VI 3s 4P 3p 4D* 2 4 4640.08 0.054 3.122e + 00 N HI 4640.64 V2 3p 2P* 3d 2D 4 6 4641.25 0.054 5.232e - 01 O II 4641.81 VI 3s 4P 3p 4D* 4 6 4646.84 0.053 2.198e - 01 C HI 4647.42 VI 3s 3S 3p 3P* 3 5 4648.55 0.052 2.108e - 01 O II 4649.13 VI 3s 4P 3p 4D* 6 8 4649.67 0.052 1.397e - 01 C HI 4650.25 VI 3s 3S 3p 3P* 3 3 A.I. Galactic Planetary Nebulae 214

-^obs (^) f(A) /(A ) Ion Ao Mult Lower term Upper term gl g2 4650.26 0.052 4.146e - 02 O II 4650.84 VI 3s 4P 3p 4D* 2 2 4650.89 0.052 5.807e - 02 C III 4651.47 VI 3s 3S 3p 3P* 3 1 4657.74 0.050 3.321e - 01 C IV 4658.10 F3 3d6 5D 3d6 3F2 9 9 4661.07 0.049 7.432e - 02 O II 4661.63 VI 3s 4P 3p 4D* 4 4 4665.25 0.048 1.417e - 02 C III 4665.90 ** 4668.90 0.047 1.219e - 02 O II 4669.27 ** 4673.36 0.046 3.101e - 02 O II 4673.73 VI 3s 4P 3p 4D* 4 2 4675.87 0.046 6.560e - 02 O II 4676.24 VI 3s 4P 3p 4D* 6 6 4679.81 0.045 6.694e - 02 7 4679.81 * * 4685.21 0.043 4.165e 4- 01 He II 4685.68 3.4 3d-t- 2D Ai+ 2F* 18 32 0.040 * O II 4699.22 V25 3p 2D* 3d 2F 4 6 4701.14 0.039 1.768e - 02 [Fe III] 4701.62 F3 3d6 5D 3d6 3F2 7 7 4710.86 0.037 6.165e -f- 00 [Ar IV] 4711.37 FI 3p3 4S* 3p3 2D* 4 6 4712.66 0.036 8.714e - 01 He I 4713.17 V12 2P 3p* 4S 3s 9 3 4723.60 0.034 5.794e - 01 [Ne IV] 4724.15 FI 2p3 2D* 2p3 2P* 4 4 4725.07 0.033 4.995e - 01 [Ne IV] 4725.62 FI 2p3 2D* 2p3 2P* 4 2 4739.73 0.030 7.626e -f 00 [Ar IV] 4740.17 FI 3p3 4S* 3p3 2D* 4 4 4754.27 0.026 8.419e - 03 [Fe HI] 4754.83 3d6 5D 3d6 3F2 7 9 4769.28 0.023 2.149e - 02 ? 4769.28 ** 4802.99 0.015 2.638e - 02 C II 4802.23 4f 2F* 8 g 2 G 14 18 4860.89 0.000 l.OOOe-t-02 H 4 4861.33 H4 2p+ 2P* 4d-t- 2D 8 32 4880.60 -0.005 3.185e - 02 [Fe HI] 4881.11 F2 3d6 5D 3d6 3H 9 9 4905.88 -0.011 9.690e - 03 O II 4906.83 V28 3p 4S* 3d 4P 4 4 4921.52 -0.015 9.359e - 01 He I 4921.93 V48 2p IP* 4d ID 3 5 -0.016 * O II 4924.53 V28 3p 4S* 3d 4P 4 6 4930.85 -0.017 2.064e - 01 [O HI] 4931.80 FI 2p2 3P 2p2 ID 1 5 4958.00 -0.024 5.489e -f 02 [O III] 4958.91 FI 2p2 3P 2p2 ID 3 5 5005.93 -0.036 1.657e -f 03 [O HI] 5006.84 FI 2p2 3P 2p2 ID 5 5 5189.94 -0.081 2.415e - 01 [Ar HI] 5191.82 * * 5197.96 -0.083 6.284e - 01 [N I] 5199.84 * * 5410.60 -0.134 3.042e -f 00 He II 5411.52 4.7 4f+ 2F* 7g+ 2G 32 98 5516.79 -0.154 5.481e - 01 [Cl III] 5517.66 FI 2p3 4S* 2p3 2D* 4 6 5536.98 -0.158 7.851e - 01 [Cl HI] 5537.60 FI 2p3 4S* 2p3 2D* 4 4 5591.58 -0.167 6.425e - 02 O HI 5592.37 V5 3s IP* 3p IP 3 3 5753.69 -0.195 1.735e -f 00 [N II] 5754.60 F3 2p2 ID 2p2 IS 5 1 5800.81 -0.202 9.789e - 02 C IV 5801.51 VI 3s 2S 3p 2P* 2 4 5874.81 -0.215 1.081e-k01 He I 5875.66 Vll 2p 3P* 3d 3D 9 15 6086.35 -0.249 l.OOOe - 01 [Fe VII] 6086.90 ** A.I. Galactic Planetary Nebulae 215

-^obs (Â) f(A) /(A) Ion Ao Mult Lower term Upper term gl g2 6101.17 -0.251 3.496e - 01 [K IV] 6101.83 FI 3p4 3P 3d4 ID 5 5 6170.07 -0.262 7.300e - 02 He II 6170.60 5.18 5g+ 2G 18h-b 2H 5 * 6233.27 -0.272 9.131e - 02 He II 6233.80 5.17 5g+ 2G 17h-b 2H 5 * 6299.76 -0.282 5.243e -f 00 [OI] 6300.34 FI 2p4 3? 2p4 ID 5 5 6311.44 -0.283 2.271e -1- 00 [S III] 6312.10 F3 2p2 ID 2p2 IS 5 1 6363.34 -0.291 1.776e-h00 [O I] 6363.78 FI 2p4 3P 2p4 ID 3 5 6406.11 -0.297 1.705e - 01 He II 6406.30 5.15 5g+ 2G 15h+ 2H 5 * 6434.64 -0.302 7.899e - 01 [ArV] 6435.10 3s2 3P 3s2 ID 3 5 6527.11 -0.315 1.530e - 01 He II 6527.11 5.14 5g+ 2G 14h+ 2H 5 * 6547.62 -0.318 2.647e + 01 [N II] 6548.10 FI 2p2 3P 2p2 ID 3 5 6562.27 -0.320 3.040e -f 02 H 3 6562.77 H3 2p+ 2P* 3d+ 2D 8 18 6583.00 -0.323 8.366e -b 01 [N II] 6583.50 FI 2p2 3P 2p2 ID 5 5 6678.00 -0.336 2.873e -b 00 He I 6678.16 V46 2p IP* 3d ID 3 5 6716.04 -0.342 2.893e -b 00 [S II] 6716.44 F2 2p3 4S* 2p3 2D* 4 6 6730.41 -0.344 5.039e -b 00 [S II] 6730.82 F2 2p3 4S* 2p3 2D* 4 4 7005.19 -0.380 1.508e -b 00 [Ar V] 7005.67 3s2 3P 3s2 ID 5 5 7064.61 -0.387 3.853c "b 00 He I 7065.25 VIO 2p 3P* 3s 3S 9 3 7134.99 -0.396 1.742e -b 01 [Ar HI] 7135.80 FI 3p4 3P 3p4 ID 5 5 7169.84 -0.401 3.087c - 01 [Ar IV] 7170.62 F2 3p3 2D* 3p3 2P* 4 4 7176.65 -0.401 3.559c - 01 He II 7177.50 5.11 5g-b 2G l lh + 2H 5 * 7236.26 -0.408 3.646c - 01 C II 7236.42 V3 3p 2P* 3d 2D 4 6 7262.28 -0.412 2,049c - 01 [Ar IV] 7262.76 F2 3p3 2D* 3p3 2P* 4 2 7280.48 -0.414 4.484c - 01 He I 7281.35 V45 2p IP* 3s IS 3 1 7318.44 -0.418 3.932c -b 00 [O II] 7318.92 F2 2p3 2D* 2p3 2P* 6 2 7329.28 -0.420 3.474c -b 00 [O II] 7329.67 F2 2p3 2D* 2p3 2P* 4 2 7418.96 -0.430 1.155c -b 00 ? 7418.96 ** 4162.78 0.168 6.749c - 02 [K V] ***

c(H/3)(Bal) = 0.700, c(H/9)(rad) = 0.569,, c(H/3)^7330/2470 = 0.586

T able A .6: N G C 5315 linelist.

-^obs (^) f(A) /(A) Ion Aq Mult Lower term Upper term gi g2

1401 0.257 6.790e - 01 O IV] 1401 1486 0.257 2.580c -b 00 N IV] 1486 1574 0.257 7.410c -b 00 [Ne V] 1574 1601 0.257 2.410c -b 00 [Ne IV] 1601 1663 0.257 4.420c + 00 O III] 1663 1751 0.257 5.260c -b 00 N III] 1663 1908 0.257 3.188c-b 01 C III] 1908 A.I. Galactic Planetary Nebulae 216

-^obs (Â) f(A) /(A) Ion Ao Mult Lower term Upper term gl g2 2424 0.257 1.920e + 00 [Ne IV] 2424 2471 0.257 8.930e + 00 [O II] 2471 0.257 * [O II] 3726.03 FI 2p3 4S* 2p3 2D* 4 4 0.256 * [O II] 3728.82 FI 2p3 4S* 2p3 2D* 4 6 3749.33 0.253 2.825e + 00 H 12 3750.15 H12 2p+ 2P* 12d+ 2D 8 * 3769.57 0.249 4.256e + 00 H 11 3770.63 H ll 2p+ 2P* lld+ 2D 8 * 3797.03 0.244 5.141e + 00 H 10 3797.90 HIO 2p+ 2P* 10d+ 2D 8 * 3817.43 0.240 1.516e + 00 He I 3819.62 V22 2p 3P* 6d 3D 9 15 3834.04 0.237 9.025e + 00 H 9 3835.39 H9 2p+ 2P* 9d+ 2D 8 * 3867.55 0.230 6.862e + 01 [Ne III] 3868.75 FI 2p4 3P 2p4 ID 5 5 3887.64 0.226 1.822e + 01 He I 3888.65 V2 2s 3S 3p 3P* 3 9 3967.15 0.211 3.651e + 01 [Ne HI] 3967.46 FI 2p4 3P 2p4 ID 3 5 4008.44 0.202 2.662e - 01 He I 4009.26 V55 2p IP* 7d ID 3 5 4025.38 0.198 3.117e + 00 He I 4026.21 V18 2p 3P* 5d 3D 9 15 4040.42 0.195 5.729e - 02 N II 4041.31 V39b 3d 3F* 4f 2|5^ 9 11 4067.63 0.189 5.699e + 00 [S II] 4068.60 FI 2p3 4S* 2p3 2P* 4 4 4068.79 0.189 4.263e - 01 O II 4069.62 VIO * * 4071.19 0.188 1.879e - 01 0 II 4072.16 VIO 3p 4D* 3d 4F 6 8 4074.89 0.187 1.412e - 0 1 O II 4075.86 VIO ** 4075.38 0.187 2.046e + 00 [S II] 4076.35 FI 2p3 4S* 2p3 2P* 2 4 4082.97 0.186 2.411e - 02 O II 4083.90 V48b 3d 4F 4f G4* 6 8 4084.18 0.185 4.371e - 02 O II 4085.11 VIO 3p 4D* 3d 4F 6 6 4086.22 0.185 3.185e - 02 O II 4087.15 V48c 3d 4F 4f G3* 4 6 4088.36 0.184 8.987e - 02 O II 4089.29 V48a 3d 4F 4f G5* 10 12 4096.68 0.183 3.088e - 01 N HI 4097.33 VI 3s 2S 3p 2P* 2 4 4100.80 0.182 2.650e + 01 H 6 4101.74 H6 2p+ 2P* 6d+ 2D 8 72 4109.04 0.180 1.417e - 02 O II 4110.78 V20 3p 4P* 3d 4D 4 2 4119.70 0.178 3.309e - 01 O II 4119.22 V20 3p 4P* 3d 4D 6 8 4121.32 0.177 5.202e - 02 He I 4120.84 V16 2p 3P* 5s 3S 9 3 0.177 * O II 4121.46 V19 3p 4P* 3d 4P 2 2 4131.68 0.175 7.292e - 02 O II 4132.80 V19 3p 4P* 3d 4P 2 4 4142.89 0.172 5.016e - 01 He I 4143.76 V53 2p IP* 6d ID 3 5 4152.27 0.170 9.108e - 02 O II 4153.30 V19 3p 4P* 3d 4P 4 6 4155.50 0.169 8.363e - 02 0 II 4156.53 V19 3p 4P* 3d 4P 6 4 4167.97 0.167 9.273e - 02 0 II 4169.22 V19 3p 4P* 3d 4P 6 6 4184.47 0.163 5.256e - 02 O II 4185.45 V36 3p’ 2F* 3d’ 2G 6 8 4185.92 0.162 1.071e - 02 C HI 4186.90 V18 4f IF* 5g IG 7 9 4188.81 0.162 6.709e - 02 O II 4189.79 V36 3p’ 2F* 3d’ 2G 8 10 A.I. Galactic Planetary Nebulae 217

■^obs (^ ) f(A) 7(A) Ion Ao Mult Lower term Upper term gl g2 4218.93 0.155 1.412e - 02 Ne II 4219.74 V52a 3d 4D 4f 2\Ai* 8 10

4236.04 0.151 1.777e - 02 N II 4237.05 V48b 3d 3D* 4f l\H 5 7 4240.99 0.150 4.211e - 02 N II 4241.24 V48a 3d 3D* 4f liS i 5 5 4252.87 0.147 3.169e - 02 O II 4254.00 VlOl 3d’ 2G 4P H5* 18 22 4266.15 0.144 7.061e - 01 C II 4267.15 V6 3d 2D 4f 2F* 10 14 4274.54 0.142 1.053e - 01 O II 4275.55 V67a 3d 4D 4f F4* 8 10 0.142 * O II 4275.99 V67b 3d 4D 41 F3* 4 6 0.142 * O II 4276.28 V67b 3d 4D 4f F3* 6 6 0.142 * O II 4276.75 V67b 3d 4D 4f F3* 6 8 0.142 * O II 4277.43 V67c 3d 4D 4f F2* 2 4 0.141 * O II 4277.89 V67b 3d 4D 4f F3* 8 8 0.135 * O II 4303.61 V65a 3d 4D 4f G5* 8 10 4302.82 0.135 3.307e - 02 O II 4303.82 V53a 3d 4P 4f D3* 6 8 4316.19 0.132 6.940e - 02 O II 4317.14 V2 3s 4P 3p 4P* 2 4 0.132 * O II 4317.70 V53a 3d 4P 4f D3* 4 6 4318.68 0.132 4.912e - 02 O II 4319.63 V2 3s 4P 3p 4P* 4 6 4324.96 0.130 2.688e - 02 O II 4325.76 V2 3s 4P 3p 4P* 2 2 4339.50 0.127 4.854e + 01 H 5 4340.47 H5 2p+ 2P* 5d+ 2D 8 50 4344.58 0.125 6.347e - 02 O II 4345.56 V2 3s 4P 3p 4P* 4 2 4348.45 0.125 7.439e - 02 O II 4349.43 V2 3s 4P 3p 4P* 6 6 4362.26 0.121 3.913e + 00 [O III] 4363.21 F2 2p2 ID 2p2 IS 5 1 4366.73 0.120 5.887e - 02 0 II 4366.89 V2 3s 4P 3p 4P* 6 4 4386.98 0.115 7.589e - 01 He I 4387.93 V51 2p IP* 5d ID 3 5 4390.89 0.115 1.684e - 02 Ne II 4391.99 V55e 3d 4F 4f 2\5i* 10 12 4396.89 0.113 1.920e - 02 Ne II 4397.99 V57b 3d 4F 4f l-,4^* 6 8 4408.71 0.110 1.195e - 02 Ne II 4409.30 V55e 3d 4F 4f 2j5i* 8 10 4412.18 0.109 2.450e - 02 Ne II 4413.22 V65 ** 4413.86 0.109 7.206e - 02 O II 4414.90 V5 3s 2P 3p 2D* 4 6 4415.93 0.109 6.200e - 02 O II 4416.97 V5 3s 2P 3p 2D* 2 4 4436.61 0.103 6.380e - 02 He I 4437.55 V50 2p IP* 5s IS 3 1 4324.76 0.100 1.455e - 02 O II 4452.37 V5 * * 4470.52 0.095 6.495e + 00 He I 4471.50 V14 2p 3P* 4d 3D 9 15 4480.30 0.093 2.091e - 02 Mg II 4481.21 V4 3d 2D 4f 2F* 10 14 4490.40 0.091 1.620e - 02 C II 4491.07 4f 2F* 9g 2G 14 18 0.091 * O II 4491.23 V86a 3d 2P 41 D3* 4 6 4529.05 0.081 5.474e - 02 N II 4530.41 V58b 3d IF* 4f 2,57 7 9 4551.59 0.076 2.694e - 02 N II 4552.53 V58a 3d IF* 4f 2\Ai 7 9

0.073 * Mg I] 4562.60 3s2 IS 3s3p 3P* 1 5 A.I. Galactic Planetary Nebulae 218

-^obs (Â) f(A) /(A) Ion Ao Mult Lower term Upper term gl g2 4570.07 0.071 2.302e - 01 Mg I] 4571.10 3s2 IS 3s3p 3P* 1 3 4590.06 0.066 7.937e - 02 O II 4590.97 V15 3s’ 2D 3p’ 2F* 6 8 4595.27 0.065 5.843e —02 O II 4596.18 V15 3s’ 2D 3p’ 2F* 4 6 4600.59 0.064 5.659e —02 N II 4601.48 V5 3s 3P* 3p 3P 3 5 4606.27 0.062 5.321c —02 N II 4607.16 V5 3s 3P* 3p 3P 1 3 4608.54 0.062 4.763e - 02 O II 4609.44 V92a 3d 2D 4f F4* 6 8 0.062 * O II 4610.20 V92c 3d 2D 4f F2* 4 6 4624.96 0.058 3.590e - 02 [ArV] 4624.54 3s2 ID 3s2 18 5 1 4629.65 0.057 1.208c - 01 N II 4630.54 V5 3s 3P’*= 3p 3P 5 5 4633.32 0.056 1.245e - 01 N III 4634.14 V2 3p 2P* 3d 2D 2 4 4638.04 0.055 1.740c - 01 O II 4638.86 VI 3s 4P 3p 4D* 2 4 4639.82 0.054 2.591c - 01 N III 4640.64 V2 3p 2P’* 3d 2D 4 6 4640.99 0.054 2.616c - 01 O II 4641.81 VI 3s 4P 3p 4D* 4 6 4648.17 0.052 3.574c - 01 O II 4649.13 VI 3s 4P 3p 4D* 6 8 4649.88 0.052 9.002c - 02 0 II 4650.84 VI 3s 4P 3p 4D* 2 2 4657.05 0.050 3.156c —01 [Fe III] 4658.10 ** 4660.58 0.049 8.969c - 02 O II 4661.63 VI 3s 4P 3p 4D* 4 4 4672.77 0.046 1.367c - 02 0 II 4673.73 VI 3s 4P 3p 4D* 4 2 4675.28 0.046 7.894c - 02 O II 4676.24 VI 3s 4P 3p 4D* 4 2 4677.18 0.045 3.121c - 02 N II 4678.14 V61b 3d IP* 4f 2j2i 3 5 4684.78 0.043 5.040c —02 He II 4685.68 3.4 3d-f- 2D 4f+ 2F* 18 32 0.041 * O II 4696.35 VI 3s 4P 3p 4D* 6 4 4698.92 0.040 2.497c - 02 O II 4699.22 V25 3p 2D* 3d 2F 4 6 4700.66 0.039 7.159c - 02 [Fe III] 4701.62 F3 3d6 5D 3d6 3F2 7 7 4712.20 0.037 8.385c —01 [Ar IV] 4711.37 FI 3p3 4S* 3p3 2D* 4 6 0.036 * He I 4713.17 V12 2P 3p* 48 3s 9 3 4739.46 0.030 1.889c - 01 [Ar IV] 4740.17 FI 3p3 4S* 3p3 2D* 4 4 4753.58 0.026 4.694c —03 [Fe III] 4754.83 3d6 5D 3d6 3F2 7 9 4802.07 0.015 3.071c - 02 C II 4802.23 4f 2F* 8g 2 G 14 18 4814.15 0.011 1.934c - 02 S II 4815.55 V9 4s 4P 4p 48* 6 4 4860.43 0.000 l.OOOe 4- 02 H 4 4861.33 H4 2p-t- 2P* 4d+ 2D 8 32 4880.05 -0.005 1.191c - 01 [Fe III] 4881.11 F2 3d6 5D 3d6 3H 9 9 4889.81 -0.007 2.275c - 02 O II 4890.86 V28 3p 4S* 3d 4P 4 2 4905.90 -0.011 3.372c - 02 O II 4906.83 V28 3p 48* 3d 4P 4 4 4921.06 -0.015 1.650c-t-00 He I 4921.93 V48 2p IP* 4d ID 3 5 -0.016 * O II 4924.53 V28 3p 48* 3d 4P 4 6 4930.28 -0.017 9.238c - 02 [O HI] 4931.80 FI 2p2 3P 2p2 ID 1 5 4957.55 -0.024 2.622c + 02 [O HI] 4958.91 FI 2p2 3P 2p2 ID 3 5 A.I. Galactic Planetary Nebulae 219

-^obs (Â) f(A) /(A) Ion Ao Mult Lower term Upper term gl g2 5005.47 -0.036 7.945e -f 02 [0 III] 5006.84 FI 2p2 3P 2p2 ID 5 5 5189.23 -0.081 1.449e - 01 [Ar III] 5191.82 * * 5197.25 -0.083 1.103e-f 00 [N I] 5199.84 * * 5516.32 -0.154 2.730e - 01 [Cl III] 5517.66 FI 2p3 48* 2p3 2D* 4 6 5536.25 -0.158 7.362e - 01 [Cl III] 5537.60 FI 2p3 48* 2p3 2D* 4 4 5665.14 -0.180 8.983e - 02 N II 5666.63 V3 3s 3P* 3p 3D 3 5 5678.13 -0.182 2.387e - 01 N II 5679.56 V3 3s 3P* 3p 3D 5 7

5753.23 -0.195 4.517e -H 00 [N II] 5754.60 F3 2p2 ID 2p2 18 5 1 5874.26 -0.215 1.892e-t-01 He I 5875.66 Vll 2p 3P* 3d 3D 9 15 6299.16 -0.282 5.367e -h 00 [O I] 6300.34 FI 2p4 3P 2p4 ID 5 5 6310.96 -0.283 3.1866 4- 00 [S III] 6312.10 F3 2p2 ID 2p2 18 5 1 6345.67 -0.289 6.217e - 02 Si II 6347.10 V2 4s 28 4p 2P* 2 4 6362.69 -0.291 1.814e-H00 [O I] 6363.78 FI 2p4 3P 2p4 ID 3 5 6460.91 -0.306 6.722e - 02 C II 6461.95 4f 2F* 6 g 2 G 14 18

6546.94 -0.318 6.528e -H 01 [N II] 6548.10 FI 2p2 3P 2p2 ID 3 5 6561.73 -0.320 2.959e -f 02 H 3 6562.77 H3 2p-t- 2P* 3d+ 2D 8 18 6582.46 -0.323 2.044e + 02 [N II] 6583.50 FI 2p2 3P 2p2 ID 5 5 6677.00 -0.336 5.036e -f 00 He I 6678.16 V46 2p IP* 3d ID 3 5 6715.28 -0.342 5.309e 4- 00 [S II] 6716.44 F2 2p3 48* 2p3 2D* 4 6 6729.66 -0.344 1.0026-K 01 [S II] 6730.82 F2 2p3 48* 2p3 2D* 4 4 7063.78 -0.387 8.170e -f 00 He I 7065.25 VIO 2p 3P* 3s 38 9 3 7134.31 -0.396 2.808e -t- 01 [Ar III] 7135.80 FI 3p4 3P 3p4 ID 5 5 7233.23 -0.408 4.090e - 01 C II 7236.42 V3 3p 2P* 3d 2D 4 6 7279.63 -0.414 7.493e - 01 He I 7281.35 V45 2p IP* 3s 18 3 1 7317.73 -0.418 6.2546 -f 00 [O II] 7318.92 F2 2p3 2D* 2p3 2P* 6 2 7328.48 -0.420 5.6266 + 00 [0 II] 7329.67 F2 2p3 2D* 2p3 2P* 4 2

c(H/3)(Bal) = 0.524, c(H/3)(rad) = 0.430, c(H/3)^^G40 _ Q.605

Table A.7: NGC 5882 linelist.

-^obs (^ ) f(A) /(A) Ion Ao Mult Lower term Upper term gl g2

1640 1.786-HOl He II 1640 1663 4.90e + 00 O III 1663 1750 4.50e-f 00 N HI] 1750 1908 3.14e-t-01 C HI] 1908 3671.78 0.266 4.6226 - 01 H 24 3671.48 H24 2p+ 2P* 24d-t- 2D 8 * 3674.04 0.265 4.8866 - 01 H 23 3673.74 H23 2p-l- 2P* 23d4- 2D 8 * 3676.66 0.265 5.2516 - 01 H 22 3676.36 H22 2p4- 2P* 22d-H 2D 8 * 3679.66 0.265 5.8126 - 01 H 21 3679.36 H21 2p+ 2P* 21d+ 2D 8 * A.I. Galactic Planetary Nebulae 2 2 0

-^obs (Â) f(A) /(A) Ion Ao Mult Lower term Upper term gl g2 3683.12 0.264 6.281e - 01 H 20 3682.81 H20 2p+ 2P* 20d+ 2D 8 * 3687.14 0.263 7.415e - 01 H 19 3686.83 H19 2p+ 2P* 19d+ 2D 8 * 3691.87 0.262 8.416e - 01 H 18 3691.56 H18 2p+ 2P* 18d+ 2D 8 * 3697.46 0.262 9.575e - 01 H 17 3697.15 H17 2p+ 2P* 17d+ 2D 8 * 3703.89 0.260 1.237e + 00 H 16 3703.86 H16 2p+ 2P* 16d+ 2D 8 * 3705.05 0.260 1.016e + 00 He I 3705.02 * * 3707.28 0.260 3.275e -01 O III 3707.25 * * 3712.33 0.259 1.610e + 00 H 15 3711.97 H15 2p+ 2P* 15d+ 2D 8 * 3715.44 0.259 2.475e - 01 O III 3715.08 V14 3p 3P 3d 3D* 5 7 0.257 * [S III] 3721.63 F2 3p2 3P 3p2 IS 3 1 3722.16 0.257 2.942e + 00 H 14 3721.94 H14 2p+ 2P* 14d+ 2D 8 * 3726.41 0.257 1.022e + 01 [O II] 3726.03 FI 2p3 4S* 2p3 2D* 4 4 3729.20 0.256 4.982e + 00 [O II] 3728.82 FI 2p3 4S* 2p3 2D* 4 6 3734.75 0.255 2.294e + 00 H 13 3734.37 H13 2p+ 2P* 13d+ 2D 8 * 3750.50 0.253 3.047e + 00 H 12 3750.15 H12 2p+ 2P* 12d+ 2D 8 * 3755.09 0.252 3.922e - 01 O III 3754.69 V2 3s 3P* 3p 3D 3 5 3757.59 0.251 2.598e - 01 O III 3757.24 V2 3s 3P* 3p 3D 1 3 3760.22 0.251 6.957e - 01 O III 3759.87 V2 3s 3P* 3p 3D 5 7 3770.98 0.249 3.665e + 00 H 11 3770.63 Hll 2p+ 2P* l ld + 2D 8 * 3774.28 0.248 1.016e - 01 O III 3774.02 V2 3s 3P* 3p 3D 3 3 0.247 * 7 3781.40 * * 3791.65 0.245 1.244e - 01 O III 3791.27 V2 3s 3P* 3p 3D 5 5 3798.29 0.244 4.991e + 00 H 10 3797.90 HIO 2p+ 2P* 10d+ 2D 8 * 3806.38 0.242 1.034e - 01 He I 3805.74 * * 0.241 * 7 3811.24 * * 0.241 * He II 3813.50 4.19 4f+ 2F* 19g+ 2G 32 * 3820.05 0.240 1.499e + 00 He I 3819.62 V22 2p 3P* 6d 3D 9 15 3835.82 0.237 7.011e + 00 H 9 3835.39 H9 2p+ 2P* 9d+ 2D 8 * 3869.20 0.230 9.314e + 01 [Ne III] 3868.75 FI 2p4 3P 2p4 ID 5 5 3889.09 0.226 6.832e + 00 He I 3888.65 * * 3889.49 0.226 1.141e+ 01 H 8 3889.05 V2 2s 3S 3p 3P* 3 9 3927.14 0.219 2.983e + 01 He I 3926.54 * * 3967.90 0.211 2.810e + 01 [Ne HI] 3967.46 FI 2p4 3P 2p4 ID 3 5 3970.51 0.210 1.532e + 01 H 7 3970.07 * * 4009.06 0.202 2.452e - 01 He I 4009.26 V55 2p IP* 7d ID 3 5 4026.02 0.198 2.693e + 00 He I 4026.21 V18 2p 3P* 5d 3D 9 15 4026.02 0.198 2.694e + 00 He I 4026.21 V18 2p 3P* 5d 3D 9 15

0.196 * N II 4035.08 V39a 3d 3F* 4f 2 < 4 > 5 7 A.I. Galactic Planetary Nebulae 221

-^obs (Â) f(A) /(A) Ion Ao Mult Lower term Upper term gl g2 4041.24 0.195 2.727e - 02 N II 4041.31 V39b 3d 3F* 4f 2 < 5 > 9 11 4067.63 0.189 1.024e - 01 C III 4067.94 41 3F* 5g 3G * * 4068.36 0.189 3.794e- 01 [S II] 4068.60 FI 2p3 48* 2p3 2P* 4 4 4068.73 0.189 1.341e - 01 C III 4068.91 ** 4069.52 0.189 1.361e - 01 O II 4069.78 * * 4070.06 0.189 1.750e - 01 C III 4070.26 ** 4071.92 0.188 2.385e -01 O II 4072.16 VIO 3p 4D* 3d 4F 6 8 4075.62 0.187 2.754e - 01 O II 4075.86 VIO 3p 4D* 3d 4F 8 10 4076.11 0.187 1.246e - 01 [S II] 4076.35 FI 2p3 48* 2p3 2P* 2 4 0.187 * O II 4078.84 VIO 3p 4D* 3d 4F 4 4 4083.45 0.186 2.104e - 02 0 II 4083.90 V48b 3d 4F 4f G4* 6 8 4084.66 0.185 2.899e - 02 O II 4085.11 VIO 3p 4D* 3d 4F 6 6 4086.70 0.185 3.588e - 02 0 II 4087.15 V48c 3d 4F 4f G3* 4 6 4088.84 0.184 1.452e - 01 O II 4089.29 V48a 3d 4F 41 G5* 10 12 0.184 * O II 4092.93 VIO 3p 4D* 3d 4F 8 8 4097.07 0.183 1.314e + 00 N III 4097.33 VI 3s 28 3p 2P* 2 4 4101.48 0.182 2.664e + 01 H 6 4101.74 H6 2p+ 2P* 6d+ 2D 8 72 4110.48 0.180 3.5926 —02 O II 4110.78 V20 3p 4P* 3d 4D 4 2 4115.79 0.178 2.806e - 02 Si IV 4116.10 ** 4118.94 0.178 9.185e - 02 O II 4119.22 V20 3p 4P* 3d 4D 6 8 4120.56 0.177 2.543e - 01 He I 4120.84 V16 2p 3P* 5s 38 9 3 4121.18 0.177 2.8926 —02 0 II 4121.46 V19 * * 4132.54 0.175 5.3536 - 02 O II 4132.80 V19 * * 4143.49 0.172 4.0826 - 01 He I 4143.76 V53 2p IP* 6d ID 3 5 4153.00 0.170 9.5506 - 02 0 II 4153.30 V19 3p 4P* 3d 4P 4 6 4156.23 0.169 7.2756 - 02 O II 4156.53 V19 3p 4P* 3d 4P 6 4 4168.91 0.167 7.5906 - 02 O II 4169.22 V19 3p 4P* 3d 4P 6 6 4185.16 0.163 4.2536 - 02 O II 4185.45 V36 3p’ 2F* 3d’ 2G 6 8 4186.61 0.162 3.7616 - 02 C HI 4186.90 V18 4f IF* 5g IG 7 9 4189.50 0.162 6.5416 - 02 O II 4189.79 V36 3p’ 2F* 3d’ 2G 8 10 4195.59 0.160 6.436e - 02 N HI 4195.76 V6 3s’ 2P* 3p’ 2D 2 4 4199.66 0.160 1.2316 - 01 He II 4199.83 4.11 4f+ 2F* l lg + 2G 32 * 4215.51 0.156 1.9996 - 02 N III 4215.69 * * 4219.38 0.155 3.7536 —02 Ne II 4219.74 V52a 3d 4D 41 2<4>* 8 10 4227.38 0.153 2.8016 — 02 N II 4227.74 V33 3p ID 4s IP* 5 3 4231.78 0.152 1.4416 —02 Ne II 4231.64 V52b 3d 4D 4f 2<3>* 6 8 4237.05 0.151 1.3936 —02 N II 4237.05 V48b 3d 3D* 4f 1< 4> 5 7 4241.38 0.150 3.0606 - 02 N II 4241.78 V48a 3d 3D* 4f 1< 3> 5 7 A.I. Galactic Planetary Nebulae 2 2 2

-^obs (Â) f(A) /(A) Ion Ao Mult Lower term Upper term g l g2 4253.59 0.147 2.570e - 02 O II 4254.00 V lO l 3d’ 2G 4 f’ H5* 18 22 4266.91 0.144 3.994e - 01 C II 4267.15 V6 3d 2D 4f 2F* 10 14 4275.38 0.142 5.785e - 02 O II 4275.55 V67a 3d 4D 4f F4* 8 10 4276.23 0.142 2.358e - 02 O II 4276.75 V67b 3d 4D 4f F3* 6 8 4277.48 0.142 2.645e - 02 0 II 4277.43 V67c 3d 4D 4f F2* 2 4 0.140 * O II 4282.96 V67c 3d 4D 4f F2* 4 6 0.140 * O II 4283.73 V67c 3d 4D 4f F2* 4 4 4291.14 0.138 2.213e - 02 O II 4291.25 V55 3d 4P 4f G3* 6 8 4292.10 0.138 7.793e - 03 O II 4292.21 V18c 3d 2F 4f F2* 6 6 4294.67 0.137 1.562e - 02 O II 4294.78 V53b 3d 4P 4f D2* 4 6 4303.51 0.135 5.725e - 02 O II 4303.61 V65a 3d 4D 4f G5* 8 10 4317.11 0.132 6.375e - 02 O II 4317.14 V2 3s 4P 3p 4P* 2 4 4319.32 0.132 3.33 le - 02 O II 4319.63 V2 3s 4P 3p 4P* 4 6 4340.19 0.127 4.778e + 01 H 5 4340.47 H5 2p+ 2P* 5d+ 2D 8 50 4345.19 0.125 l.lOOe - 01 O II 4345.56 V2 3s 4P 3p 4P* 4 2 4349.06 0.125 9.113e - 02 O II 4349.43 V2 3s 4P 3p 4P* 6 6 4353.18 0.124 1.493e - 02 O II 4353.59 V76c 3d 2F 4f G3* 6 8 4362.95 0.121 5.569e + 00 [0 III] 4363.21 F2 2p2 ID 2p2 IS 5 1 4366.54 0.120 6.046e - 02 O II 4366.89 V2 3s 4P 3p 4P* 6 4 4378.86 0.118 1.361e - 01 N III 4379.11 V18 4f 2F* 5g 2G 14 18 4385.79 0.116 5.310e - 02 O II 4386.01 * * 4387.71 0.115 6.009e - 01 He I 4387.93 V51 2p IP* 5d ID 3 5 4391.77 0.115 3.441e - 02 Ne II 4391.99 V55e 3d 4F 4f 2< 5> * 10 12 4409.02 0.110 2.324e - 02 Ne II 4409.30 V55e 3d 4F 4f 2<5>* 8 10 4414.62 0.109 4.390e - 02 0 II 4414.90 V5 3s 2P 3p 2D* 4 6 4416.69 0.109 3.987e - 02 O II 4416.97 V5 3s 2P 3p 2D* 2 4 4437.33 0.103 5.770e - 02 He I 4437.55 V50 2p IP* 5s IS 3 1 4465.38 0.097 3.047e - 02 O II 4466.42 V86b 3d 2P 4f D2* 4 6 4471.23 0.095 5.650e + 00 He I 4471.50 V14 2p 3P* 4d 3D 9 15 4491.10 0.091 2.034e - 02 O II 4491.23 V86a 3d 2P 4f D3* 4 6 4510.71 0.086 7.873e - 02 N HI 4510.91 V3 3s’ 4P* 3p’ 4D 4 6 4514.66 0.085 2.576e - 02 N HI 4514.86 V3 3s’ 4P* 3p’ 4D 6 8 4517.95 0.084 2.984e - 02 N III 4518.15 V3 3s’ 4P* 3p’ 4D 2 2 4523.38 0.083 2.395e - 02 N HI 4523.58 V3 3s’ 4P* 3p’ 4D 4 4 4530.21 0.081 1.795e - 02 N II 4530.41 V58b 3d IF* 4f 2 < 5 > 7 9 4534.38 0.080 1.952e - 02 N HI 4534.58 V3 3s’ 4P* 3p’ 4D 6 6 4541.31 0.078 6.485e - 02 He II 4541.59 4.9 41+ 2F* 9g+ 2G 32 * 4544.52 0.078 3.312e - 02 N HI 4544.80 * * A.I. Galactic Planetary Nebulae 223

-^obs {N) f(A) /(A) Ion Ao Mult Lower term Upper term g i g2 0.073 * Mg I] 4562.60 3s2 IS 3s3p 3P* 1 5 4570.81 0.071 2.908e - 02 Mg I] 4571.10 3s2 IS 3s3p 3P* 1 3 4590.89 0.066 5.708e - 02 O II 4590.97 V15 3s’ 2D 3p’ 2F* 6 8 4596.10 0.065 4.221e - 02 O II 4596.18 V15 3s’ 2D 3p’ 2F* 4 6 4601.56 0.064 2.063e - 02 N II 4601.48 V5 3s 3P* 3p 3P 3 5 0.062 * N II 4607.16 V5 3s 3P* 3p 3P 1 3 4609.52 0.062 6.087e - 02 O II 4609.44 V92a 3d 2D 4f F4* 6 8 0.062 * O II 4610.20 V92c 3d 2D 4f F2* 4 6 4620.48 0.059 2.880e - 02 N II 4621.39 V5 3s 3P* 3p 3P 3 1 4630.55 0.057 3.812e - 02 N II 4630.54 V5 3s 3P* 3p 3P 5 5 4633.91 0.056 5.735e - 01 N III 4634.14 V2 3p 2P* 3d 2D 2 4 4638.63 0.055 1.690e- 01 O II 4638.86 VI 3s 4P 3p 4D* 2 4 4640.41 0.054 1.125e + 00 N III 4640.64 V2 3p 2P* 3d 2D 4 6 4641.58 0.054 3.669e - 01 O II 4641.81 V I 3s 4P 3p 4D* 4 6 4647.21 0.053 9.597e - 02 C III 4647.42 VI 3s 3S 3p 3P* 3 5 4648.92 0.052 3.447e - 01 O II 4649.13 VI 3s 4P 3p 4D* 6 8 4650.04 0.052 5.753e - 02 C III 4650.25 VI 3s 3S 3p 3P* 3 3 4650.63 0.052 1.141e - 01 O II 4650.84 VI 3s 4P 3p 4D* 2 2 4651.26 0.052 1.917e - 02 C III 4651.47 VI 3s 3S 3p 3P* 3 1 4657.92 0.050 1.567e - 01 [Fe III] 4658.10 F3 3d6 5D 3d6 3F2 9 9 4661.45 0.049 1.188e - 01 O II 4661.63 VI 3s 4P 3p 4D* 4 4 4673.59 0.046 1.878e - 02 O II 4673.73 V I 3s 4P 3p 4D* 4 2 4676.10 0.046 9.064e - 02 O II 4676.24 VI 3s 4P 3p 4D* 6 6 4685.54 0.043 2.751e-kOO He II 4685.68 3.4 3d4- 2D 4f-t- 2F* 18 32 4688.69 0.042 3.771e - 02 7 4688.69 * * 0.040 * O II 4699.22 V25 3p 2D* 3d 2F 4 6 4701.45 0.039 3.694e - 02 [Fe III] 4701.62 F3 3d6 5D 3d6 3F2 7 7 4711.18 0.037 2.608e -f 00 [Ar IV] 4711.37 FI 3p3 4S* 3p3 2D* 4 6 4712.98 0.036 6.619e - 01 He I 4713.17 vl2 2P 3p* 4S 3s 9 3 4733.83 0.031 2.259e - 02 [Fe HI] 4733.93 ** 4740.05 0.030 2.995e -f 00 [Ar IV] 4740.17 FI 3p3 4S* 3p3 2D* 4 4 4754.60 0.026 2.717e - 02 [Fe HI] 4754.83 ** 4769.49 0.022 1.746e - 02 [Fe HI] 4769.60 * * 4803.00 0.015 2.322e - 02 C II 4802.23 4f 2F* 8g 2G 14 18 4860.80 0.000 l.OOOe -f 02 H 4 4861.33 H4 2p+ 2P* 4d+ 2D 8 32 4881.03 -0.005 6.343e - 02 [Fe HI] 4881.11 F2 3d6 5D 3d6 3H 9 9 4891.60 -0.007 8.668e - 03 O II 4890.86 V28 3p 4S* 3d 4P 4 2 4899.28 -0.009 8.684e - 03 ? 4899.28 * * A.I. Galactic Planetary Nebulae 224

-^obs ( Â ) f(A) /(A) Ion Ao Mult Lower term Upper term gi g2 4902.59 -0.010 2.848e - 02 Si II 4902.65 * * 4906.53 -0.011 3.621e - 02 O II 4906.83 V28 3p 4S* 3d 4P 4 4 4921.84 -0.015 1.406e + 00 He I 4921.93 V48 2p IP* 4d ID 3 5 4924.44 -0.016 6.563c — 02 O II 4924.53 V28 3p 48* 3d 4P 4 6 4931.02 -0.017 1.379e - 01 [0 III] 4931.80 FI 2p2 3P 2p2 ID 1 5 4958.41 -0.024 3.535e-h 02 [O III] 4958.91 FI 2p2 3P 2p2 ID 3 5 5006.33 -0.036 1.057e -b 03 [O III] 5006.84 FI 2p2 3P 2p2 ID 5 5 5411.19 -0.134 1.685e - 01 He II 5411.52 4.7 4f-b 2F* 7g+ 2 0 32 98 5517.46 -0.154 4.228c - 01 [Cl HI] 5517.66 FI 2p3 48* 2p3 2D* 4 6 5537.40 -0.158 5.122c - 01 [Cl III] 5537.60 FI 2p3 48* 2p3 2D* 4 4 5754.21 -0.195 2.682c - 01 [N II] 5754.60 F3 2p2 ID 2p2 18 5 1 5800.76 -0.202 1.545c - 01 C IV 5801.51 VI 3s 28 3p 2P* 2 4 5811.39 -0.204 1.176c - 01 C IV 5812.14 VI 3s 28 3p 2P* 2 2 5875.26 -0.215 1.600c -b 01 He I 5875.66 Vll 2p 3P* 3d 3D 9 15 6101.74 -0.251 1.200c - 01 [K IV] 6101.83 FI 3p4 3P 3d4 ID 5 5 6300.31 -0.282 3.829c - 01 [0 I] 6300.34 FI 2p4 3P 2p4 ID 5 5 6312.07 -0.283 1.202c -b 00 [S HI] 6312.10 F3 2p2 ID 2p2 18 5 1 6364.00 -0.291 1.585c - 01 [OI] 6363.78 FI 2p4 3P 2p4 ID 3 5 6462.90 -0.306 7.430c - 02 C II 6461.95 4f 2F* 6g 2 0 14 18 6548.22 -0.318 4.877e -b 00 [N II] 6548.10 FI 2p2 3P 2p2 ID 3 5 6562.85 -0.320 2.958c -b 02 H 3 6562.77 H3 2p-b 2P* 3d-b 2D 8 18 6583.40 -0.323 1.445c-b 01 [N II] 6583.50 FI 2p2 3P 2p2 ID 5 5 6678.27 -0.336 4.314c -b 00 He I 6678.16 V46 2p IP* 3d ID 3 5 6716.55 -0.342 1.119c-b 00 [S II] 6716.44 F2 2p3 48* 2p3 2D* 4 6 6730.93 -0.344 1.911c + 00 [S II] 6730.82 F2 2p3 48* 2p3 2D* 4 4 7064.98 -0.387 4.923c -b 00 He I 7065.25 VIO 2p 3P* 3s 38 9 3 7111.38 -0.393 4.206c - 02 C I 7111.50 * * 7135.52 -0.396 1.420c -b 01 [Ar III] 7135.80 FI 3p4 3P 3p4 ID 5 5 7229.89 -0.408 7.600c - 02 ? 7229.89 ** 7236.02 -0.408 1.654c - 01 C II 7236.42 V3 3p 2P* 3d 2D 4 6 7262.36 -0.412 8.257c - 02 [Ar IV] 7262.76 F2 3p3 2D* 3p3 2P* 4 2 7280.95 -0.414 6.097c - 01 He I 7281.35 V45 2p IP* 3s 18 3 1 7318.96 -0.418 9.383c - 01 [0 II] 7318.92 F2 2p3 2D* 2p3 2P* 6 2 7329.71 -0.420 8.569c - 01 [O II] 7329.67 F2 2p3 2D* 2p3 2P* 4 2

c(H/3)(Bal) = 1.44, c(H/3)(rad) = 1.47

Table A.8: NGC6302 linelist.

A o b s (^) f(A) /(A) Ion Ao Mult Lower term Upper term gi g2

1240 0.260 1.382c -b 03 N V 1240 A.I. Galactic Planetary Nebulae 225

-^obs (Â) f(A) /(A) Ion Ao Mult Lower term Upper term gl g2 1401 0.260 1.119e + 02 O IV] 1401 1484 0.260 5.336e + 02 N IV] 1484 1549 0.260 5.516e + 02 C IV 1549 1601 0.257 2.661e + 00 [Ne IV] 1601 1640 0.260 4.094e + 02 He II 1640 1663 0.260 1.292e + 02 O III] 1663 1750 0.260 4.077e + 02 N III] 1750 1908 0.260 4.629e + 02 C III] 1908 2423 0.260 7.946e + 02 [Ne IV] 2423 3118.22 0.453 5.249e + 00 ? * ** 3120.54 0.452 5.8616 + 00 O III 3121.71 V12 3p 3S 3d 3P* 3 3 3131.16 0.447 7.339e + 00 O III 3132.79 V12 3p 3S 3d 3P* 3 5 3184.38 0.424 4.964e - 01 S illl 3185.12 V8 4p IP* 5s IS 3 1 3187.00 0.423 3.3816 + 00 He I 3187.74 V3 2s 3S 4p 3P* 3 9 3202.37 0.416 2.8286 + 01 He II 3203.10 3.5 3d+ 2D 5f+ 2F* 18 50 3239.26 0.402 2.6726 - 01 7 * * * 3241.15 0.401 1.4926 + 00 Silll 3241.62 V6 4p 3P* 5s 3S 5 3 3298.89 0.379 2.8666 + 00 O III 3299.36 V3 3s 3P* 3p 38 1 3 3311.83 0.374 7.8776 + 00 O HI 3312.30 V3 3s 3P* 3p 3S 3 3 3340.26 0.364 1.2486 + 01 O HI 3340.74 V3 3s 3P* 3p 38 5 3 3345.38 0.362 8.3306 + 01 [Ne V] 3345.86 2s2 3P 2s2 ID 3 5 3354.22 0.359 1.8476 - 01 ? * ** 3361.83 0.356 4.869e - 01 ? * ** 3380.59 0.350 2.4076 - 01 O IV 3381.21 2s 4P* 2s 4D 4 6 3384.90 0.348 1.8386 - 01 O IV 3385.52 2s 4P* 2s 4D 6 8 3405.33 0.341 4.3866 - 01 ? * * * 3411.35 0.339 1.9736 - 01 O IV 3411.69 2s2 2P* 2s2 2D 4 6 3414.92 0.338 8 .2 1 3 6 -0 1 7 * ** 3425.45 0.334 2.3316 + 02 [Ne V] 3425.86 2s2 3P 2s2 ID 5 5 3429.65 0.333 1.8886 + 00 O HI 3428.65 V15 3s 3P* 3d 3P* 3 5 3432.71 0.332 3.372e - 01 7 * ** 3443.66 0.329 2.0536 + 01 0 HI 3444.07 V15 3s 3P* 3d 3P* 5 5 3466.13 0.322 1.1736 + 00 7 * ** 3478.27 0.318 1.8166 - 01 7 * * * 3553.98 0.295 2.2466 - 01 7 ** * 3567.92 0.290 2.1456 - 01 Ne II 3568.50 V9 3s' 2D 3p’ 2F* 6 8 3586.70 0.285 4.3536 - 01 He I 3587.28 V31 2p 3P* 9d 3D 9 15 3613.18 0.278 3.0276 - 01 ? * ** A.I. Galactic Planetary Nebulae 226

-^obs (Â) f(A) /(A) Ion Ao Mult Lower term Upper term gl g2 3633.66 0.272 4.177e - 01 He I 3634.25 V28 2p 3P* 8d 3D 9 15 3671.16 0.266 1.738e - 01 H 24 3671.48 H24 2p+ 2P* 24d+ 2D 8 * 3673.38 0.265 3.375e - 01 H 23 3673.74 H23 2p+ 2P* 23d+ 2D 8 * 3676.03 0.265 4.816e-01 H 22 3676.36 H22 2p+ 2P* 22d+ 2D 8 * 3678.99 0.265 6.251e - 01 H 21 3679.36 H21 2p+ 2P* 21d+ 2D 8 * 3682.46 0.264 7.437e - 01 H 20 3682.81 H20 2p+ 2P* 20d+ 2D 8 * 3686.46 0.263 9.336e - 01 H 19 3686.83 H19 2p+ 2P* 19d+ 2D 8 * 3691.18 0.262 1.074e + 00 H 18 3691.56 H18 2p+ 2P* 18d+ 2D 8 * 3696.81 0.262 1.548e + 00 H 17 3697.15 H17 2p+ 2P* 17d+ 2D 8 * 3703.62 0.260 2.846e + 00 H 16 3703.86 H16 2p+ 2P* 16d+ 2D 8 * 3711.64 0.259 2.526e + 00 H 15 3711.97 H15 2p+ 2P* 15d+ 2D 8 * 3714.54 0.259 2.239e - 01 O III 3715.08 V14 3p 3P 3d 3D* 5 7 3721.14 0.257 4.966e + 00 [S III] 3721.63 F2 3p2 3P 3p2 IS 3 1 3725.54 0.257 3.003e + 01 [O II] 3726.03 FI 2p3 4S* 2p3 2D* 4 4 3728.33 0.256 1.405e + 01 [O II] 3728.82 FI 2p3 4S* 2p3 2D* 4 6 3733.88 0.255 1.799e + 00 H 13 3734.37 H13 2p+ 2P* 13d+ 2D 8 * 3749.69 0.253 2.322e + 00 H 12 3750.15 H12 2p+ 2P* 12d+ 2D 8 * 3754.19 0.252 6.991e - 01 O III 3754.69 V2 3s 3P* 3p 3D 3 5 3756.74 0.251 5.165e - 01 O III 3757.24 V2 3s 3P* 3p 3D 1 3 3759.37 0.251 3.762e + 00 O III 3759.87 V2 3s 3P* 3p 3D 5 7 3770.17 0.249 2.868e + 00 H 11 3770.63 Hll 2p+ 2P* l ld + 2D 8 * 3773.31 0.248 8.268e - 02 O III 3774.02 V2 3s 3P* 3p 3D 3 3 3781.28 0.247 1.338e - 01 He II * * * 3790.86 0.245 1.810e - 01 O HI 3791.27 V2 3s 3P* 3p 3D 5 5 3797.46 0.244 4.007e + 00 H 10 3797.90 HIO 2p+ 2P* 10d+ 2D 8 * 3805.53 0.242 5.540e - 02 7 * * * 3813.12 0.241 2.067e - 01 He II 3813.50 4.19 4f+ 2F* 19g+ 2G 32 * 3819.27 0.240 l.lOle + 00 He I 3819.62 V22 2p 3P* 6d 3D 9 15 3834.98 0.237 6.273e + 00 H 9 3835.39 H9 2p+ 2P* 9d+ 2D 8 * 3839.10 0.236 5.077e - 02 [Fe V] 3839.27 3d4 5D 3d4 3F2 7 7 3855.83 0.233 3.226e - 01 ? * ** 3857.65 0.233 3.286e - 01 He II 3858.07 4.17 4f+ 2F* 17g+ 2G 32 * 3868.38 0.230 * [Ne HI] 3868.75 FI 2p4 3P 2p4 ID 5 5 3888.47 0.226 1.816e + 01 He I 3888.65 V2 2s 3S 3p 3P* 3 9 3895.42 0.225 2.132e - 02 ? * ** 3923.18 0.220 4.454e - 01 He II 3923.48 4.15 4f+ 2F* 15g+ 2G 32 * 3926.23 0.219 1.048e + 00 ? *** 3967.11 0.211 5.567e 4- 01 [Ne HI] 3967.46 FI 2p4 3P 2p4 ID 3 5 A.I. Galactic Planetary Nebulae 227

-^obs (.N) f(A) /(A) Ion Ao Mult Lower term Upper term gl g2 3975.44 0.209 4.479e - 02 ? * * * 4008.91 0.202 1.828e - 01 He I 4009.26 V55 2p IP* 7d ID 3 5 4025.65 0.198 3.002e + 00 He I 4026.21 V18 2p 3P* 5d 3D 9 15 4059.90 0.191 6.773e - 02 O II 4060.60 2s2 2F 2s2 2 [4] 8 10 4068.23 0.189 1.453e + 01 [S II] 4068.60 FI 2p3 4S* 2p3 2P* 4 4 4075.98 0.187 4.808e + 00 [S II] 4076.35 FI 2p3 4S* 2p3 2P* 2 4 4096.89 0.183 3.090e + 00 N HI 4097.33 VI 3s 2S 3p 2P* 2 4 4101.31 0.182 2.684e + 01 H 6 4101.74 H6 2p+ 2P* 6d + 2D 8 72 4120.43 0.177 4.038e - 01 He I 4120.84 V16 2p 3P* 5s 3S 9 3 4122.08 0.177 3.828e - 01 ? 4122.08 * * 4143.35 0.172 3.461e - 01 He I 4143.76 V53 2p IP* 6d ID 3 5 4162.88 0.168 4.078e - 01 ? 4162.88 * * 4199.43 0.160 1.481e + 00 He II 4199.83 4.11 4f+ 2F* l l g + 2G 32 * 4266.67 0.144 1.625e - 01 C II 4267.15 V6 3d 2D 4f 2F* 10 14 4335.09 0.129 1.268e - 01 O II 4334.09 V63b 3d 4D 4f D2* 4 4 4340.01 0.127 4.839e + 01 H 5 4340.47 H5 2p+ 2P* 5d + 2D 8 50 4362.75 0.121 3.967e +01 [O III] 4363.21 F2 2p2 ID 2p2 IS 5 1 4366.39 0.120 1.136e - 01 O II 4366.89 V2 3s 4P 3p 4P* 6 4 4378.63 0.118 1.819e - 01 N III 4379.11 V18 4f 2F* 5g 2G 14 18 4387.45 0.115 6.233e - 01 He I 4387.93 V51 2p IP* 5d ID 3 5 4415.15 0.109 1.369e - 01 O II 4414.90 V5 3s 2P 3p 2D* 4 6 4468.12 0.095 5.925e + 00 He I 4471.50 V14 2p 3P* 4d 3D 9 15 4510.45 0.086 1.742e - 01 N HI 4510.91 V3 3s' 4P* 3p' 4D 4 6 4519.30 0.084 1.449e - 01 N III 4518.15 V3 3s' 4P* 3p' 4D 2 2 4541.20 0.078 2.838e + 00 He II 4541.59 4.9 4f+ 2F* 9g+ 2G 32 * 4562.43 0.073 2.486e - 01 Mg I] 4562.60 3s2 IS 3s3p 3P* 1 5 4570.71 0.071 2.474e - 01 Mg I] 4571.10 3s2 IS 3s3p 3P* 1 3 4590.91 0.066 4.237e - 02 O II 4590.97 V15 3s' 2D 3p' 2F* 6 8 4603.52 0.063 1.657e - 0 1 ? 4603.52 * * 4606.06 0.063 6.675e - 01 [Fe HI] 4606.50 * * 4606.72 0.062 6.073e - 02 N II 4607.16 V5 3s 3P* 3p 3P 1 3 4625.04 0.058 4.164e - 01 ? ** * 4629.68 0.057 7.433e - 02 N II 4630.54 V5 3s 3P 3p 3P 5 5 4631.93 0.056 2.280e - 01 7 4631.93 * * 4633.79 0.056 2.892e + 00 N HI 4634.14 V2 3p 2P* 3d 2D 2 4 4638.49 0.055 4.571e - 01 O II 4638.86 VI 3s 4P 3p 4D* 2 4 4640.27 0.054 5.466e + 00 N HI 4640.64 V2 3p 2P* 3d 2D 4 6 4641.44 0.054 7.555e - 01 O II 4641.44 * * A.I. Galactic Planetary Nebulae 228

-^obs (-À) f(A) /(A) Ion Ao Mult Lower term Upper term gl g2 4647.24 0.053 7.545e - 02 C III 4647.42 VI 3s 3S 3p 3P* 3 5 4648.95 0.052 1.229e - 01 O II 4649.13 VI 3s 4P 3p 4D* 6 8 4675.41 0.050 5.562e - 02 [Fe III] 4658.10 F3 3d6 5D 3d6 3F2 9 9 4675.82 0.046 1.148e - 01 O II 4676.24 VI 3s 4P 3p 4D* 6 6 4685.42 0.043 7.530c -j- 01 He II 4685.68 3.4 3d-f- 2D 4f+ 2F* 18 32 4701.28 0.039 5.428e - 02 [Fe III] 4701.62 F3 3d6 5D 3d6 3F2 7 7 4707.10 0.038 1.685e - 01 ? * ** 4711.05 0.037 1.269e-h01 [Ar IV] 4711.37 FI 3p3 4S* 3p3 2D* 4 6 4713.65 0.036 2.966e 00 He I 4713.17 V12 2P 3p* 4S 3s 9 3 4723.85 0.034 2.086e -f 00 [Ne IV] 4724.15 FI 2p3 2D* 2p3 2P* 4 4 4725.32 0.033 1.829e -f 00 [Ne IV] 4725.62 FI 2p3 2D* 2p3 2P* 4 2 4739.93 0.030 2.091e4-01 [Ar IV] 4740.17 FI 3p3 4S* 3p3 2D* 4 4 4802.41 0.015 6.543e - 02 C II 4802.23 4f 2F* 8g 2 0 14 18 4831.07 0.007 2.768e - 02 ? ** * 4861.11 0.000 l.OOOe -f 02 H 4 4861.33 H4 2p+ 2P* 4d-t- 2D 8 32 4880.73 -0.005 8.977e - 02 [Fe III] 4881.11 F2 3d6 5D 3d6 3H 9 9 4893.40 -0.008 5.308e — 02 ? 4893.40 * * 4921.75 -0.015 1.454e 00 He I 4921.93 * * 4930.63 -0.017 5.286e - 01 O III 4931.80 * * 4934.69 -0.018 5.566e — 02 ? ** * 4944.42 -0.020 2.253e - 01 O II 4943.00 * * 4957.76 -0.024 4.295e + 02 [O III] 4958.91 * * 5005.68 -0.036 1.290e -H 03 [0 III] 5006.84 FI 2p2 3P 2p2 ID 5 5 5039.66 -0.044 3.060e -t- 00 Si II 5039.66 * * 5054.61 -0.047 1.314e-H00 Si II 5054.61 * * 5190.01 -0.081 1.521e-H 00 [Ar HI] 5191.82 * * 5198.03 -0.083 1.081e-H01 [N I] 5199.84 * * 5409.63 -0.134 6.045e -|- 00 He II 5411.52 4.7 4f-t- 2F* 7g+ 2 0 32 98 5516.74 -0.154 3.394e - 01 [Cl HI] 5517.66 FI 2p3 4S* 2p3 2D* 4 6 5536.67 -0.158 8.610e - 01 [Cl HI] 5537.60 FI 2p3 4S* 2p3 2D* 4 4 5576.41 -0.164 2.328e - 01 [O I] 5577.34 * * 5668.07 -0.180 7.350e - 02 N II 5666.63 V3 3s 3P 3p 3D 3 5 5677.46 -0.181 1.580e - 01 [Fe VI] 5677.00 * * 5700.47 -0.185 7.720e - 02 ? 5700.47 * * 5720.00 -0.189 2.416e - 01 [Fe VII] 5720.00 * * 5753.69 -0.195 2.011e-H 01 [N II] 5754.60 F3 2p2 ID 2p2 IS 5 1 5874.74 -0.215 1.754e-k01 He I 5875.66 Vll 2p 3P* 3d 3D 9 15 5912.59 -0.221 7.143e- 02 He II 5913.26 * * A.I. Galactic Planetary Nebulae 229

-^obs (Â) f(A) /(A) Ion Ao Mult Lower term Upper term gl g2 5931.10 -0.224 1.048e - 01 N II 5931.78 V28 3p 3P 3d 3D 3 5 5952.26 -0.227 9.774e - 02 He II 5952.94 * * 5976.35 -0.231 1.176e - 01 He II 5977.03 * * 6004.77 -0.236 1.132e-t-00 He II 6004.73 5.22 5g+ 2G 22h4- 2H* 50 * 6037.15 -0.241 1.932e - 01 He II 6036.70 * * 6046.60 -0.242 3.502e - 01 7 6046.60 ** 6074.62 -0.247 3.765e - 01 He II 6074.10 * * 6085.33 -0.248 5.221e - 01 [Fe VII] 6086.90 ** 6100.80 -0.251 8.732e - 01 [KIV] 6101.83 FI 3p4 3P 3d4 ID 5 5 6117.65 -0.254 1.751e - 01 He II 6118.20 5.19 5g+ 2G 19h-f 2H* 50 * 6169.79 -0.262 2.024e - 01 He II 6170.60 5.18 5g+ 2G 18h+ 2H 5 * 6200.27 -0.266 5.818e - 02 ? 6200.27 * * 6227.73 -0.271 2.152e - 01 [KIV] 6228.40 * * 6232.99 -0.272 2.326e - 01 He II 6233.80 5.17 5g+ 2G 17h+ 2H 5 * 6299.93 -0.282 2.430e-h 01 [0 I] 6300.34 FI 2p4 3P 2p4 ID 5 5 6311.69 -0.283 5.8456 -j- 00 [S HI] 6312.10 F3 2p2 ID 2p2 IS 5 1 6346.73 -0.292 3.710e - 01 Si II 6371.38 ** 6363.41 -0.291 7.970e -f- 00 [0 I] 6363.78 FI 2p4 3P 2p4 ID 3 5 6371.01 -0.292 8.711e - 01 Si II 6371.01 * * 6393.32 -0.296 4.160e - 02 [Mn V] 6393.60 ** 6406.02 -0.297 3.596e - 01 He II 6406.30 5.15 5g4- 2G 15h+ 2H 5 * 6434.52 -0.302 5.025e -f 00 [ArV] 6435.10 3s2 3P 3s2 ID 3 5 6476.86 -0.308 8.994e - 02 7 * * * 6499.95 -0.311 8.125e - 02 7 * * * 6516.44 -0.313 7.622e - 02 ? * * * 6526.80 -0.315 5.275e - 01 He II 6527.11 5.14 5g4- 2G 14h+ 2H 5 * 6547.73 -0.318 1.740e -f 02 [N II] 6548.10 FI 2p2 3P 2p2 ID 3 5 6562.40 -0.320 2.837e -k 02 H 3 6562.77 H3 2p+ 2P* 3d+ 2D 8 18 6583.13 -0.323 5.220e 4- 02 [N II] 6583.50 FI 2p2 3P 2p2 ID 5 5 6677.97 -0.336 1.026e -k 00 He I 6678.16 V46 2p IP* 3d ID 3 5 6716.15 -0.342 l.OlOe-k 01 [S II] 6716.44 F2 2p3 4S* 2p3 2D* 4 6 6730.53 -0.344 1.9886 4-01 [S II] 6730.82 F2 2p3 4S* 2p3 2D* 4 4 6744.88 -0.345 6.9486 - 02 7 *** 6794.58 -0.352 1.2046 - 01 7 *** 6826.21 -0.356 4.047e - 01 ? * * * 6890.67 -0.365 5.5126 - 01 He II 6890.88 7 * * 7005.26 -0.380 9.677e + 00 [Ar V] 7005.67 3s2 3P 3s2 ID 5 5 7064.85 -0.387 8.2216 -f 00 He I 7065.25 VIO 2p 3P* 3s 3S 9 3 A.I. Galactic Planetary Nebulae 230

■^obs ( Â ) f(A) /(A) Ion Ao Mult Lower term Upper term g l g2 7135.26 -0.396 2.371e-f01 [Ar III] 7135.80 FI 3p4 3P 3p4 ID 5 5 7154.48 -0.398 1.553e - 01 ? * * * 7170.06 -0.401 1.255e 4- 00 [Ar IV] 7170.62 F2 3p3 2D* 3p3 2P* 4 4 7176.94 -0.401 7.441e - 01 He II 7177.50 5.11 5g+ 2G llh+ 2H 5 * 7236.17 -0.408 8.777e - 01 C II 7236.42 V3 3p 2P* 3d 2D 4 6 7262.51 -0.412 9.688e - 01 [Ar IV] 7262.76 F2 3p3 2D* 3p3 2P* 4 2 7281.09 -0.414 8.937e - 01 He I 7281.35 V45 2p IP* 3s IS 3 1 7319.04 -0.418 8.577e + 00 [O II] 7318.92 F2 2p3 2D* 2p3 2P* 6 2 7329.79 -0.420 7.285e-kOO [O II] 7329.67 F2 2p3 2D* 2p3 2P* 4 2

c(H/3)(Bal) = 0.473, c(H,5)(rad) = 0.361

Table A.9: NGC6818 linelist.

•^obs (^ ) f(A) /(A) Ion Ao Mult Lower term Upper term g l g2

1240 0.260 L590e4-01 N V 1240 1401 0.260 3.260e -f 01 O IV] 1401 1484 0.260 3.160e4-01 N IV] 1484 1549 0.260 9.670e -f 01 C IV 1549 1574 0.257 9.840e 4- 00 [Ne V] 1574 1601 0.257 6.940e 4- 00 [Ne IV] 1601 1640 0.260 3.180e-t-02 He II 1640 1663 0.260 2.980e 4- 01 O III] 1663 1750 0.260 2.120e4-01 N III] 1750 1908 0.260 4.525e 4- 02 C III] 1908 2326 0.260 2.600e 4- 01 C II] 1908 2423 0.260 5.074e 4- 01 [Ne IV] 2423 2470 0.260 3.000e 4- 00 [O II] 1908 3120.58 0.452 1.799e4-00 O HI 3121.71 V12 3p 3S 3d 3P* 3 3 3131.91 0.447 8.378e 4- 01 O HI 3132.79 V12 3p 3S 3d 3P* 3 5 3187.11 0.423 2.203e 4- 00 He I 3187.74 V3 2s 3S 4p 3P* 3 9 3202.55 0.416 3.049e4-01 He II 3203.10 3.5 3d+ 2D 514- 2F* 18 50 3241.35 0.401 5.879e - 01 S illl 3241.62 V6 4p 3P* 5s 38 5 3 3299.06 0.379 2.750e 4- 00 O HI 3299.36 V3 3s 3P* 3p 3S 1 3 3312.03 0.374 7.860e 4- 00 O HI 3312.30 V3 3s 3P* 3p 38 3 3 3340.55 0.364 1.202e4-01 O HI 3340.74 V3 3s 3P* 3p 38 5 3 3345.54 0.362 1.613e + 01 [Ne V] 3345.86 2s2 3P 2s2 ID 3 5 3380.95 0.350 3.073e - 01 O IV 3381.21 2s 4P* 2s 4D 4 6 3385.40 0.348 2.483e - 01 O IV 3385.52 2s 4P* 2s 4D 6 8 A.I. Galactic Planetary Nebulae 231

-^obs {N) f(A) /(A) Ion Ac Mult Lower term Upper term gl g2 3411.40 0.339 1.736e - 01 O IV 3411.69 2s2 2P* 2s2 2D 4 6 3425.59 0.334 5.008e + 01 [Ne V] 3425.86 2s2 3P 2s2 ID 5 5 3428.20 0.333 2.890e + 00 0 III 3428.65 V15 3s 3P* 3d 3P* 3 5 3443.77 0.329 2.347e + 01 O III 3444.07 V15 3s 3P* 3d 3P* 5 5 3568.12 0.290 1 .9 2 1 e -0 1 Ne II 3568.50 V9 3s’ 2D 3p’ 2F* 6 8 3586.73 0.285 1.637e - 0 1 He I 3587.28 V31 2p 3P* 9d 3D 9 15 3633.91 0.272 1.748e - 01 He I 3634.25 V28 2p 3P* 8d 3D 9 15 3671.13 0.266 5.008e - 01 H 24 3671.48 H24 2p+ 2P* 24d+ 2D 8 * 3673.39 0.265 5.220e - 01 H 23 3673.74 H23 2p+ 2P* 23d+ 2D 8 * 3676.01 0.265 5.694e - 01 H 22 3676.36 H22 2p+ 2P* 22d+ 2D 8 * 3679.01 0.265 6.212e - 01 H 21 3679.36 H21 2p+ 2P* 21d+ 2D 8 * 3682.46 0.264 6.541e - 01 H 20 3682.81 H20 2p+ 2P* 20d+ 2D 8 * 3686.48 0.263 7.205e - 01 H 19 3686.83 H19 2p+ 2P* 19d+ 2D 8 * 3691.21 0.262 8.119e - 01 H 18 3691.56 H18 2p+ 2P* 18d+ 2D 8 * 3696.72 0.262 9.666e - 01 H 17 3697.15 H17 2p+ 2P* 17d+ 2D 8 * 3703.43 0.260 1.594e + 00 H 16 3703.86 H16 2p+ 2P* 16d+ 2D 8 * 3711.62 0.259 1.641e + 00 H 15 3711.97 H15 2p+ 2P* 15d+ 2D 8 * 3714.73 0.259 2.640e - 01 O III 3715.08 V14 3p 3P 3d 3D* 5 7 3721.47 0.257 5.443e + 00 H 14 3721.94 H14 2p+ 2P* 14d+ 2D 8 * 3725.70 0.257 4.456e + 01 [O II] 3726.03 FI 2p3 4S* 2p3 2D* 4 4 3728.49 0.256 3.012e + 01 [O II] 3728.82 FI 2p3 4S* 2p3 2D* 4 6 3734.04 0.255 2.437e + 00 H 13 3734.37 H13 2p+ 2P* 13d+ 2D 8 * 3749.78 0.253 3.171e + 00 H 12 3750.15 H12 2p+ 2P* 12d+ 2D 8 * 3754.33 0.252 1.272e + 00 O HI 3754.69 V2 3s 3P* 3p 3D 3 5 3756.88 0.251 9.638e - 01 O HI 3757.24 V2 3s 3P* 3p 3D 1 3 3759.51 0.251 4.126e + 00 O HI 3759.87 V2 3s 3P* 3p 3D 5 7 3770.29 0.249 4.009e + 00 H 11 3770.63 Hll 2p+ 2P* lld+ 2D 8 * 3773.66 0.248 3.669e - 01 O HI 3774.02 V2 3s 3P* 3p 3D 3 3 3781.40 0.247 1.440e - 01 7 3781.40 * * 3791.02 0.245 3.657e - 01 O HI 3791.27 V2 3s 3P* 3p 3D 5 5 3797.58 0.244 5.439e + 00 H 10 3797.90 HIO 2p+ 2P* 10d+ 2D 8 * 3811.24 0.241 7.675e - 02 O VI 3811.35 * * 3813.24 0.241 2.470e - 01 He II 3813.50 4.19 4f+ 2F* 19g+ 2G 32 * 3819.38 0.240 5.843e - 01 He I 3819.62 V22 2p 3P* 6d 3D 9 15 3835.11 0.237 7.014e + 00 H 9 3835.39 H9 2p+ 2P* 9d+ 2D 8 * 3839.13 0.236 5.950e - 02 [Fe V] 3839.27 3d4 5D 3d4 3F2 7 7 3857.86 0.233 2.834e - 01 He II 3858.07 4.17 4f+ 2F* 17g+ 2G 32 * 3868.51 0.230 1.421e + 02 [Ne HI] 3868.75 FI 2p4 3P 2p4 ID 5 5 A.I. Galactic Planetary Nebulae 232

-^obs (Â) f(A) /(A) Ion Ao Mult Lower term Upper term gl g2 3888.36 0.226 4.917e + 00 He I 3888.65 V2 2s 3S 3p 3P* 3 9 3888.76 0.226 1.216e + 01 H 8 3889.05 * * 3923.29 0.220 3.811e - 01 He II 3923.48 4.15 4f+ 2F* 15g+ 2G 32 * 3967.23 0.211 6.014e + 01 [Ne HI] 3967.46 FI 2p4 3P 2p4 ID 3 5 4008.82 0.202 9.219e - 02 He I 4009.26 V55 2p IP* 7d ID 3 5 4025.71 0.198 1.886e + 00 He I 4026.21 V18 2p 3P* 5d 3D 9 15 4059.90 0.191 4.432e - 02 [FIV] 4060.23 2s2 2F 2s2 2 [4] 8 10 4067.64 0.189 2.380e - 01 C HI 4067.94 * * 4068.37 0.189 1.056e + 00 [S II] 4068.60 FI 2p3 4S* 2p3 2P* 4 4 4068.74 0.189 3.117e - 01 C III 4068.91 * * 4070.07 0.189 4.068e - 01 C HI 4070.26 * * 4071.93 0.188 1.697e - 01 O II 4072.16 VIO 3p 4D* 3d 4F 6 8 4075.63 0.187 2.420e - 01 O II 4075.86 VIO 3p 4D* 3d 4F 8 10 4076.12 0.187 3.192e - 01 [S II] 4076.35 FI 2p3 4S* 2p3 2P* 2 4 4078.61 0.187 6.256e - 02 O II 4078.84 ** 4083.46 0.186 6.945e - 02 O II 4083.90 V48b 3d 4F 4f G4* 6 8 4085.05 0.185 5.620e - 02 O II 4085.11 VIO 3p 4D* 3d 4F 6 6 4086.71 0.185 3.946e - 02 O II 4087.15 V48c 3d 4F 4f G3* 4 6 4088.84 0.184 9.842e - 02 O II 4089.29 V48a 3d 4F 4f G5* 10 12 4097.01 0.183 1.684e + 00 N HI 4097.33 VI 3s 2S 3p 2P* 2 4 4101.40 0.182 2.554e + 01 H 6 4101.74 H6 2p+ 2P* 6d4" 2D 8 72 4120.45 0.177 1.572e - 01 He I 4120.84 V16 2p 3P* 5s 3S 9 3 4122.14 0.177 6.557e - 02 O II 4121.46 V19 ** 4143.24 0.172 2.159e - 01 He I 4143.76 V53 2p IP* 6d ID 3 5 4185.06 0.163 5.233e - 02 O II 4185.45 ** 4186.51 0.162 1.184e-01 C III 4186.90 V18 4f IF* 5g IG 7 9 4195.91 0.160 2.977e - 02 N III 4195.76 V6 3s’ 2P* 3p’ 2D 2 4 4199.50 0.160 1.106e + 00 He II 4199.83 4.11 4f+ 2F* l lg + 2G 32 * 4219.16 0.155 5.025e - 02 Ne II 4219.74 V52a 3d 4D 4f 2|4^* 8 10 4227.32 0.153 2.818e - 01 N II 4227.74 V33 3p ID 4s IP* 5 3 4250.98 0.148 3.400e - 02 Ne II 4250.65 * * 4266.89 0.144 4.485e - 01 C II 4267.15 V6 3d 2D 4f 2F* 10 14 4340.13 0.127 4.686e + 01 H 5 4340.47 H5 2p+ 2P* 5d4" 2D 8 50 4362.87 0.121 2.323e + 01 [O HI] 4363.21 F2 2p2 ID 2p2 IS 5 1 4366.81 0.120 9.975e - 02 O II 4366.89 V2 3s 4P 3p 4P* 6 4 4368.89 0.120 6.994e - 02 Ne II 4369.86 * * 4378.87 0.118 2.028e - 01 N III 4379.11 V18 4f 2F* 5g 2G 14 18 4387.69 0.115 3.186e - 01 He I 4387.93 V51 2p IP* 5d ID 3 5 A.I. Galactic Planetary Nebulae 233

-^obs (Â) f(A) 7(A) Ion Ao Mult Lower term Upper term gl g2 4414.82 0.109 5.944e - 02 0 II 4414.90 V5 3s 2P 3p 2D* 4 6 4452.67 0.100 6.432e - 02 O II 4452.37 V5 3s 2P 3p 2D* 4 4 4471.20 0.095 2.583e -f 00 He I 4471.50 V14 2p 3P* 4d 3D 9 15 4510.69 0.086 9.996e - 02 N III 4510.91 V3 3s’ 4P* 3p’ 4D 4 6 4534.44 0.080 4.838e - 02 N III 4534.58 V3 3s’ 4P* 3p’ 4D 6 6 4541.35 0.078 2.091e + 00 He II 4541.59 4.9 41+ 2F* 9g+ 2G 32 * 4562.43 0.073 2.112e - 01 Mg I] 4562.60 3s2 IS 3s3p 3P* 1 5 4570.94 0.071 3.134e- 01 Mg I] 4571.10 3s2 IS 3s3p 3P* 1 3 4590.86 0.066 5.409e - 02 O II 4590.97 V15 3s’ 2D 3p’ 2F* 6 8 4606.52 0.062 4.603e - 02 N II 4607.16 V5 3s 3P* 3p 3P 1 3 4633.92 0.056 1.139e + 00 N HI 4634.14 V2 3p 2P* 3d 2D 2 4 4632.05 0.056 1.382e - 01 N HI 4634.14 * * 4638.64 0.055 2.160e - 01 O II 4638.86 VI 3s 4P 3p 4D* 2 4 4640.42 0.054 2.251e-b00 N HI 4640.64 V2 3p 2P* 3d 2D 4 6 4641.59 0.054 5.232e - 01 O II 4641.81 VI 3s 4P 3p 4D* 4 6 4647.21 0.053 1.873e - 01 C HI 4647.42 VI 3s 3S 3p 3P* 3 5 4648.92 0.052 1.687e - 01 O II 4649.13 VI 3s 4P 3p 4D* 6 8 4650.04 0.052 1.234e - 01 C HI 4650.25 VI 3s 3S 3p 3P* 3 3 4650.63 0.052 2.617e - 02 O II 4650.84 VI 3s 4P 3p 4D* 2 2 4651.26 0.052 7.851e - 02 C HI 4651.47 V I 3s 3S 3p 3P* 3 1 4657.29 0.050 1.573e - 01 ? 4657.29 * * 4658.45 0.050 1.8006 - 01 [Fe HI] 4658.10 F3 3d6 5D 3d6 3F2 9 9 4661.36 0.049 8.505e - 02 O II 4661.63 VI 3s 4P 3p 4D* 4 4 4676.20 0.046 2.417e - 02 O II 4676.24 VI 3s 4P 3p 4D* 6 6 4685.60 0.043 6.112e-b01 He II 4685.68 3.4 3d+ 2D 4f+ 2F* 18 32 4711.25 0.037 8.235e -H 00 [Ar IV] 4711.37 FI 3p3 4S* 3p3 2D* 4 6 0.037 8.3396 + 00 [Ar IV] 4711.37 FI 3p3 4S* 3p3 2D* 4 6 4713.85 0.036 8.2206 - 01 He I 4713.17 V12 2P 3p* 4S 3s 9 3 4724.04 0.034 6.7236 - 01 [Ne IV] 4724.15 FI 2p3 2D* 2p3 2P* 4 4 4725.51 0.033 5.3826 - 01 [Ne IV] 4725.62 FI 2p3 2D* 2p3 2P* 4 2 0.030 6.7876 + 00 [Ar IV] 4740.17 FI 3p3 4S* 3p3 2D* 4 4 4740.12 0.030 7.3356 + 00 [Ar IV] 4740.17 FI 3p3 4S* 3p3 2D* 4 4 4796.03 0.016 4.1726 - 02 7 4796.03 * * 4831.07 0.007 2.7226 - 02 ? 4831.07 * * 4861.31 0.000 1.0006 + 02 H 4 4861.33 H4 2p+ 2P* 4d+ 2D 8 32 4921.97 -0.015 6.8876 - 01 He I 4921.93 V48 2p IP* 4d ID 3 5 4931.48 -0.017 1.5706 - 01 [O HI] 4931.80 FI 2p2 3P 2p2 ID 1 5 4935.82 -0.018 1.1766 - 01 ? 4935.82 * * A.I. Galactic Planetary Nebulae 234

-^obs (Â) /(A) Ion Ao Mult Lower term Upper term gl g2 4958.92 -0.024 5.568e -f 02 [0 III] 4958.91 FI 2p2 3P 2p2 ID 3 5 -0.024 5.483e 4- 02 [O III] 4958.91 FI 2p2 3P 2p2 ID 3 5 5006.10 -0.036 1.545e-H03 [O III] 5006.84 FI 2p2 3P 2p2 ID 5 5 5190.37 -0.081 3.104e - 01 [Ar III] 5191.82 ** 5198.39 -0.083 3.217e - 01 [N I] 5199.84 * * 5410.75 -0.134 4.314e + 00 He II 5411.52 4.7 4f-h 2F* 7g+ 2 0 32 98 5517.02 -0.154 8.286e - 01 [Cl III] 5517.66 FI 2p3 4S* 2p3 2D* 4 6 5537.15 -0.158 8.676e - 01 [Cl III] 5537.60 FI 2p3 4S* 2p3 2D* 4 4 5592.33 -0.167 2.085e - 01 O III 5592.37 V5 3s IP* 3p IP 3 3 5753.99 -0.195 1.238e + 00 [N II] 5754.60 F3 2p2 ID 2p2 IS 5 1 5801.33 -0.202 5.615e - 02 C IV 5801.51 VI 3s 2S 3p 2P* 2 4 5875.16 -0.215 7.464e -t- 00 He I 5875.66 Vll 2p 3P* 3d 3D 9 15 6004.63 -0.236 1.083e - 01 He II 6004.73 5.22 5g+ 2 0 22h+ 2H* 50 * 6046.22 -0.242 2.766e - 01 7 6046.22 * * 6101.26 -0.251 3.500e - 01 [K IV] 6101.83 FI 3p4 3P 3d4 ID 5 5 6117.41 -0.254 7.587e - 02 He II 6118.20 5.19 5g4- 2 0 19h+ 2H* 50 * 6170.06 -0.262 1.278e - 01 He II 6170.60 5.18 5g+ 20 18h4- 2H 5 * 6233.58 -0.272 2.009e - 01 He II 6233.80 5.17 5g-f- 2 0 17h+ 2H 5 * 6300.23 -0.282 2.215e -K 00 [O I] 6300.34 FI 2p4 3P 2p4 ID 5 5 6311.85 -0.283 3.305c -|- 00 [S HI] 6312.10 F3 2p2 ID 2p2 IS 5 1 6363.78 -0.291 7.283e - 01 [O I] 6363.78 FI 2p4 3P 2p4 ID 3 5 6406.33 -0.297 2.749e - 01 He II 6406.30 5.15 5g-t- 2 0 15h+ 2H 5 * 6435.08 -0.302 7.200e - 01 [Ar V] 6435.10 3s2 3P 3s2 ID 3 5 6527.32 -0.315 3.335e - 01 He II 6527.11 5.14 5g+ 20 14h4- 2H 5 * 6548.22 -0.318 2.038e -f 01 [N II] 6548.10 FI 2p2 3P 2p2 ID 3 5 6562.79 -0.320 2.742e -4- 02 H 3 6562.77 H3 2p+ 2P* 3d4- 2D 8 18 6583.55 -0.323 6.257e 4- 01 [N II] 6583.50 FI 2p2 3P 2p2 ID 5 5 6678.45 -0.336 2.170c-k 00 He I 6678.16 V46 2p IP* 3d ID 3 5 6716.64 -0.342 5.301c-h 00 [S II] 6716.44 F2 2p3 4S* 2p3 2D* 4 6 6730.99 -0.344 7.071c 4- 00 [S II] 6730.82 F2 2p3 4S* 2p3 2D* 4 4 6825.87 -0.356 1.831c - 01 7 6825.87 * * 6890.92 -0.365 4.301c - 01 He II 6890.88 7 ** 7005.89 -0.380 1.369c 4-00 [ArV] 7005.67 3s2 3P 3s2 ID 5 5 7065.08 -0.387 1.876c + 00 He I 7065.25 VIO 2p 3P* 3s 3S 9 3 7135.66 -0.396 1.824c-h 01 [Ar HI] 7135.80 FI 3p4 3P 3p4 ID 5 5 7170.72 -0.401 2.625c - 01 [Ar IV] 7170.62 F2 3p3 2D* 3p3 2P* 4 4 7177.60 -0.401 5.205c - 01 He II 7177.50 5.11 5g+ 20 llh+ 2H 5 * 7236.26 -0.408 2.948c - 01 C II 7236.42 V3 3p 2P* 3d 2D 4 6 A.I. Galactic Planetary Nebulae 235

-^obs (Â) f(A) /(A) Ion Ao Mult Lower term Upper term g l g2 7262.79 -0.412 2.299e - 01 [Ar IV] 7262.76 F2 3p3 2D* 3p3 2P* 4 2 7281.29 -0.414 2.549e - 01 He I 7281.35 V45 2p IP* 3s IS 3 1 7319.74 -0.418 2.175e + 00 [0 II] 7318.92 F2 2p3 2D* 2p3 2P* 6 2 7329.94 -0.420 1.817e-kOO [O II] 7329.67 F2 2p3 2D* 2p3 2P* 4 2

c(H;0)(Bal) = 1.24

Table A.10: My Cn 18 linelist.

Aobs {N) f(A) /(A) Ion Ao Mult Lower term Upper term gl g2

* 0.257 * H 14 3721.94 H14 2p-b 2P* 14d-b 2D 8 * 0.257 * [S III] 3721.63 F2 3p2 3P 3p2 IS 3 1 3724.96 0.257 6.361e -b 01 [O II] 3726.03 FI 2p3 4S* 2p3 2D* 4 4 0.256 * [O II] 3728.82 FI 2p3 4S* 2p3 2D* 4 6 0.255 * H 13 3734.37 H13 2p+ 2P* 13d+ 2D 8 * 0.253 * H 12 3750.15 H12 2p-b 2P* 12d-b 2D 8 * 3833.28 0.237 4.198e-bOO H 9 3835.39 H9 2p-b 2P* 9d+ 2D 8 * 3866.32 0.230 1.057e-b01 [Ne III] 3868.75 FI 2p4 3P 2p4 ID 5 5 3886.40 0.226 1.303e-b01 He I 3888.65 V2 2s 3S 3p 3P* 3 9 3967.24 0.211 1.526e-b01 [Ne III] 3967.46 FI 2p4 3P 2p4 ID 3 5 0.202 * He I 4009.26 V55 2p IP* 7d ID 3 5 4025.33 0.198 2.039e -b 00 He I 4026.21 V18 2p 3P* 5d 3D 9 15 4040.83 * 5.148e - 02 N II 4040.83 * * 4056.61 0.191 1.422e - 01 ? 4056.61 * * 4067.73 0.189 2.238e -b 00 [S II] 4068.60 FI 2p3 4S* 2p3 2P* 4 4 0.189 * O II 4069.62 VIO 3p 4D* 3d 4F 2 4 0.189 * O II 4069.89 VIO 3p 4D* 3d 4F 4 6 0.188 * O II 4072.16 VIO 3p 4D* 3d 4F 6 8 0.187 * O II 4075.86 VIO 3p 4D* 3d 4F 8 10 4075.48 0.187 7.631e - 01 [S II] 4076.35 FI 2p3 4S* 2p3 2P* 2 4 0.186 * O II 4083.90 V48b 3d 4F 4f G4* 6 8 0.185 * O II 4085.11 VIO 3p 4D* 3d 4F 6 6 4087.91 0.185 1.895e - 01 O II 4087.15 V48c 3d 4F 4f G3* 4 6 0.184 * O II 4089.29 V48a 3d 4F 4f G5* 10 12 4100.84 0.182 2.123e-b 01 H 6 4101.74 H6 2p+ 2P* 6d+ 2D 8 72 4119.95 0.178 1.385e - 01 O II 4119.22 V20 3p 4P* 3d 4D 6 8 0.177 * He I 4120.84 V16 2p 3P* 5s 3S 9 3 4142.60 0.172 3.261e - 01 He I 4143.76 * * A.I. Galactic Planetary Nebulae 236

-^obs (•'^) f(A) /(A) Ion Ao Mult Lower term Upper term g l g2 4238.12 * 5.587e - 02 N II 4238.12 * * 0.150 * N II 4241.78 V48a 3d 3D* 4f lj3i 5 7 4266.22 0.144 2.524e - 01 C II 4267.15 V6 3d 2D 4f 2F* 10 14 4339.51 0.127 4.162e + 01 H 5 4340.47 H5 2p+ 2P* 5d-k 2D 8 50 0.125 * O II 4345.56 V2 3s 4P 3p 4P* 4 2 4362.22 0.121 5.707e - 01 [O III] 4363.21 F2 2p2 ID 2p2 IS 5 1 4386.93 0.115 5.412e - 01 He I 4387.93 V51 2p IP* 5d ID 3 5 0.115 * Ne II 4391.99 V55e 3d 4F 41 2j5i* 10 12 4470.54 0.095 5.497e -f 00 He I 4471.50 V14 2p 3P* 4d 3D 9 15 0.073 * Mg I] 4562.60 3s2 IS 3s3p 3P* 1 5 4570.24 0.071 8.830e - 02 Mg I] 4571.10 3s2 IS 3s3p 3P* 1 3 4606.37 0.062 8.417e - 02 N II 4607.16 * * 0.062 * O II 4609.44 V92a 3d 2D 4f F4* 6 8 0.062 * O II 4610.20 V92c 3d 2D 4f F2* 4 6 4620.60 0.059 7.307e - 02 N II 4621.39 * * 4629.75 0.057 2.099e - 01 N II 4630.54 * * 4638.77 0.055 1.081e - 01 0 II 4638.86 VI 3s 4P 3p 4D* 2 4 4640.55 0.054 2.396e - 01 N HI 4640.64 V2 3p 2P* 3d 2D 4 6 4641.72 0.054 1.169e - 01 O II 4641.44 * * 4648.09 0.052 2.294e - 01 O II 4649.13 VI 3s 4P 3p 4D* 6 8 4649.79 0.052 1.531e - 01 O II 4650.84 VI 3s 4P 3p 4D* 2 2 4657.32 0.050 6.180e - 01 [Fe HI] 4658.10 * * 0.049 * O II 4661.63 VI 3s 4P 3p 4D* 4 4 0.046 * O II 4676.24 VI 3s 4P 3p 4D* 6 6 4685.46 0.043 5.741e - 01 He II 4685.68 3.4 3d-t- 2D 4f+ 2F* 18 32 4700.83 0.039 1.836e - 01 [Fe HI] 4701.62 * * 0.037 * [Ar IV] 4711.37 FI 3p3 4S* 3p3 2D* 4 6 4712.38 0.036 6.454e — 01 He I 4713.17 V12 2P 3p* 4S 3s 9 3 4733.12 0.032 4.954e — 02 * * * 0.030 * [Ar IV] 4740.17 FI 3p3 4S* 3p3 2D* 4 4 4754.05 0.026 1.355e - 01 ** * 4860.61 0.000 100.* H 4 4861.33 H4 2p-t- 2P* 4d+ 2D 8 32 4880.22 -0.005 2.802e - 01 [Fe HI] 4881.11 * * 4905.77 -0.011 7.5356 — 02 ** * 4921.23 -0.015 1.463e + 00 He I 4921.93 * * 4923.96 -0.015 4.662e - 02 * * * 4930.21 -0.017 5.700e - 02 * * * 4958.19 -0.024 9.033e 4- 01 [O HI] 4958.91 * * A.I. Galactic Planetary Nebulae 237

-^obs (À) f(A) /(A) Ion Ao Mult Lower term Upper term gl g2 4956.96 -0.024 9.802e -h 01 [O III] 4958.91 * * 5004.88 -0.036 2.881e 4- 02 [O III] 5006.84 FI 2p2 3P 2p2 ID 5 5 5099.97 -0.058 1.977e + 01 ? * ** 5196.78 -0.082 7.384e - 01 7 * * * 5268.56 -0.099 4.616e - 01 ? * * * 5515.94 -0.154 3.160e - 01 [Cl III] 5517.66 FI 2p3 48* 2p3 2D* 4 6 5535.87 -0.158 6.047e - 01 [Cl III] 5537.60 FI 2p3 48* 2p3 2D* 4 4 5677.11 -0.182 3.561e - 01 N II 5679.56 * * 5752.91 -0.195 3.172e-H00 [N II] 5754.60 F3 2p2 ID 2p2 18 5 1 5873.93 -0.215 1.530e4-01 He I 5875.66 V ll 2p 3P* 3d 3D 9 15 6006.39 -0.236 2.298e - 01 7 ** * 6048.71 -0.243 4.399e - 01 ? * * * 6299.13 -0.282 1.572e -b 00 [O I] 6300.34 FI 2p4 3P 2p4 ID 5 5 6310.89 -0.283 1.354e-b00 [S III] 6312.10 F3 2p2 ID 2p2 18 5 1 6346.02 -0.289 1.171e - 01 Si II 6347.10 * * 6362.70 -0.291 4.916e - 01 [O I] 6363.78 FI 2p4 3P 2p4 ID 3 5 6547.09 -0.318 5.7936 -f- 01 [N II] 6548.10 FI 2p2 3P 2p2 ID 3 5 6561.75 -0.320 2.297e + 02 H 3 6562.77 H3 2p+ 2P* 3d+ 2D 8 18 6582.48 -0.323 1.7606 4-02 [N II] 6583.50 FI 2p2 3P 2p2 ID 5 5 6677.03 -0.336 3.7986 4- 00 He I 6678.16 V46 2p IP* 3d ID 3 5 6715.42 -0.342 2.8606 + 00 [S II] 6716.44 F2 2p3 48* 2p3 2D* 4 6 6729.80 -0.344 5.0476 4- 00 [S II] 6730.82 F2 2p3 48* 2p3 2D* 4 4 7063.98 -0.387 5.1126-b 00 He I 7065.25 VIO 2p 3P* 3s 38 9 3 7134.51 -0.396 1.1076-b 01 [Ar HI] 7135.80 FI 3p4 3P 3p4 ID 5 5 7176.65 -0.401 3.2346 - 01 He II 7177.50 * * 7230.02 -0.408 7.7176 - 02 C II 7231.32 * * 7235.12 -0.408 1.9336 - 01 C II 7236.42 * * 7279.85 -0.414 5.0806 - 01 He I 7281.35 V45 2p IP* 3s 18 3 1 7318.18 -0.418 3.6196-b 00 [O II] 7318.92 F2 2p3 2D* 2p3 2P* 6 2 7328.93 -0.420 3.0216 4- 00 [O II] 7329.67 F2 2p3 2D* 2p3 2P* 4 2

c(H/3)(Bal) = 0.26-0.44, c(H/3)(rad) = 0.300, c (H/3)^^33°/^2470 ^ 0.640

Table A .11: IC 4406 linelist.

•^obs {N) f(A) f(A) Ion Aq Mult Lower term Upper term gi g2

1549 0.260 1.9906 -b 01 C IV 1549 1640 0.260 8.3766 4- 01 He II 1640 1663 0.260 7.8006 4- 00 O HI] 1663 1751 0.260 7.8006 4- 00 N HI] 1750 A.I. Galactic Planetary Nebulae 238

-^obs (^) f(A) 7(A) Ion Ao Mult Lower term Upper term gl g2 1908 0.260 1.401e + 02 C III] 1908 2326 0.260 6.890e + 01 C II] 1908 2423 0.260 9.300e + 00 [Ne IV] 2423 2470 0.260 1.039e + 01 [O II] 1908

0.257 * H 14 3721.94 H14 2p+ 2P* 14d+ 2D 8 * 3721.14 0.257 * [S III] 3721.63 F2 3p2 3P 3p2 IS 3 1 3726.29 0.257 4.223e + 02 [O II] 3726.03 FI 2p3 4S* 2p3 2D* 4 4 3728.33 0.256 * [O II] 3728.82 FI 2p3 4S* 2p3 2D* 4 6 3733.88 0.255 * H 13 3734.37 H13 2p+ 2P* 13d+ 2D 8 * 3749.18 0.253 4.267e + 00 H 12 3750.15 H12 2p+ 2P* 12d+ 2D 8 * 3756.27 0.251 1.522e + 00 O III 3757.24 V2 3s 3P* 3p 3D 1 3 3758.90 0.251 2.187e + 00 O III 3759.87 V2 3s 3P* 3p 3D 5 7 3769.66 0.249 4.574e + 00 H 11 3770.63 Hll 2p+ 2P* lld+ 2D 8 * 3797.28 0.244 6.126e + 00 H 10 3797.90 HIO 2p+ 2P* 10d+ 2D 8 * 3835.34 0.237 7.774e + 00 H 9 3835.39 H9 2p+ 2P* 9d+ 2D 8 * 3867.94 0.230 1.290e + 02 [Ne III] 3868.75 FI 2p4 3P 2p4 ID 5 5 3888.24 0.226 2.292e + 01 He I 3888.65 V2 2s 3S 3p 3P* 3 9 3966.63 0.211 5.196e + 01 [Ne III] 3967.46 FI 2p4 3P 2p4 ID 3 5 4008.78 0.202 1.798e - 01 He I 4009.26 V55 2p IP* 7d ID 3 5 4025.88 0.198 2.589e + 00 He I 4026.21 V18 2p 3P* 5d 3D 9 15 4068.18 0.189 1.432e + 00 [S II] 4068.60 FI 2p3 4S* 2p3 2P* 4 4 4069.34 0.189 4.106e - 01 O II 4069.62 VIO 3p 4D* 3d 4F 2 4 0.189 * O II 4069.89 VIO 3p 4D* 3d 4F 4 6 4071.74 0.188 3.776e - 01 O II 4072.16 VIO 3p 4D* 3d 4F 6 8 4075.44 0.187 2.530e - 01 O II 4075.86 VIO 3p 4D* 3d 4F 8 10 4075.93 0.187 5.187e - 01 [S II] 4076.35 FI 2p3 4S* 2p3 2P* 2 4 0.186 * 0 II 4083.90 V48b 3d 4F 4f G4* 6 8 0.185 * O II 4085.11 VIO 3p 4D* 3d 4F 6 6 0.185 * O II 4087.15 V48c 3d 4F 4f G3* 4 6 0.184 * O II 4089.29 V48a 3d 4F 4f G5* 10 12 4096.97 0.183 9.972e - 01 N III 4097.33 VI 3s 2S 3p 2P* 2 4 4101.38 0.182 2.404e + 01 H 6 4101.74 H6 2p+ 2P* 6d+ 2D 8 72 0.178 * O II 4119.22 V20 3p 4P* 3d 4D 6 8 4120.68 0.177 2.717e - 01 He I 4120.84 V16 2p 3P* 5s 3S 9 3 4143.37 0.172 3.652e - 01 He I * * 4186.73 0.162 9.124e - 02 c m 4186.90 V18 4flF* 5glG 7 9 4199.56 0.160 3.556e - 01 He II 4199.83 4.11 4f+ 2F* l lg + 2G 32 * 4266.82 0.144 8.334e - 01 C II 4267.15 V6 3d 2D 4f 2F* 10 14 A.I. Galactic Planetary Nebulae 239

-^obs (Â) f(A) /(A) Ion Ao Mult Lower term Upper term gl g2 0.142 * O II 4275.55 V67a 3d 4D 41 F4* 8 10 0.142 * O II 4275.99 V67b 3d 4D 4f F3* 4 6 0.140 * 0 II 4282.96 V67c 3d 4D 4f F2* 4 6 0.140 * O II 4283.73 V67c 3d 4D 4f F2* 4 4 0.140 * O II 4285.69 V78b 3d 2F 4f F3* 6 8 0.139 * O II 4288.82 V53c 3d 4P 4f D l* 2 4 0.138 * O II 4291.25 V55 3d 4P 4f G3* 6 8 0.132 * O II 4319.63 V2 3s 4P 3p 4P* 4 6 4325.24 0.130 1.231e - 01 O II 4325.76 V2 3s 4P 3p 4P* 2 2 4340.09 0.127 4.571e-h01 H 5 4340.47 H5 2p+ 2P* 5d+ 2D 8 50 0.125 * O II 4345.56 V2 3s 4P 3p 4P* 4 2 4362.83 0.121 6.796e -H 00 [O III] 4363.21 F2 2p2 ID 2p2 IS 5 1 0.121 * [O III] 4363.21 F2 2p2 ID 2p2 IS 5 1 4378.89 0.118 1.717e - 01 N III 4379.11 V18 4f 2F* 5g 2G 14 18 4387.53 0.115 6.997e - 01 He I 4387.93 V51 2p IP* 5d ID 3 5 4391.23 0.115 * Ne II 4391.99 V55e 3d 4F 4f 2\5i* 10 12 4471.12 0.095 4.684e + 00 He I 4471.50 V14 2p 3P* 4d 3D 9 15 4541.25 0.079 4.941e - 01 ? 4541.25 ** 4562.26 0.073 3.175e - 01 Mg I] 4562.60 3s2 IS 3s3p 3P* 1 5 4570.75 0.071 3.205e - 01 Mg I] 4571.10 3s2 IS 3s3p 3P* 1 3 4631.29 0.056 2.318e - 01 N II 4630.54 V5 3s 3P* 3p 3P 5 5 4633.83 0.056 4.500e - 01 N HI 4634.14 V2 3p 2P* 3d 2D 2 4 4638.55 0.055 1.817e - 01 O II 4638.86 VI 3s 4P 3p 4D* 2 4 4640.33 0.054 8.048e - 01 N HI 4640.64 V2 3p 2P* 3d 2D 4 6 4641.50 0.054 3.527e - 01 O II 4641.44 * * 4647.08 0.053 1.137e - 01 C HI 4647.42 VI 3s 3S 3p 3P* 3 5 4648.79 0.052 2.585e - 01 O II 4649.13 VI 3s 4P 3p 4D* 6 8 4649.91 0.052 1.017e - 01 C HI 4650.25 * * 4650.50 0.052 9.590e - 02 O II 4650.84 VI 3s 4P 3p 4D* 2 2 4651.13 0.052 2.219e - 02 C III 4651.47 ** 4661.37 0.049 9.884e - 02 O II 4661.63 VI 3s 4P 3p 4D* 4 4 4675.89 0.046 4.808e - 02 O II 4676.24 VI 3s 4P 3p 4D* 6 6 4685.43 0.043 1.467e + 01 He II 4685.68 3.4 3d-|- 2D 4f+ 2F* 18 32 4711.13 0.037 1.454e -h 00 [Ar IV] 4711.37 FI 3p3 4S* 3p3 2D* 4 6 4712.93 0.036 5.335e - 01 He I 4713.17 V12 2P 3p* 4S 3s 9 3 4739.92 0.030 1.186e + 00 [Ar IV] 4740.17 FI 3p3 4S* 3p3 2D* 4 4 4861.07 0.000 l.OOOe -t- 02 H 4 4861.33 H4 2p+ 2P* 4d+ 2D 8 32 4921.68 -0.015 1.341e-t-00 He I 4921.93 * * A.I. Galactic Planetary Nebulae 240

-^obs (Â) f(A) /(A) Ion Ao Mult Lower term Upper term gl g2 4931.02 -0.017 1.436e - 01 O III 4931.80 ** 4934.52 -0.018 1.081e - 01 ? 4934.52 * * 4958.66 -0.024 3.617e + 02 [O III] 4958.91 ** 4958.16 -0.024 3.640e -b 02 [O III] 4958.91 ** 5006.09 -0.036 1.046e -b 03 [O III] 5006.84 FI 2p2 3P 2p2 ID 5 5 5198.10 -0.083 6.501e -b 00 [NI] 5199.84 * * 5410.55 -0.134 1.051e-b 00 He II 5411.52 4.7 4f+ 2F* 7g+ 2 0 32 98 5517.28 -0.154 5.246e - 01 [Cl III] 5517.66 FI 2p3 4S* 2p3 2D* 4 6 5537.22 -0.158 6.406e — 01 [Cl III] 5537.60 FI 2p3 4S* 2p3 2D* 4 4 5576.96 -0.164 3.703e - 01 [OI] 5577.34 * * 5753.81 -0.195 5.778e + 00 [N II] 5754.60 F3 2p2 ID 2p2 IS 5 1 5875.01 -0.215 1.4096 -b 01 He I 5875.66 Vll 2p 3P* 3d 3D 9 15 6299.99 -0.282 3.1966 -b 01 [0 1] 6300.34 FI 2p4 3P 2p4 ID 5 5 6311.75 -0.283 6.7216 - 01 [S III] 6312.10 F3 2p2 ID 2p2 IS 5 1 6363.41 -0.291 1.5586-b 01 [OI] 6363.78 FI 2p4 3P 2p4 ID 3 5 6547.95 -0.318 1.2606-b 02 [N II] 6548.10 FI 2p2 3P 2p2 ID 3 5 6562.62 -0.320 2.6636 + 02 H 3 6562.77 H3 2p+ 2P* 3d+ 2D 8 18 6583.35 -0.323 3.7646 -b 02 [N II] 6583.50 FI 2p2 3P 2p2 ID 5 5 6678.03 -0.336 3.7336 + 00 He I 6678.16 V46 2p IP* 3d ID 3 5 6716.41 -0.342 8.2156 + 00 [S II] 6716.44 F2 2p3 4S* 2p3 2D* 4 6 6730.79 -0.344 9.2736 + 00 [S II] 6730.82 F2 2p3 4S* 2p3 2D* 4 4 7064.96 -0.387 3.1226 + 00 He I 7065.25 VIO 2p 3P* 3s 3S 9 3 7135.47 -0.396 1.9716 + 01 [Ar III] 7135.80 FI 3p4 3P 3p4 ID 5 5 7281.53 -0.414 5.2216 - 01 He I 7281.35 V45 2p IP* 3s IS 3 1 7319.11 -0.418 6.4216 + 00 [O II] 7318.92 F2 2p3 2D* 2p3 2P* 6 2 7329.86 -0.420 5.1616 + 00 [0 II] 7329.67 F2 2p3 2D* 2p3 2P* 4 2

c(H/3)(Bal) = 0.610

Table A.12: IC4191 linelist (fixed-slit observations).

-^obs (Â) f(A) /(A) Ion Ao Mult Lower term Upper term gl g2

3704.18 0.260 1.034e + 00 He I 3705.02 ** 3711.95 0.259 8.638e - 01 H 15 3711.97 * * 0.257 * H 14 3721.94 H14 2p+ 2P* 14d+ 2D 8 * 0.257 * [S HI] 3721.63 F2 3p2 3P 3p2 IS 3 1 3726.19 0.257 4.655e + 01 [O II] 3726.03 FI 2p3 4S* 2p3 2D* 4 4 0.256 * [O II] 3728.82 FI 2p3 4S* 2p3 2D* 4 6 0.255 * H 13 3734.37 H13 2p+ 2P* 13d+ 2D 8 * 3749.91 0.253 2.944e + 00 H 12 3750.15 H12 2p+ 2P* 12d+ 2D 8 * A.I. Galactic Planetary Nebulae 241

■^obs (Â) f(A) /(A) Ion Ao Mult Lower term Upper term gl g2 3757.00 0.251 6.168e - 01 0 III 3757.24 V2 3s 3P* 3p 3D 1 3 3759.63 0.251 1.632e + 00 O III 3759.87 V2 3s 3P* 3p 3D 5 7 3770.39 0.249 3.362e + 00 H 11 3770.63 H ll 2p+ 2P* l l d + 2D 8 * 3797.47 0.244 4.739e + 00 H 10 3797.90 HIO 2p+ 2P* 10d+ 2D 8 * 3819.19 0.240 1.251e + 00 He I 3819.62 * * 3834.96 0.237 6.720e -t- 00 H 9 3835.39 H9 2p+ 2P* 9d+ 2D 8 * 3868.41 0.230 1.309e + 02 [Ne III] 3868.75 FI 2p4 3P 2p4 ID 5 5 3888.71 0.226 1.661e + 01 He I 3888.65 V2 2s 3S 3p 3P* 3 9 3967.11 0.211 5.185e + 01 [Ne HI] 3967.46 FI 2p4 3P 2p4 ID 3 5 4009.04 0.202 2.105e - 01 He I 4009.26 V55 2p IP* 7d ID 3 5 4025.97 0.198 2.654e + 00 He I 4026.21 V18 2p 3P* 5d 3D 9 15 4068.35 0.189 4.121e + 00 [S II] 4068.60 FI 2p3 4S* 2p3 2P* 4 4 4069.51 0.189 5.998e - 01 O II 4069.62 VIO 3p 4D* 3d 4F 2 4 0.189 * O II 4069.89 VIO 3p 4D* 3d 4F 4 6 4071.91 0.188 3.565e - 01 O II 4072.16 VIO 3p 4D* 3d 4F 6 8 4075.61 0.187 3.452e - 01 O II 4075.86 VIO 3p 4D* 3d 4F 8 10 4076.10 0.187 1.486e + 00 [S II] 4076.35 FI 2p3 4S* 2p3 2P* 2 4 4083.71 0.186 7.007e - 02 O II 4083.90 V48b 3d 4F 4f G4* 6 8 4084.92 0.185 2.658e - 02 O II 4085.11 VIO 3p 4D* 3d 4F 6 6 4086.96 0.185 6.842e - 02 O II 4087.15 V48c 3d 4F 4f G3* 4 6 4089.10 0.184 2.061e - 01 O II 4089.29 V48a 3d 4F 4f G5* 10 12 4097.11 0.183 1.649e + 00 N HI 4097.33 VI 3s 2S 3p 2P* 2 4 4101.52 0.182 2.480e + 01 H 6 4101.74 H6 2p+ 2P* 6d+ 2D 8 72 4119.16 0.178 1.564e - 01 O II 4119.22 V20 3p 4P* 3d 4D 6 8 4120.78 0.177 2.575e - 01 He I 4120.84 V16 2p 3P* 5s 3S 9 3 4143.48 0.172 3.713e - 01 He I 4143.76 ** 4185.26 0.163 6.316e - 02 0 II 4185.45 * * 4186.71 0.162 9.864e - 02 C HI 4186.90 V18 4flF* 5glG 7 9 4189.60 0.162 9.228e - 02 O II 4189.79 * * 4199.73 0.160 2.950e - 01 He II 4199.83 4.11 4f+ 2F* llg + 2G 32 * 4219.75 0.155 6.878e - 02 Ne II 4219.74 * * 4267.00 0.144 5.693e - 01 C II 4267.15 V6 3d 2D 4f 2F* 10 14 4275.40 0.142 1.041e - 01 O II 4275.55 V67a 3d 4D 4f F4* 8 10 0.142 * 0 II 4275.99 V67b 3d 4D 4f F3* 4 6 4276.60 0.142 9.630e - 02 O II 4276.75 * * 0.140 * O II 4282.96 V67c 3d 4D 4f F2* 4 6 0.140 * O II 4283.73 V67c 3d 4D 4f F2* 4 4

0.140 * O II 4285.69 V78b 3d 2F 4f F3* 6 8 A.I. Galactic Planetary Nebulae 242

-^obs (Â) f(A) /(A) Ion Ao Mult Lower term Upper term gl g2 0.139 * O II 4288.82 V53c 3d 4P 4f D l* 2 4 0.138 * O II 4291.25 V55 3d 4P 4f G3* 6 8 0.137 * O II 4294.78 V53b 3d 4P 4f D2* 4 6 0.135 * O II 4303.61 V65a 3d 4D 4f G5* 8 10 4303.63 0.135 6.839e - 02 0 II 4303.82 ** 0.133 * O II 4312.11 V78a 3d 2F 4f F4* 8 8 0.133 * O II 4313.44 V78a 3d 2F 41 F4* 8 10 4317.01 0.132 6.755e - 02 O II 4317.14 V2 3s 4P 3p 4P* 2 4 4319.50 0.132 5.508e - 02 O II 4319.63 V2 3s 4P 3p 4P* 4 6 0.130 * O II 4325.76 V2 3s 4P 3p 4P* 2 2 4340.27 0.127 4.691e + 01 H 5 4340.47 H5 2p+ 2P* 5d+ 2D 8 50 4345.40 0.125 1.210e - 01 O II 4345.56 V2 3s 4P 3p 4P* 4 2 4349.27 0.125 2.324e - 01 O II 4349.43 V2 3s 4P 3p 4P* 6 6 0.124 * O II 4353.59 V76c 3d 2F 4f G3* 6 8 4363.01 0.121 1.084e + 01 [O III] 4363.21 F2 2p2 ID 2p2 IS 5 1 0.121 * [O III] 4363.21 F2 2p2 ID 2p2 IS 5 1 4366.78 0.120 1.064e - 01 O II 4366.89 V2 3s 4P 3p 4P* 6 4 0.119 * 0 II 4371.62 V76b 3d 2F 4f G4* 8 10 4378.94 0.118 2.769e - 01 N III 4379.11 V18 4f 2F* 5g 2G 14 18 4387.76 0.115 5.599e - 01 He I 4387.93 V51 2p IP* 5d ID 3 5 4391.63 0.115 5.527e - 02 Ne II 4391.99 V55e 3d 4F 4f 2j5i* 10 12 0.113 * Ne II 4397.99 V57b 3d 4F 4f lj4i* 6 8 0.110 * Ne II 4409.30 V55e 3d 4F 4f 2j5^* 8 10 4414.81 0.109 5.461e - 02 O II 4414.90 V5 3s 2P 3p 2D* 4 6 4416.88 0.109 4.438e - 02 O II 4416.97 V5 3s 2P 3p 2D* 2 4 0.106 * Ne II 4428.64 V60c 3d 2F 4f 1[3]* 6 8 4437.40 0.103 4.609e - 02 He I 4437.55 V50 2p IP* 5s IS 3 1 0.101 * O II 4448.19 V35 3p’ 2F* 3d’ 2F 8 8 4465.51 0.097 3.869e - 02 7 4465.51 * * 4471.33 0.095 5.019e + 00 He I 4471.50 V14 2p 3P* 4d 3D 9 15 0.093 * Mg II 4481.21 V4 3d 2D 41 2F* 10 14 0.091 * O II 4488.20 V104 3d’ 2P 4f’ D2* 4 6 0.091 * C II 4491.07 4f 2F* 9g 2G 14 18 0.091 * O II 4491.23 V86a 3d 2P 4f D3* 4 6 4510.78 0.086 1.228e - 01 N III 4510.91 ** 4541.46 0.078 4.816e-01 He II 4541.59 * * 0.078 * N III 4544.80 * * 0.076 * N II 4552.53 V58a 3d IF* 4f 2\4i 7 9 A.I. Galactic Planetary Nebulae 243

-^obs (Â) f(A) /(A) Ion Ao Mult Lower term Upper term gl g2 0.073 * Mg I] 4562.60 3s2 IS 3s3p 3P* 1 5 4570.97 0.071 1.510e- 01 Mg I] 4571.10 3s2 IS 3s3p 3P* 1 3 4590.82 0.066 8.042e - 02 O II 4590.97 * * 4596.03 0.065 6.177e - 02 0 II 4596.18 V15 3s’ 2D 3p’2F* 4 6 0.064 * N II 4601.48 V5 3s 3P* 3p 3P 3 5 4609.75 0.062 6.032e - 02 O II 4609.44 V92a 3d 2D 4f F4* 6 8 0.062 * O II 4610.20 V92c 3d 2D 4f F2* 4 6 4631.93 0.056 8.157e - 02 7 4631.93 * * 4634.03 0.056 l.OlOe-bOO N III 4634.14 V2 3p 2P* 3d 2D 2 4 4638.75 0.055 2.426e - 01 O II 4638.86 VI 3s 4P 3p 4D* 2 4 4640.53 0.054 1.970e + 00 N III 4640.64 V2 3p 2P* 3d 2D 4 6 4641.70 0.054 5.282e - 01 O II 4641.44 ** 4647.30 0.053 1.729e - 01 C III 4647.42 VI 3s 3S 3p 3P* 3 5 4649.01 0.052 5.757e - 01 O II 4649.13 VI 3s 4P 3p 4D* 6 8 4650.13 0.052 1.221e - 01 C III 4650.25 * * 4650.72 0.052 1.168e - 01 0 II 4650.84 VI 3s 4P 3p 4D* 2 2 4651.35 0.052 1.554e - 02 C III 4651.47 ** 4658.06 0.050 8.454e - 02 [Fe III] 4658.10 ** 4661.59 0.049 1.442e - 01 O II 4661.63 VI 3s 4P 3p 4D* 4 4 4676.20 0.046 1.084e - 01 O II 4676.24 VI 3s 4P 3p 4D* 6 6 4685.64 0.043 1.452e + 01 He II 4685.68 3.4 3d4- 2D 4f+ 2F* 18 32 0.038 * O II 4705.35 * * 4711.32 0.037 3.128e-kOO [Ar IV] 4711.37 FI 3p3 4S* 3p3 2D* 4 6 4713.12 0.036 9.902e - 01 He I 4713.17 V12 2P 3p* 4S 3s 9 3 4724.14 0.034 1.931e - 01 [Ne IV] 4724.15 ** 4725.61 0.033 1.533e - 01 [Ne IV] 4725.62 * * 4740.17 0.030 5.614e + 00 [Ar IV] 4740.17 FI 3p3 4S* 3p3 2D* 4 4 4861.30 0.000 * H 4 4861.33 H4 2p-f- 2P* Ad+ 2D 8 32 -0.011 * O II 4906.83 * * 4921.95 -0.015 1.403e-H00 He I 4921.93 * * 4924.55 -0.016 4.662e - 02 0 II 4924.53 ** 4931.29 -0.017 1.811e - 01 [O III] 4931.80 * * -0.018 * ? 4934.52 * * 4958.88 -0.024 4.402e + 02 [O HI] 4958.91 ** 4958.20 -0.024 4.844e + 02 [O III] 4958.91 * * 5006.12 -0.036 1.501e + 03 [O III] 5006.84 FI 2p2 3P 2p2 ID 5 5 5041.04 -0.044 1.644e - 01 Si II 5041.03 * * 5047.73 -0.046 1.396e - 01 He I 5047.74 ** A.I. Galactic Planetary Nebulae 244

-^obs (-^) f(A) /(A) Ion Ao Mult Lower term Upper term gl g2 5055.99 -0.048 1.758e - 01 Si II 5055.98 * * 5190.21 -0.081 1.559e - 01 [Ar III] 5191.82 **

5198.23 -0.083 4.107e - 01 [NI] 5199.84 * * 5343.41 -0.117 3.215e - 02 C II 5342.38 * * 5411.18 -0.134 1.230e 4- 00 He II 5411.52 4.7 4f-k 2F* 7g+ 2 0 32 98 5517.43 -0.154 4.077e - 01 [Cl III] 5517.66 FI 2p3 4S* 2p3 2D* 4 6 5537.37 -0.158 8.547e - 01 [Cl III] 5537.60 FI 2p3 4S* 2p3 2D* 4 4 5577.11 -0.164 1.323e - 01 [0 I] 5577.34 * * 5591.89 -0.167 9.776e - 02 O III 5592.37 * * 5666.14 -0.180 9.206e - 02 N II 5666.63 * * 5679.07 -0.182 1.682e - 01 N II 5679.56 * * 5701.48 -0.185 6.311e - 02 7 * * 5754.43 -0.195 2.139e-H00 [N II] 5754.60 F3 2p2 ID 2p2 IS 5 1 5801.53 -0.202 9.141e - 02 C IV 5801.51 * * 5812.16 -0.204 8.365e - 02 C IV 5812.14 ** 5861.49 -0.212 6.331e - 02 ? * * 5875.49 -0.215 1.672e -K 01 He I 5875.66 V ll 2p 3P* 3d 3D 9 15 5930.58 -0.224 5.935e - 02 He II 5930.58 ** 5953.67 -0.227 4.524e - 02 He II 5953.67 * * 5977.58 -0.231 3.473e - 02 He II 5977.58 * * 6004.34 -0.235 1.302e - 01 7 * * 6045.58 -0.242 4.680e - 01 7 * * 6074.35 -0.247 5.390e - 02 7 * * 6083.53 -0.248 4.466e — 02 7 * * 6101.77 -0.251 3.133e - 01 [KIV] 6101.83 * * 6117.67 -0.254 4.358e - 02 He II 6118.20 * * 6151.97 -0.259 3.134e - 02 C II 6151.43 * * 6158.15 -0.260 3.681e - 02 [Mn V] 6157.60 * * 6170.35 -0.262 5.396e - 02 *** 6234.02 -0.272 6.720e - 02 He II 6233.80 * * 6300.39 -0.282 8.033e + 00 [OI] 6300.34 FI 2p4 3P 2p4 ID 5 5 6312.15 -0.283 2.651e -t- 00 [S HI] 6312.10 * * 6347.23 -0.289 1 .8 4 5 e - 01 Si II 6347.10 * * 6363.91 -0.291 2.706e -t- 00 [O I] 6363.78 FI 2p4 3P 2p4 ID 3 5 6371.51 -0.292 1.483e - 01 Si II 6371.38 * * 6393.98 -0.296 7.782e - 02 [Mn V] 6393.60 * * 6401.95 -0.297 6.317e - 02 ? 6401.95 * * 6406.68 -0.297 8.332e - 02 He II 6406.30 ** A.I. Galactic Planetary Nebulae 245

^obs (À) f(A) /(A) Ion Ao Mult Lower term Upper term gl g2 6435.11 -0.302 3.538e - 01 [Ar V] 6435.10 * 6461.96 -0.306 8.224e - 02 C II 6461.95 ** 6527.29 -0.315 1.346e - 01 He II 6527.11 * * 6547.97 -0.318 2.439e 4- 01 [N II] 6548.10 FI 2p2 3P 2p2 ID 3 5 6562.64 -0.320 2.763e + 02 H 3 6562.77 H3 2p+ 2P* 3d+ 2D 8 18

6583.37 -0.323 7.559e -|- 01 [N II] 6583.50 FI 2p2 3P 2p2 ID 5 5 6678.24 -0.336 4.553e -f 00 He I 6678.16 V46 2p IP* 3d ID 3 5 6716.52 -0.342 3.775e 4- 00 [S II] 6716.44 F2 2p3 4S* 2p3 2D* 4 6 6730.91 -0.344 7.497e -f 00 [S II] 6730.82 F2 2p3 4S* 2p3 2D* 4 4 6794.84 -0.352 6.015e - 02 [K IV] 6795.00 * * 6891.13 -0.365 1.436e - 01 ** * 7005.74 -0.380 6.639e - 01 [Ar V] 7005.74 * * 7109.63 -0.387 3.940e - 02 He I 7065.25 VIO 2p 3P* 3s 3S 9 3 * * * * * 7135.79 -0.396 1.899e -t- 01 [Ar III] 7135.80 FI 3p4 3P 3p4 ID 5 5 7170.70 -0.401 1.756e - 01 [Ar IV] 7170.62 * * 7177.58 -0.401 1.594e - 01 He II 7177.50 * * 7231.56 -0.408 7.401e - 02 ? 7231.56 * * 7236.60 -0.408 2.704e - 01 C II 7236.42 * * 7262.94 -0.468 1.236e - 01 [Ar IV] 7762.76 * * 7281.53 -0.414 6.664e - 01 He I 7281.35 V45 2p IP* 3s IS 3 1 7297.49 -0.416 3.082e - 02 * * * 7319.82 -0.418 5.798e -t 00 [O II] 7318.92 F2 2p3 2D* 2p3 2P* 6 2 7330.57 -0.420 4.787e + 00 [O II] 7329.67 F2 2p3 2D* 2p3 2P* 4 2

c(H/3)(Bal) = 0.580

T able A. 13: IC4191 linelist (scanning-slit observations).

Aobs (Â) f(A) /(A) Ion Ao Mult Lower term Upper term gl g2

* 0.257 * H 14 3721.94 H14 2p+ 2P* 14d+ 2D 8 * 0.257 * [S HI] 3721.63 F2 3p2 3P 3p2 IS 3 1 3726.10 0.257 5.917e-h01 [0 II] 3726.03 FI 2p3 4S* 2p3 2D* 4 4 0.256 * [O II] 3728.82 FI 2p3 4S* 2p3 2D* 4 6 0.255 * H 13 3734.37 H13 2p4- 2P* 13d+ 2D 8 * 3749.66 0.253 3.417e 4- 00 H 12 3750.15 H12 2p4- 2P* 12d4- 2D 8 * 3756.75 0.251 6.356e - 01 O III 3757.24 V2 3s 3P* 3p 3D 1 3 3759.38 0.251 1.782e 4- 00 O III 3759.87 V2 3s 3P* 3p 3D 5 7 3770.14 0.249 3.931e -f 00 H 11 3770.63 H ll 2p+ 2P* Hd4- 2D 8 * 3796.95 0.244 5.462e -F 00 H 10 3797.90 HIO 2p+ 2P* lOd-t- 2D 8 * A.I. Galactic Planetary Nebulae 246

■^obs (Â) f(A) /(A) Ion Aq Mult Lower term Upper term gl g2 3818.67 0.240 1.504e + 00 He I 3819.62 * * 3834.43 0.237 7.409e + 00 H 9 3835.39 H9 2p+ 2P* 9d+ 2D 8 * 3867.88 0.230 1.380e + 02 [Ne III] 3868.75 FI 2p4 3P 2p4 ID 5 5 3888.18 0.226 1.690e + 01 He I 3888.65 V2 2s 3S 3p 3P* 3 9 3966.57 0.211 5.409e + 01 [Ne HI] 3967.46 FI 2p4 3P 2p4 ID 3 5 4008.59 0.202 2.078e - 01 He I 4009.26 V55 2p IP* 7d ID 3 5 4025.59 0.198 2.460e + 00 He I 4026.21 V18 2p 3P* 5d 3D 9 15 4067.88 0.189 3.963e + 00 [S II] 4068.60 FI 2p3 4S* 2p3 2P* 4 4 0.189 * O II 4069.89 VIO 3p 4D* 3d 4F 4 6 4069.04 0.189 7.056e - 01 O II 4069.78 * * 4071.44 0.188 3.654e - 01 O II 4072.16 VIO 3p 4D* 3d 4F 6 8 4075.14 0.187 2.883e - 01 O II 4075.86 VIO 3p 4D* 3d 4F 8 10 4075.63 0.187 1.405e + 00 [S II] 4076.35 FI 2p3 4S* 2p3 2P* 2 4 4083.14 0.186 4.173e - 0 2 O II 4083.90 V48b 3d 4F 4f G4* 6 8 4084.35 0.185 5.0286 — 02 O II 4085.11 VIO 3p 4D* 3d 4F 6 6 4086.39 0.185 5.796e - 02 O II 4087.15 V48c 3d 4F 4f G3* 4 6 4088.52 0.184 1.476e - 01 O II 4089.29 V48a 3d 4F 4f G5* 10 12 4096.64 0.183 1.5416 + 00 N HI 4097.33 VI 3s 2S 3p 2P* 2 4 4101.05 0.182 2.4586 + 01 H 6 4101.74 H6 2p+ 2P* 6d+ 2D 8 72 4118.70 0.178 1.4476 - 01 O II 4119.22 V20 3p 4P* 3d 4D 6 8 4120.32 0.177 2.4806 - 01 He I 4120.84 V16 2p 3P* 5s 3S 9 3 4143.07 0.172 3.9976 - 01 He I 4143.76 ** 4184.56 0.163 4.4806 — 02 0 II 4185.45 ** 4186.01 0.162 7.1126 - 02 C III 4186.90 V18 4flF* 5glG 7 9 4188.90 0.162 7.092e - 02 O II 4189.79 ** 4199.27 0.160 2.7666 - 01 He II 4199.83 4.11 4f+ 2F* l lg + 2G 32 * 4266.49 0.144 5.4626 - 01 C II 4267.15 V6 3d 2D 4f 2F* 10 14 4274.89 0.142 7.2386 - 02 0 II 4275.55 V67a 3d 4D 4f F4* 8 10 0.142 * O II 4275.99 V67b 3d 4D 4f F3* 4 6 4276.09 0.142 9.274e - 02 0 II 4276.75 * * 4282.49 0.140 5.8736 - 02 0 II 4282.96 V67c 3d 4D 4f F2* 4 6 4284.45 0.140 7.9476 - 02 O II 4283.73 V67c 3d 4D 41 F2* 4 4 0.140 * O II 4285.69 V78b 3d 2F 4f F3* 6 8 4303.26 0.135 9.2966 - 02 0 II 4303.61 V65a 3d 4D 4f G5* 8 10 0.132 8.2696 - 02 O II 4317.14 V2 3s 4P 3p 4P* 2 4 0.132 * O II 4319.63 V2 3s 4P 3p 4P* 4 6 4339.73 0.127 4.6846 + 01 H 5 4340.47 H5 2p+ 2P* 5d + 2D 8 50 0.125 * O II 4345.56 V2 3s 4P 3p 4P* 4 2 A.I. Galactic Planetary Nebulae 247

-^obs (Â) f(A) /(A) Ion Ao Mult Lower term Upper term gl g2 4348.84 0.125 1.663e - 01 O II 4349.43 V2 3s 4P 3p 4P* 6 6 0.124 * O II 4353.59 V76c 3d 2F 4f G3* 6 8 4362.46 0.121 1.014e + 01 [O III] 4363.21 F2 2p2 ID 2p2 IS 5 1 0.120 * O II 4366.89 V2 3s 4P 3p 4P* 6 4 4378.43 0.118 2.701e - 01 N III 4379.11 V18 4f 2F* 5g 2G 14 18 4387.25 0.115 5.701e - 01 He I 4387.93 V51 2p IP* 5d ID 3 5 4391.03 0.115 4.316e - 02 Ne II 4391.99 V55e 3d 4F 4f 2j5i* 10 12 4414.17 0.109 7.566e - 02 O II 4414.90 V5 3s 2P 3p 2D* 4 6 4416.24 0.109 5.307e - 02 O II 4416.97 V5 3s 2P 3p 2D* 2 4 4436.83 0.103 4.697e - 02 He I 4437.55 V50 2p IP* 5s IS 3 1 4470.78 0.095 5.026e + 00 He I 4471.50 V14 2p 3P* 4d 3D 9 15 4510.22 0.086 1.207e - 01 N HI 4510.91 * * 4540.90 0.078 4.452e - 01 He II 4541.59 * * 0.073 * Mg I] 4562.60 3s2 IS 3s3p 3P* 1 5 4570.40 0.071 1.630e - 01 Mg I] 4571.10 3s2 IS 3s3p 3P* 1 3 4590.35 0.066 7.579e - 02 0 II 4590.97 * * 4595.56 0.065 5.858e - 02 O II 4596.18 V15 3s’ 2D 3p’2F* 4 6 4608.91 0.062 8.519e - 02 O II 4609.44 V92a 3d 2D 4f F4* 6 8 0.062 * O II 4610.20 V92c 3d 2D 4f F2* 4 6 4630.10 0.057 5.494e - 02 N II 4630.54 V5 3s 3P* 3p 3P 5 5 4633.47 0.056 9.185e - 01 N HI 4634.14 V2 3p 2P* 3d 2D 2 4 4638.19 0.055 1.912e - 01 O II 4638.86 VI 3s 4P 3p 4D* 2 4 4639.97 0.054 1.870e + 00 N HI 4640.64 V2 3p 2P* 3d 2D 4 6 4641.14 0.054 5.055e — 01 0 II 4641.44 * * 4646.76 0.053 1.568e - 01 C HI 4647.42 VI 3s 3S 3p 3P* 3 5 4648.47 0.052 5.439e - 01 O II 4649.13 VI 3s 4P 3p 4D* 6 8 4649.59 0.052 1.283e - 01 C HI 4650.25 * * 4650.18 0.052 9.970e - 02 O II 4650.84 VI 3s 4P 3p 4D* 2 2 4650.81 0.052 7.409e - 02 C HI 4651.47 * * 4657.49 0.050 1.061e - 0 1 [Fe HI] 4658.10 * * 4661.02 0.049 1.539e - 01 O II 4661.63 V I 3s 4P 3p 4D* 4 4 4675.63 0.046 1.226e - 01 0 II 4676.24 VI 3s 4P 3p 4D* 6 6 4685.07 0.043 1.322e + 01 He II 4685.68 3.4 3d+ 2D 4f+ 2F* 18 32 4704.93 0.038 1.766e - 02 O II 4705.35 ** 4710.74 0.037 2.929e + 00 [Ar IV] 4711.37 FI 3p3 4S* 3p3 2D* 4 6 4712.54 0.036 8.365e - 01 He I 4713.17 V12 2P 3p* 4S 3s 9 3 4723.61 0.034 1.703e - 01 [Ne IV] 4724.15 ** 4725.08 0.033 1.253e - 01 [Ne IV] 4725.62 * * A.I. Galactic Planetary Nebulae 248

-^obs (Â) f(A) /(A ) Ion Ao Mult Lower term Upper term gl g2 4739.61 0.030 5.074e -1- 00 [Ar IV] 4740.17 FI 3p3 4S* 3p3 2D* 4 4 4861.08 0.000 l.OOe -h 02 H 4 4861.33 H4 2p+ 2P* 4d-b 2D 8 32 4905.89 -0.011 4.048e - 02 ? 4905.89 * * 4907.15 -0.011 3.350e - 02 O II 4906.83 * * 4921.39 -0.015 1.271e + 00 He I 4921.93 * * 4923.98 -0.016 6.851e - 02 O II 4924.53 * * 4930.73 -0.017 1.576e - 01 [O III] 4931.80 * * 4958.22 -0.024 5.068e 4- 02 [O III] 4958.91 * * 5006.14 -0.036 1.508e + 03 [0 III] 5006.84 FI 2p2 3? 2p2 ID 5 5 5189.59 -0.081 1.724e - 01 [Ar III] 5191.82 * * 5197.61 -0.083 5.843e - 01 [NI] 5199.84 * * 5410.44 -0.134 1.086e -f 00 He II 5411.52 4.7 4f-b 2F* 7g+ 2 0 32 98 5516.71 -0.154 4.466e — 01 [Cl HI] 5517.66 FI 2p3 4S* 2p3 2D* 4 6 5536.64 -0.158 8.572e - 01 [Cl III] 5537.60 FI 2p3 4S* 2p3 2D* 4 4 5576.38 -0.164 1.467e - 01 [OI] 5577.34 * * 5677.93 -0.181 1.300e - 01 N II 5676.02 V3 3s 3P* 3p 3D 1 3 5753.56 -0.195 2.327e + 00 [N II] 5754.60 F3 2p2 ID 2p2 IS 5 1 5874.60 -0.215 1.726e -b 01 He I 5875.66 V ll 2p 3P* 3d 3D 9 15 5930.99 -0.224 7.772e - 02 N II 5931.78 V3 3p 3P 3d 3D* 3 5 6101.07 -0.251 2.648e - 01 [K IV] 6101.83 * * 6299.58 -0.282 8.923e + 00 [O I] 6300.34 FI 2p4 3P 2p4 ID 5 5 6311.34 -0.283 2.59le "b 00 [S HI] 6312.10 * * 6346.41 -0.289 1.019e - 01 Si II 6347.10 * * 6363.09 -0.291 3.020e -b 00 [OI] 6363.78 FI 2p4 3P 2p4 ID 3 5 6370.69 -0.292 1.869e - 01 Si II 6371.38 * * 6434.24 -0.302 3.032e - 01 [Ar V] 6435.10 * * 6461.32 -0.306 9.096e - 02 C II 6461.95 * * 6501.68 -0.311 5.482e - 02 ? 6501.68 * * 6526.50 -0.315 1.660e - 01 He II 6527.11 * * 6547.75 -0.318 2.680e -b 01 [N II] 6548.10 FI 2p2 3P 2p2 ID 3 5 6562.42 -0.320 2.752e + 02 H 3 6562.77 H3 2p+ 2P* 3d-b 2D 8 18 6583.15 -0.323 8.224e + 01 [N II] 6583.50 FI 2p2 3P 2p2 ID 5 5 6677.51 -0.336 4.600e -b 00 He I 6678.16 V46 2p IP* 3d ID 3 5 6715.79 -0.342 4.952e -b 00 [S II] 6716.44 F2 2p3 4S* 2p3 2D* 4 6 6730.17 -0.344 9.293e -b 00 [S II] 6730.82 F2 2p3 4S* 2p3 2D* 4 4 6794.17 -0.352 5.875e - 02 [K IV] 6795.00 * * 7064.20 -0.387 8.123e-b 00 He I 7065.25 VIO 2p 3P* 3s 3S 9 3 7134.74 -0.396 2.143e-b01 [Ar HI] 7135.80 FI 3p4 3P 3p4 ID 5 5 A.I. Galactic Planetary Nebulae 249

Aobs (•'^) f(A) 7(A) Ion Ao Mult Lower term Upper term gl g2 7169.45 -0.401 1.668e - 01 [Ar IV] 7170.62 ** 7176.32 -0.401 1.147e - 01 He II 7177.50 ** 7235.22 -0.408 3.162e - 01 C II 7236.42 * * 7261.56 -0.468 1.462e - 01 [Ar IV] 7762.76 * * 7280.14 -0.414 7.849e - 01 He I 7281.35 V45 2p IP* 3s IS 3 1 7318.17 -0.418 6.731e-h00 [O II] 7318.92 F2 2p3 2D* 2p3 2P* 6 2 7328.92 -0.420 5.997e -1- 00 [O II] 7329.67 F2 2p3 2D* 2p3 2P* 4 2 A.2. Magellanic Cloud Planetary Nebulae 250

A.2 Magellanic Cloud Planetary Nebulae c(H/3)(gal) = 0.03, c(H/3)(smc) =: 0.02

Table A.14: SMC N87 linelist.

'^obs (Â) f(A) /(A) Ion Ao Mult Lower term Upper term gl g2

3721.72 0.401 1.616e + 00 H 14 3721.94 * * 3725.79 0.400 4.795e + 00 [0 II] 3726.03 * * 3728.58 0.398 2.737e + 00 [O II] 3728.82 * * 3734.06 0.396 1.202e + 00 H 13 3734.37 ** 3749.98 0.389 1.798e + 00 H 12 3750.15 H12 2p+ 2P* 12d+ 2D 8 * 3770.46 0.379 1.933e + 00 H 11 3770.63 H ll 2p+ 2P* l ld + 2D 8 * 3797.73 0.367 3.007e + 00 H 10 3797.90 HIO 2p+ 2P* 10d+ 2D 8 * 3835.22 0.351 4.769e + 00 H 9 3835.39 H9 2p+ 2P* 9d+ 2D 8 * 3868.47 0.336 2.107e + 01 [Ne III] 3868.75 FI 2p4 3P 2p4 ID 5 5 3888.77 0.328 1.227e + 01 He I 3888.65 V2 2s 3S 3p 3P* 3 9 3967.17 0.296 7.883e + 00 [Ne III] 3967.46 FI 2p4 3P 2p4 ID 3 5 3969.78 0.294 1.168e + 01 H 7 3970.07 * * 4025.99 0.272 2.266e + 00 He I 4026.21 V18 2p 3P* 5d 3D 9 15 4070.30 0.255 3.615e - 0 1 [S II] 4068.60 FI 2p3 4S* 2p3 2P* 4 4 4071.46 0.255 1.713e - 01 O II 4069.62 VIO 3p 4D* 3d 4F 2 4 0.255 * O II 4069.89 VIO 3p 4D* 3d 4F 4 6 4073.86 0.254 6.482e - 02 O II 4072.16 VIO 3p 4D* 3d 4F 6 8 0.253 * O II 4075.86 VIO 3p 4D* 3d 4F 8 10 4078.06 0.253 1.630e - 01 [S II] 4076.35 FI 2p3 4S* 2p3 2P* 2 4 4080.55 0.252 2.691e - 02 O II 4078.84 VIO 3p 4D* 3d 4F 4 4 4086.82 0.250 3.617e - 02 O II 4083.90 V48b 3d 4F 4f G4* 6 8 4088.03 0.249 1.273e - 02 O II 4085.11 VIO 3p 4D* 3d 4F 6 6 4090.07 0.248 2.866e - 02 O II 4087.15 V48c 3d 4F 4f G3* 4 6 0.248 * O II 4089.29 V48a 3d 4F 4f G5* 10 12 4103.46 0.243 2.656e + 01 H 6 4101.74 H6 2p+ 2P* 6d+ 2D 8 72 0.236 * O II 4119.22 V20 3p 4P* 3d 4D 6 8 0.236 * O II 4120.28 V20 3p 4P* 3d 4D 6 6 0.236 * 0 II 4120.54 V20 3p 4P* 3d 4D 6 4 4122.46 0.236 3.014e - 01 He I 4120.84 V16 2p 3P* 5s 3S 9 3 0.236 * O II 4121.46 V19 3p 4P* 3d 4P 2 2 4145.37 0.227 3.569e - 01 He I 4143.76 V53 2p IP* 6d ID 3 5 4188.46 0.211 1.081e - 01 C HI 4186.90 V18 4f IF* 5g IG 7 9 4268.66 0.183 6.722e - 01 C II 4267.15 V6 3d 2D 4f 2F* 10 14 A.2,. Magellanic Cloud Planetary Nebulae 251

-^obs (Â) f(A) /(A) Ion Ao Mult Lower term Upper term gl g2 0.174 * O II 4291.25 V55 3d 4P 4f G3* 6 8 4294.11 0.174 4.653e - 02 O II 4292.21 V78c 3d 2F 4f F2* 6 6 4327.58 0.162 6.279e - 02 O II 4325.76 V2 3s 4P 3p 4P* 2 2 4342.02 0.158 4.641e-h01 H 5 4340.47 H5 2p+ 2P* 5d-t- 2D 8 50 4351.07 0.154 5.001e - 02 O II 4349.43 V2 3s 4P 3p 4P* 6 6 4359.53 0.152 3.899e - 02 O II 4357.25 V63a * * 4364.79 0.150 6.782e -H 00 [O III] 4363.21 F2 2p2 ID 2p2 IS 5 1 4389.55 0.141 6.238e - 01 He I 4387.93 V51 2p IP* 5d ID 3 5 4439.03 0.125 1.248e - 01 He I 4437.55 V50 2p IP* 5s IS 3 1 4458.31 0.119 2.516e - 02 Ne II 4457.24 V61d 3d 2D 4f 2]2i* 4 4 4473.10 0.114 5.336e 00 He I 4471.50 V14 2p 3P* 4d 3D 9 15 4641.10 0.063 8.323e - 02 O II 4638.86 VI 3s 4P 3p 4D* 2 4 4643.26 0.062 1.349e - 02 N III 4640.64 V2 3p 2P* 3d 2D 4 6 4644.06 0.062 8.319e - 02 O II 4641.81 VI 3s 4P 3p 4D* 4 6 4649.59 0.060 7.817e - 02 C HI 4647.42 VI 3s 3S 3p 3P* 3 5 4651.30 0.060 7.814e - 02 O II 4649.13 VI 3s 4P 3p 4D* 6 8 4652.42 0.060 4.688e - 02 C III 4650.25 VI 3s 3S 3p 3P* 3 3 4653.01 0.059 7.811e - 02 O II 4650.84 VI 3s 4P 3p 4D* 2 2 4653.64 0.059 1.562e - 02 C III 4651.47 VI 3s 3S 3p 3P* 3 1 4675.16 0.057 8.253e - 03 C IV 4658.10 F3 3d6 5D 3d6 3F2 9 9 4677.67 0.056 2.428e - 02 O II 4661.63 VI 3s 4P 3p 4D* 4 4 4677.66 0.052 2.481e - 02 O II 4676.24 VI 3s 4P 3p 4D* 6 6 4713.41 0.042 2.722e - 01 [Ar IV] 4711.37 FI 3p3 4S* 3p3 2D* 4 6 4715.22 0.041 9.326e - 01 He I 4713.17 v l2 2P 3p* 4S 3s 9 3 4742.22 0.034 4.059e - 01 [Ar IV] 4740.17 FI 3p3 4S* 3p3 2D* 4 4 4755.96 0.029 2.835e - 02 [Fe HI] 4755.96 * * 4758.33 0.028 1.370e - 02 7 4758.33 ** 4802.77 0.017 2.837e - 02 ? 4797.67 ** 4804.79 0.016 3.509e - 02 N II 4802.29 V20 3p 3D 3d 3D* 7 7 0.000 * H 4 4861.33 H4 2p4- 2P* 4d+ 2D 8 32 4923.89 -0.016 1.289e-kOO He I 4921.93 V48 2p IP* 4d ID 3 5 4926.49 -0.017 6.483e - 02 O II 4924.53 V28 3p 4S* 3d 4P 4 6 4933.09 -0.019 5.794e - 02 [O III] 4931.80 FI 2p2 3P 2p2 ID 1 5 4960.99 -0.026 1.905e -f 02 [O HI] 4958.91 FI 2p2 3P 2p2 ID 3 5 6557.81 -0.338 2.031e-h 00 [N II] 6548.10 FI 2p2 3P 2p2 ID 3 5 6571.25 -0.340 2.850e -k 02 H 3 6562.77 H3 2p+ 2P* 3d-k 2D 8 18 6590.48 -0.343 5.498e -f 00 [N II] 6583.50 FI 2p2 3P 2p2 ID 5 5 7068.72 -0.357 4.142e-t-00 He I 6678.16 V46 2p IP* 3d ID 3 5 A.2. Magellanic Cloud Planetary Nebulae 252

-^obs(^) f(A) /(A) Ion Ao Mult Lower term Upper term gl g2 6717.16 -0.362 4.437e - 01 [S II] 6716.44 F2 2p3 4S* 2p32D* 4 6 6731.54 -0.364 7.188e - 01 [S II] 6730.82 F2 2p3 4S* 2p32D* 4 4 7068.72 -0.409 1.043e + 01 He I 7065.25 VIO 2p 3P* 3s 3S 9 3 7141.32 -0.418 3.269e + 00 [Ar III] 7135.80 FI 3p4 3P 3p4ID 5 5 7291.42 -0.436 1.186e + 00 He I 7281.35 V45 2p IP* 3s IS 3 1 7330.40 -0.440 2.386e + 00 [0 II] 7318.92 F2 2p3 2D* 2p3 2P* 6 2 7341.16 -0.442 1.920e + 00 [O II] 7329.67 F2 2p3 2D* 2p32P* 4 2

Table A.15: LMC N66 linelist.

-^o b s ( ^ ) FWHM F(A) Ion Ao Mult Lower term Upper term gl g2

3727.04 0.75 1.130e-h01 [O II] 3726.03 * * 3729.83 2.59 1.941e-h01 [O II] 3728.82 * * 3870.43 2.09 1.090e-t-02 [Ne III] 3868.75 FI 2p4 3P 2p4 ID 5 5 3890.74 2.79 1.957e-t-01 He I 3888.65 V2 2s 38 3p 3P* 3 9 3969.18 2.39 4.589e-j-01 [Ne HI] 3967.46 FI 2p4 3P 2p4 ID 3 5 3971.79 2.06 1.305e4-01 H 7 3970.07 * * 4105.68 2.60 3.141e+01 H 6 4101.74 H6 2p+ 2P* 6d-b 2D 8 72 4231.02 2.94 4.231e-|-00 [FeV] 4227.20 F2 ** 4344.38 2.83 4.960e-t-01 H 5 4340.47 H5 2p+ 2P* 5d4- 2D 8 50 4367.17 2.84 3.347e-h01 [O III] 4363.21 F2 2p2 ID 2p2 IS 5 1 4475.38 2.50 3.719e-t-00 He I 4471.50 V14 2p 3P* 4d 3D 9 15 4545.70 3.22 2.152e-kOO He II 4541.59 4.9 4f+ 2F* 9g+ 2G 32 * 4559.22 2.00 2.947e-01 [Fe II] 4555.00 * * 4600.94 2.00 2.987e-01 O II 4596.18 V15 3s’ 2D 3p’2F* 4 6 4616.38 1.74 3.405e-01 ? 4616.38 * * 4618.55 1.74 1.542e-01 N II 4613.87 V5 * * 4637.28 1.79 1.975e-01 ? 4632.99 ** 4638.96 1.79 3.628e-01 N III 4634.14 V2 3p 2P* 3d 2D 2 4 4645.05 1.86 3.798e-01 N III 4640.64 V2 3p 2P* 3d 2D 4 6 4646.25 1.86 3.798e-01 0 II 4641.81 VI * * 4653.47 1.67 2.272e-01 O II 4649.13 VI 3s 4P 3p 4D* 6 8 4655.18 1.67 2.272e-01 O II 4650.84 VI * * 4662.13 1.65 4.314e-01 [Fe III] 4658.10 F3 3d6 5D 3d6 3F2 9 9 4665.04 1.65 1.699e-01 O II 4661.63 VI 3s 4P 3p 4D* 4 4 4669.65 2.55 5.246e-01 [Fe III] 4667.01 3F * * 4673.60 1.40 2.976e-01 O II 4669.27 V89b 3d 2D 4f D2* 4 6 4676.74 1.80 4.670e-01 0 II 4673.73 VI 3s 4P 3p 4D* 4 2 A.2. Magellanic Cloud Planetary Nebulae 253

-^obs(Â) FWHM F(A) Ion Ao Mult Lower term Upper term gl g2 4679.25 1.80 4.666e-01 O II 4676.24 VI 3s 4P 3p 4D* 6 6 4681.15 1.80 2.741e-01 N II 4678.14 * * 4690.02 2.85 6.730e+01 He II 4685.68 3.4 3d+ 2D 4f+ 2F* 18 32 4701.92 5.41 2.931e+00 WR? 4701.92 * * 4703.74 2.45 6.440e-01 [Fe VII] 4698.25 * * 4707.11 2^ 8 5.466e-01 [Fe HI] 4701.62 F3 3d6 5D 3d6 3F2 7 7 4715.37 2.46 1.175e+01 [Ar IV] 4711.37 FI 3p3 4S* 3p3 2D* 4 6 4717.18 2.46 4.180e+00 He I 4713.17 v l2 2P 3p* 4S 3s 9 3 4728.12 3.00 9.641e-01 [Ne IV] 4724.15 FI 2p3 2D* 2p3 2P* 4 4 4729.59 3.00 1.499e-f00 [Ne IV] 4725.62 FI 2p3 2D* 2p3 2P* 4 2 4737.79 2.00 3.231e-01 [Fe HI] 4733.93 3F * * 4744.47 2.87 1.285e+01 [Ar IV] 4740.17 FI 3p3 4S* 3p3 2D* 4 4 4759.58 1.49 1.335e-01 [Fe III] 4754.83 ** 4761.27 1.49 1.721e-01 ? 4756.52 * * 4769.48 1.49 2.179e-01 ? 4769.48 * * 4772.56 1.49 1.692e-01 [Fe HI] 4769.40 3F ** 4777.70 1.49 1.899e-01 Ne II 4772.93 ** 4782.65 1.49 2.079e-01 [Fe III] 4777.88 V20 ** 4800.75 2.50 3.736e-01 7 4797.67 * * 4865.50 3.19 l.OOOe+02 H4+HeII8 4861.33 H4 2p+ 2P* 4d+ 2D 8 32 4898.94 3.12 5.599e-01 ? * ** 4926.32 3.40 9.145e-01 He I 4921.93 V48 2p IP* 4d ID 3 5 4992.32 13.1 3.787e+02 [O HI] 4958.91 FI 2p2 3P 2p2 ID 3 5 5040.29 12.3 1.124e+03 [O HI] 5006.84 FI 2p2 3P 2p2 ID 5 5 5444.24 11.5 5.040e+00 He II 5411.52 4.7 4f+ 2F* 7g+ 2G 32 98 5786.96 17.7 1.979e+00 [N II] 5754.60 F3 2p2 ID 2p2 IS 5 1 5907.72 12.2 7.706e+00 He I 5875.66 V l l 2p 3P* 3d 3D 9 15 6335.11 24.0 7.234e+00 [O I] 6300.34 FI 2p4 3P 2p4 ID 5 5 6395.70 15.0 1.467e+00 [OI] 6363.78 ** 6560.37 4.11 2.630e+01 [N II] 6548.10 FI 2p2 3P 2p2 ID 3 5 6574.10 4.37 2.782e+02 H 3 6562.77 H3 2p+ 2P* 3d+ 2D 8 18 6593.34 4.01 7.511e+01 [N II] 6583.50 FI 2p2 3P 2p2 ID 5 5 6683.88 5.48 3.300e+00 Hel+WR 6678.16 V46 2p IP* 3d ID 3 5 6720.26 4.24 5.914e+00 [S II] 6716.44 F2 2p3 4S* 2p3 2D* 4 6 6734.37 3.79 7.533e+00 [S II] 6730.82 F2 2p3 4S* 2p3 2D* 4 4 7011.20 3.24 3.370e+00 [Ar V] 7005.74 * * 7072.05 3.99 3.828e+00 He I 7065.25 VIO 2p 3P* 3s 38 9 3 7144.66 3.98 1.556e+01 [Ar HI] 7135.80 FI 3p4 3P 3p4 ID 5 5 A.2. Magellanic Cloud Planetary Nebulae 254

■^obs (À) FWHM F^(A) Ion Ao Mult Lower term Upper term g l g2 7333.53 3.10 3.346e+00 [O II] 7318.92 F2 2p3 2D* 2p3 2P* 6 2

7344.02 3.74 3.659g+00 [O II] 7329.67 F2 2p3 2D* 2p3 2P* 4 2

c(H/3)(gal) = 0.110, c(H/3)(lmc) = 0.03

Table A .16: LMC N141 linelist.

-^obs f(A) /(A) Ion Aq Mult Lower term Upper term gl g2

3727.44 0.375 1.066e + 01 [O II] 3726.03 ** 3730.23 0.374 6.528e + 00 [O II] 3728.82 * * 3751.71 0.365 1.867e + 00 H 12 3750.15 H12 2p+ 2P* 12d+ 2D 8 * 3772.20 0.356 1.685e + 00 H 11 3770.63 Hll 2p+ 2P* lld+ 2D 8 * 3799.48 0.344 2.457e + 00 H 10 3797.90 HIO 2p+ 2P* 10d+ 2D 8 * 3836.99 0.329 5.324e + 00 H 9 3835.39 H9 2p+ 2P* 9d+ 2D 8 * 3870.33 0.315 4.279e + 01 [Ne III] 3868.75 FI 2p4 3P 2p4 ID 5 5 3890.64 0.306 1.108e+ 01 He I 3888.65 V2 2s 3S 3p 3P* 3 9 3969.08 0.276 1.533e + 01 [Ne HI] 3967.46 FI 2p4 3P 2p4 ID 3 5 3971.69 0.274 1.224e + 01 H 7 3970.07 * * 4027.91 0.253 2.856e + 00 He I 4026.21 V18 2p 3P* 5d 3D 9 15 4072.63 0.237 1.308e + 00 [S II] 4068.60 FI 2p3 4S* 2p3 2P* 4 4 4073.79 0.237 1.879e - 01 O II 4069.62 VIO 3p 4D* 3d 4F 2 4 4076.19 0.236 1.238e - 01 O II 4072.16 VIO 3p 4D* 3d 4F 6 8 4079.90 0.235 2.536e - 01 O II 4075.86 VIO 3p 4D* 3d 4F 8 10 4080.39 0.235 3.234e - 01 [S II] 4076.35 FI 2p3 4S* 2p3 2P* 2 4 4101.60 0.227 3.267e - 01 N HI 4097.33 VI 3s 2S 3p 2P* 2 4 4105.78 0.225 2.918e + 01 H 6 4101.74 H6 2p+ 2P* 6d+ 2D 8 72 4125.03 0.219 3.568e — 01 He I 4120.84 V16 2p 3P* 5s 3S 9 3 4147.69 0.211 3.973e - 01 He I 4143.76 V53 2p IP* 6d ID 3 5 4190.46 0.196 5.219e - 02 C HI 4186.90 V18 4f IF* 5g IG 7 9 4193.36 0.195 6.678e - 02 O II 4189.79 V36 3p’ 2F* 3d’ 2 0 8 10 4271.09 0.169 6.829e - 01 C II 4267.15 V6 3d 2D 4f 2F* 10 14 4344.45 0.145 5.354e + 01 H 5 4340.47 H5 2p+ 2P* 5d+ 2D 8 50 4367.25 0.138 1.023e + 01 [O HI] 4363.21 F2 2p2 ID 2p2 IS 5 1 4391.96 0.130 6.964e - 01 He I 4387.93 V51 2p IP* 5d ID 3 5 4442.17 0.115 1.016e -01 He I 4437.55 V50 2p IP* 5s IS 3 1 4475.61 0.105 5.771e + 00 He I 4471.50 V14 2p 3P* 4d 3D 9 15 4556.98 0.081 1.918e - 01 N II 4552.53 * * 4566.39 0.078 3.140e - 01 Mg I] 4562.60 3s2 IS 3s3p 3P* 1 5 4574.90 0.076 1.221e + 00 Mg I] 4571.10 3s2 IS 3s3p 3P* 1 3 A.2. Magellanic Cloud Planetary Nebulae 255

-^obs (Â) f(A) 7(A) Ion Ao Mult Lower term Upper term gl g2 4594.73 0.070 1.749e - 01 O II 4590.97 V15 3s’ 2D 3p’ 2F* 6 8 4638.01 0.059 9.424e - 02 N III 4634.14 V2 3p 2P* 3d 2D 2 4 4642.74 0.057 2.549e - 02 O II 4638.86 VI 3s 4P 3p 4D* 2 4 4644.52 0.057 1.877e - 01 N III 4640.64 V2 3p 2P* 3d 2D 4 6 4645.69 0.056 1.202e - 01 O II 4641.81 VI 3s 4P 3p 4D* 4 6 4651.05 0.055 7.688e - 02 C III 4647.42 VI 3s 3S 3p 3P* 3 5 4652.76 0.054 1.267e - 01 O II 4649.13 VI 3s 4P 3p 4D* 6 8 4653.88 0.054 6.938e - 02 C III 4650.25 VI 3s 3S 3p 3P* 3 3 4654.47 0.054 8.360e - 02 0 II 4650.84 VI 3s 4P 3p 4D* 2 2 4655.10 0.054 3.476e - 02 C III 4651.47 V I 3s 3S 3p 3P* 3 1 4662.12 0.052 1.041e - 01 C IV 4658.10 F3 3d6 5D 3d6 3F2 9 9 4665.65 0.051 5.177e - 02 O II 4661.63 VI 3s 4P 3p 4D* 4 4 4680.62 0.047 4.646e - 02 O II 4676.24 VI 3s 4P 3p 4D* 6 6 4689.66 0.045 3.130e - 01 He II 4685.68 3.4 3d+ 2D 4f-b 2F* 18 32 4705.54 0.040 5.787e - 02 [Fe III] 4701.62 F3 3d6 5D 3d6 3F2 7 7 4715.34 0.038 8.511e - 01 [Ar IV] 4711.37 FI 3p3 4S* 3p3 2D* 4 6 4717.14 0.037 8.814e - 01 He I 4713.17 vl2 2P 3p* 4S 3s 9 3 4744.05 0.030 1.222e4-00 [Ar IV] 4740.17 FI 3p3 4S* 3p3 2D* 4 4 4818.38 0.011 3.189e - 02 [Fe II] 4814.55 * * 4864.89 0.000 1.000e-k02 H 4 4861.33 H4 2p-b 2P* 4d+ 2D 8 32 4884.98 -0.005 7.757e - 02 [Fe HI] 4881.11 ** 4894.31 -0.008 3.700e - 02 O II 4890.86 * * 4925.79 -0.015 1.407e -b 00 He I 4921.93 V48 2p IP* 4d ID 3 5 4935.06 -0.018 1.037e - 01 [O HI] 4931.80 FI 2p2 3P 2p2 ID 1 5 4963.05 -0.024 3.154e-b02 [O HI] 4958.91 FI 2p2 3P 2p2 ID 3 5 6560.35 -0.307 5.918e-f 00 [N II] 6548.10 FI 2p2 3P 2p2 ID 3 5 6574.07 -0.309 2.850e + 02 H 3 6562.77 H3 2p-b 2P* 3d-b 2D 8 18

6593.29 -0.312 1.756e -b 01 [N II] 6583.50 FI 2p2 3P 2p2 ID 5 5 6683.36 -0.325 4.704e -b 00 He I 6678.16 V46 2p IP* 3d ID 3 5 6720.29 -0.330 1.474e -b 00 [S II] 6716.44 F2 2p3 4S* 2p3 2D* 4 6

6734.37 -0.333 2.718e-b 00 [S II] 6 7 3 0 .8 2 F2 2p3 4S* 2p3 2D* 4 4 7071.99 -0.377 1.052e -b 01 He I 7065.25 VIO 2p 3P* 3s 3S 9 3 7144.58 -0.386 1.115e-b01 [Ar HI] 7135.80 FI 3p4 3P 3p4 ID 5 5 7294.46 -0.404 1.0386 -b 00 He I 7281.35 V45 2p IP* 3s IS 3 1 7333.61 -0.409 4.3596 -b 00 [0 II] 7318.92 F2 2p3 2D* 2p3 2P* 6 2 7344.37 -0.410 3.6286 -b 00 [O II] 7329.67 F2 2p3 2D* 2p3 2P* 4 2 7752.96 -0.456 2.4296 -b 00 C IV? 7736.00 0 0 A.s. Galactic H II Regions 256

A.3 Galactic H ii Regions c(H/?)(Bal) = 1.84

T able A .17: M 17 lin elist.

-^obs ( Â ) f(A) /(A) Ion Ao Mult Lower term Upper term gl g2

* 0.257 * H 14 3721.94 H14 2p+ 2P* 14d+ 2D 8 * 0.257 * [S III] 3721.63 F2 3p2 3P 3p2 IS 3 1 3726.99 0.257 8.7297e + 01 [O II] 3726.03 FI 2p3 4S* 2p3 2D* 4 4 0.256 * [O II] 3728.82 FI 2p3 4S* 2p3 2D* 4 6 0.255 * H 13 3734.37 H13 2p+ 2P* 13d+ 2D 8 * 3749.86 0.253 3.324e + 00 H 12 3750.15 H12 2p+ 2P* 12d+ 2D 8 * 3770.34 0.249 3.786e + 00 H 11 3770.63 H ll 2p+ 2P* lld+ 2D 8 * 3797.30 0.244 4.536e + 00 H 10 3797.90 HIO 2p+ 2P* 10d+ 2D 8 * 3819.01 0.240 1.144e + 00 He I 3819.62 ** 3834.78 0.237 6.591e + 00 H 9 3835.39 H9 2p+ 2P* 9d+ 2D 8 * 3868.84 0.230 1.836e + 01 [Ne III] 3868.75 FI 2p4 3P 2p4 ID 5 5 3889.14 0.226 1.628e + 01 He I 3888.65 V2 2s 3S 3p 3P* 3 9 3967.56 0.211 2.013e + 01 [Ne HI] 3967.46 FI 2p4 3P 2p4 ID 3 5 4009.31 0.202 2.979e - 01 He I 4009.26 V55 2p IP* 7d ID 3 5 4026.29 0.198 2.254e + 00 He I 4026.21 V18 2p 3P* 5d 3D 9 15 4069.00 0.189 5.795e - 01 [S II] 4068.60 FI 2p3 4S* 2p3 2P* 4 4 0.189 * O II 4069.62 VIO 3p 4D* 3d 4F 2 4 0.189 * O II 4069.89 VIO 3p 4D* 3d 4F 4 6 0.187 * O II 4075.86 VIO 3p 4D* 3d 4F 8 10 4076.23 0.187 1.562e - 01 [S II] 4076.35 FI 2p3 4S* 2p3 2P* 2 4 0.185 * O II 4085.11 VIO 3p 4D* 3d 4F 6 6 0.185 * O II 4087.15 V48c 3d 4F 4f G3* 4 6 0.184 * O II 4089.29 V48a 3d 4F 4f G5* 10 12 0.183 * N III 4097.33 V I 3s 2S 3p 2P* 2 4 4101.76 0.182 2.570e + 01 H 6 4101.74 H6 2p+ 2P* 6d-|- 2D 8 72 4119.71 0.178 1.156e - 01 O II 4119.22 V20 3p 4P* 3d 4D 6 8 4121.33 0.177 1.911e - 01 He I 4120.84 V16 2p 3P* 5s 3S 9 3 4143.87 0.172 2.806e - 01 He I 4143.76 * * 4267.19 0.144 4.817e - 01 C II 4267.15 V6 3d 2D 4f 2F* 10 14 4317.27 0.132 1.115e - 01 O II 4317.14 V2 3s 4P 3p 4P* 2 4 4319.48 0.132 3.304e - 02 O II 4319.63 V2 3s 4P 3p 4P* 4 6 0.130 * O II 4325.76 V2 3s 4P 3p 4P* 2 2 4340.43 0.127 4.779e + 01 H 5 4340.47 H5 2p+ 2P* 5dT 2D 8 50 A.3. Galactic H 11 Regions 257

-^obs (Â) f(A) /(A) Ion Ao Mult Lower term Upper term gl g2 0.125 6.700e - 02 O II 4345.56 V2 3s 4P 3p 4P* 4 2 4349.01 0.125 1.070e - 01 0 II 4349.43 V2 3s 4P 3p 4P* 6 6 4363.22 0.121 1.023e -f 00 [O III] 4363.21 F2 2p2 ID 2p2 IS 5 1 4388.02 0.115 6.633e - 01 He I 4387.93 V51 2p IP* 5d ID 3 5 4437.83 0.103 8.337e - 02 He I 4437.55 V50 2p IP* 5s IS 3 1 4471.49 0.095 4.918e + 00 He I 4471.50 V14 2p 3P* 4d 3D 9 15 4607.46 0.062 5.954e - 02 N II 4607.16 ** 4630.76 0.057 6.647e - 02 N II 4630.54 ** 4634.26 0.056 4.344e - 02 N III 4634.14 V2 3p 2P* 3d 2D 2 4 4638.98 0.055 1.056e - 01 O II 4638.86 VI 3s 4P 3p 4D* 2 4 4640.76 0.054 3.486e - 02 N III 4640.64 V2 3p 2P* 3d 2D 4 6 4641.93 0.054 1.611e - 01 O II 4641.44 * * 4649.38 0.052 1.375e - 01 O II 4649.13 VI 3s 4P 3p 4D* 6 8 4651.09 0.052 8.713e - 02 O II 4650.84 VI 3s 4P 3p 4D* 2 2 4658.49 0.050 2.791e - 01 [Fe III] 4658.10 * * 4662.02 0.049 1.043e - 01 O II 4661.63 VI 3s 4P 3p 4D* 4 4 4676.91 0.046 6.298e - 02 O II 4676.24 V I 3s 4P 3p 4D* 6 6 4711.60 0.037 8.205e - 02 [Ar IV] 4711.37 FI 3p3 4S* 3p3 2D* 4 6 4713.40 0.036 5.267e - 01 He I 4713.17 V12 2P 3p* 4S 3s 9 3 4740.37 0.030 7.040e - 02 [Ar IV] 4740.17 FI 3p3 4S* 3p3 2D* 4 4 4861.55 0.000 l.OOOe-h 02 H 4 4861.33 H4 2p+ 2P* 4d-k 2D 8 32 4881.20 -0.005 1.025e - 01 [Fe HI] 4881.11 ** 4906.87 -0.011 5.174e - 02 O II 4906.83 ** 4922.27 -0.015 1.340e -f 00 He I 4921.93 ** 4931.47 -0.017 4.180e - 02 [O III] 4931.80 * * 4959.13 -0.024 1.179e + 02 [O HI] 4958.91 * * 5006.43 -0.036 3.550e -t- 02 [0 III] 5006.84 FI 2p2 3P 2p2 ID 5 5 * 5198.36 -0.083 2.846e - 01 [N I] 5199.84 * 5517.48 -0.154 4.546e - 01 [Cl III] 5517.66 FI 2p3 4S* 2p3 2D* 4 6 5537.71 -0.158 3.976e - 01 [01 HI] 5537.60 FI 2p3 4S* 2p3 2D* 4 4 5754.74 -0.195 3.245e - 01 [N II] 5754.60 F3 2p2 ID 2p2 IS 5 1 5875.44 -0.215 1.456e-H 01 He I 5875.66 V l l 2p 3P* 3d 3D 9 15 6312.09 -0.283 1.436e 00 [S HI] 6312.10 F3 2p2 ID 2p2 IS 5 1 6548.71 -0.318 9.156e + 00 [N II] 6548.10 FI 2p2 3P 2p2 ID 3 5 6563.11 -0.320 3.012e-F02 H 3 6562.77 H3 2p+ 2P* 3d 4- 2D 8 18 6583.91 -0.323 2.745e -f 01 [N II] 6583.50 FI 2p2 3P 2p2 ID 5 5 6678.59 -0.336 3.896e -f 00 He I 6678.16 V46 2p IP* 3d ID 3 5 6717.02 -0.342 4.049e -I- 00 [S II] 6716.44 F2 2p3 4S* 2p3 2D* 4 6 A.3. Galactic H ll Regions 258

-^obs (^ ) f(A) /(A) Ion Ao Mult Lower term Upper term gl g2 6731.39 -0.344 4.119e + 00 [S II] 6730.82 F2 2p3 48* 2p3 2D* 4 4 7065.48 -0.387 3.510e -f 00 He I 7065.25 VIO 2p 3P* 3s 38 9 3 7135.90 -0.396 1.071e4-01 [Ar III] 7135.80 FI 3p4 3P 3p4 ID 5 5 7280.02 -0.414 9.505e - 01 He I 7281.35 V45 2p IP* 3s 18 3 1 7319.06 -0.418 1.123e-f-00 [O II] 7318.92 F2 2p3 2D* 2p3 2P* 6 2 7329.62 -0.420 8.729e - 01 [O II] 7329.67 F2 2p3 2D* 2p3 2P* 4 2

c(H/3)(Bal) - 1.25

Table A. 18: NGC 3576 AAT linelist.

-^obs (^ ) f(A) /(A) Ion Aq Mult Lower term Upper term gl g2

3531.25 0.302 9.701e - 01 He I 3530.50 V36 * * 3553.94 0.295 3.717e - 01 He I 3554.42 V34 * * 3566.73 0.291 L003e + 00 7 ** * 3613.22 0.277 4.920e - 01 He I 3613.64 V6 ** 3634.00 0.272 5.131e - 01 He I 3634.25 V28 2p 3P* 8d 3D 9 15 3671.17 0.266 5.132e - 01 H 24 3671.48 H24 2p+ 2P* 24d4- 2D 8 * 3673.39 0.265 4.908e - 01 H 23 3673.74 H23 2p4- 2P* 23d4- 2D 8 * 3676.01 0.265 6.119e - 01 H 22 3676.36 H22 2p4- 2P* 22d4- 2D 8 * 3679.01 0.265 6.266e - 01 H 21 3679.36 H21 2p4- 2P* 21d+ 2D 8 * 3682.46 0.264 6.920e - 01 H 20 3682.81 H20 2p+ 2P* 20d4- 2D 8 * 3686.48 0.263 8.046e - 01 H 19 3686.83 H19 2p-^ 2P* 19d4- 2D 8 * 3691.21 0.262 9.599e - 01 H 18 3691.56 H18 2p+ 2P* 18d4- 2D 8 * 3696.80 0.262 1.062e 4- 00 H 17 3697.15 H17 2p4- 2P* 17d4- 2D 8 * 3703.50 0.260 L314e + 00 H 16 3703.86 H16 2p+ 2P* 16d4- 2D 8 * 3704.66 0.260 6.495e - 01 He I 3705.02 V25 * * 3711.61 0.259 L588e + 00 H 15 3711.97 H15 2p+ 2P* 15d4- 2D 8 * 0.257 * [S HI] 3721.63 F2 3p2 3P 3p2 18 3 1 3721.49 0.257 2.798e -h 00 H 14 3721.94 H14 2p4- 2P* 14d4- 2D 8 * 3725.66 0.257 7.512e4-01 [O II] 3726.03 FI 2p3 48* 2p3 2D* 4 4 3728.45 0.256 5.498e4-01 [O II] 3728.82 FI 2p3 48* 2p3 2D* 4 6 3733.99 0.255 2.235e + 00 H 13 3734.37 H13 2p4- 2P* 13d4- 2D 8 * 3749.77 0.253 2.908e -f 00 H 12 3750.15 H12 2p4- 2P* 12d4- 2D 8 * 3770.25 0.249 3.765e -|- 00 H 11 3770.63 Hll 2p4- 2P* lld 4 - 2D 8 * 3797.52 0.244 5.167e + 00 H 10 3797.90 HIO 2p+ 2P* 10d+ 2D 8 * 3819.21 0.240 L189e + 00 He I 3819.62 V22 2p 3P* 6d 3D 9 15 3835.00 0.237 7.549e 4- 00 H 9 3835.39 H9 2p4- 2P* 9d+ 2D 8 * 3855.69 0.233 2.441e - 01 Si II 3856.02 VI * * A.3. Galactic H II Regions 259

-^obs (Â) f(A) /(A) Ion A q Mult Lower term Upper term gl g2 3862.27 0.232 2.839e - 01 Si II 3862.60 VI * * 3868.34 0.230 2.069e + 01 [Ne III] 3868.75 FI 2p4 3P 2p4 ID 5 5 3871.33 0.230 1.766e - 01 He I 3871.82 V60 * * 3888.64 0.226 1.648e + 01 He I 3888.65 V2 2s 38 3p 3P* 3 9 3912.83 0.222 7.594e - 02 ? 3912.83 * * 3914.02 0.221 6.529e - 02 ? 3914.02 ** 3918.66 0.220 1.034e - 01 C II 3918.98 V4 * * 3920.37 0.220 9.244e - 02 C II 3920.69 V4 * * 3926.03 0.219 1.527e - 01 He I 3926.54 V58 ** 3964.36 0.211 8.977e - 01 He I 3964.73 V5 ** 3967.04 0.211 5.957e + 00 [Ne HI] 3967.46 FI 2p4 3P 2p4 ID 3 5 3969.70 0.210 1.492e + 01 H 7 3970.07 * * 4008.84 0.202 2.314e - 01 He I 4009.26 V55 2p IP* 7d ID 3 5 4025.83 0.198 2.007e + 00 He I 4026.21 V18 2p 3P* 5d 3D 9 15 4037.19 0.196 2.224e - 02 ? 4037.19 * * 4042.91 0.195 4.546e - 02 ? 4042.91 * * 4059.17 0.191 2.417e - 02 ? 4059.17 * * 4060.28 0.191 1.611e - 02 [F IV] 4060.23 2s2 2F 2s2 2 [4] 8 10 4065.35 0.190 4.148e - 02 7 4065.35 * * 4068.21 0.189 1.283e + 00 [S II] 4068.60 FI 2p3 48* 2p3 2P* 4 4 4069.37 0.189 1.409e - 01 O II 4069.62 VIO 3p 4D* 3d 4F 2 4 0.189 * O II 4069.89 VIO 3p 4D* 3d 4F 4 6 4071.77 0.188 1.207e - 01 O II 4072.16 VIO 3p 4D* 3d 4F 6 8 4075.47 0.187 l.OlOe - 01 O II 4075.86 VIO 3p 4D* 3d 4F 8 10 4075.96 0.187 4.401e - 01 [S II] 4076.35 FI 2p3 48* 2p3 2P* 2 4 4080.81 0.186 3.387e - 02 ? 4080.81 * * 4084.90 0.185 1.700e - 02 O II 4085.11 VIO 3p 4D* 3d 4F 6 6 4086.93 0.185 4.147e - 02 O II 4087.15 V48c 3d 4F 41 G3* 4 6

4089.07 0.184 3.201e - 02 O II 4089.29 V48a 3d 4F 41 G5* 10 12 4096.94 0.183 6.5096 — 02 N HI 4097.33 VI 3s 28 3p 2P* 2 4 4101.33 0.182 2.564e + 01 H 6 4101.74 H6 2p+ 2P* 6d + 2D 8 72 4118.83 0.178 2.719e - 02 O II 4119.22 V20 * * 0.178 * O II 4120.28 V20 3p 4P* 3d 4D 6 6 0.177 * O II 4120.54 V20 3p 4P* 3d 4D 6 4 4120.45 0.177 2.138e - 01 He I 4120.84 V16 2p 3P* 5s 38 9 3 4132.64 0.175 5.233e - 02 O II 4132.80 V19 * * 4143.33 0.172 3.073e - 01 He I 4143.76 V53 2p IP* 6d ID 3 5 4152.88 0.170 3.612e - 02 0 II 4153.30 V19 3p 4P* 3d 4P 4 6 A.3. Galactic H II Regions 260

-^obs (-^) f(A) /(A) Ion Ao Mult Lower term Upper term gl g2 4156.11 0.169 2.086e - 02 O II 4156.53 V19 3p 4P* 3d 4P 6 4 4168.62 0.167 6.589e - 02 He I 4168.97 V52 ** 0.167 * O II 4169.22 V19 3p 4P* 3d 4P 6 6 4170.28 0.166 3.168e - 02 ? 4170.28 ** 4185.44 0.163 3.304e - 02 O II 4185.45 V36 3p' 2F* 3d’ 2G 6 8

0.163 * O II 4185.45 ** 4189.78 0.162 3.148e - 02 O II 4189.79 V36 3p’ 2F* 3d’ 2G 8 10 4253.66 0.147 4.209e - 02 O II 4254.00 V lO l 3d’ 2G 4 f’ H5* 18 22 4266.73 0.144 3.248e - 01 C II 4267.15 V6 3d 2D 4f 2F* 10 14 4276.78 0.142 1.339e - 01 O II 4275.55 V67a 3d 4D 4f F4* 8 10 4287.00 0.139 1.115e - 0 1 [Fe II] 4287.40 7F ** 0.139 * O II 4288.82 V53c 3d 4P 4f Dl* 2 4 4340.04 0.127 4.764e -f 01 H 5 4340.47 H5 2p+ 2P* 5d+ 2D 8 50 4345.30 0.125 1.083e - 01 O II 4345.56 V2 3s 4P 3p 4P* 4 2 4349.17 0.125 7.772e - 02 O II 4349.43 V2 3s 4P 3p 4P* 6 6 4357.14 0.123 7.565e - 02 O II 4357.25 V63a * * 4358.98 0.122 1.188e - 01 ? 4358.98 * * 4362.80 0.121 1.513e + 00 [O III] 4363.21 F2 2p2 ID 2p2 IS 5 1 4368.06 0.120 1.068e - 01 C II? 4368.14 V45 ** 4387.49 0.115 5.500e - 01 He I 4387.93 V51 2p IP* 5d ID 3 5 4413.23 0.109 1.047e - 01 [Fe II] 4413.78 7F * * 4471.07 0.095 4.846e + 00 He I 4471.50 V14 2p 3P* 4d 3D 9 15 4603.76 0.063 4.014e - 02 ? 4603.76 * * 4606.98 0.062 5.750e - 02 N II 4607.16 V5 3s 3P* 3p 3P 1 3 4608.66 0.062 5.347e - 02 ? 4608.66 * * 0.062 * O II 4609.44 V92a 3d 2D 4f F4* 6 8 4610.02 0.062 5.403e - 02 O II 4610.20 V92c 3d 2D 4f F2* 4 6 4614.92 0.061 4.838e - 02 7 4614.92 ** 4620.88 0.059 4.554e - 02 N II 4621.39 V5 ** 4630.03 0.057 3.580e - 02 N II 4630.54 V5 3s 3P* 3p 3P 5 5 0.056 * N HI 4634.14 V2 3p 2P* 3d 2D 2 4 4638.46 0.055 6.924e - 02 O II 4638.86 VI * * 4640.24 0.054 2.570e - 02 N HI 4640.64 V2 3p 2P* 3d 2D 4 6 4641.42 0.054 1.081e - 01 O II 4641.81 VI 3s 4P 3p 4D* 4 6 4648.60 0.052 1.609e - 01 O II 4649.13 VI 3s 4P 3p 4D* 6 8 4650.31 0.052 1.527e - 01 O II 4650.84 VI 3s 4P 3p 4D* 2 2 0.050 [Fe HI] 4658.10 F3 3d6 5D 3d6 3F2 9 9 4660.02 0.049 3.407e - 02 ? 4660.02 ** A.3. Galactic H II Regions 261

-^obs (Â) f(A) 7(A) Ion Ao Mult Lower term Upper term gl g2 4661.23 0.049 8.619e - 02 O II 4661.63 VI 3s 4P 3p 4D* 4 4 4675.84 0.046 5.858c — 02 0 II 4676.24 V I 3s 4P 3p 4D* 6 6 4696.03 0.041 3.162e - 02 0 II 4696.35 V I 3s 4P 3p 4D* 6 4 4701.16 0.039 2.165c - 01 [Fe III] 4701.62 F3 3d6 5D 3d6 3F2 7 7 4705.72 0.038 4.219c - 02 O II 4705.35 V25 ** 4958.86 -0.024 1.227c+ 02 [O III] 4958.91 FI 2p2 3P 2p2 ID 3 5 5006.82 -0.036 3.713c+ 02 [O III] 5006.84 FI 2p2 3P 2p2 ID 5 5

5200.35 -0.083 6.403c - 01 [NI] 5199.84 ** 5270.91 -0.100 3.363c - 01 [Fe III] 5270.40 IF * * 5518.20 -0.154 4.707c - 01 [Cl III] 5517.66 FI 2p3 4S* 2p3 2D* 4 6 5538.14 -0.158 5.337c - 01 [Cl III] 5537.60 FI 2p3 4S* 2p3 2D* 4 4 5577.88 -0.164 2.577c + 00 [OI] 5577.34 F3 ** 5754.86 -0.195 5.929c - 01 [N II] 5754.60 F3 2p2 ID 2p2 IS 5 1 5875.76 -0.215 1.380c + 01 He I 5875.66 V ll 2p 3P* 3d 3D 9 15 5913.00 -0.221 3.427c - 02 He II 5913.00 ** 5931.52 -0.224 1.094c - 01 N II 5931.52 * * 5957.35 -0.228 1.467c - 01 Si II 5957.35 * * 5978.71 -0.231 1 .0 1 1 c - 01 Si II 5978.71 ** 6562.97 -0.320 2.8 8 1 c+ 02 H 3 6562.77 H3 2p+ 2P* 3d+ 2D 8 18 -0.323 * [N II] 6583.50 FI 2p2 3P 2p2 ID 5 5 6678.74 -0.336 3.681c + 00 He I 6678.16 V46 2p IP* 3d ID 3 5 6717.02 -0.342 6.947c + 00 [S II] 6716.44 F2 2p3 4S* 2p3 2D* 4 6 6731.40 -0.344 8.764c + 00 [S II] 6730.82 F2 2p3 4S* 2p3 2D* 4 4 7065.67 -0.387 5.455c + 00 He I 7065.25 VIO 2p 3P* 3s 3S 9 3 7136.12 -0.396 1 .515c+ 01 [Ar HI] 7135.80 FI 3p4 3P 3p4 ID 5 5 7319.06 -0.418 4 .1 4 7 c+ 00 [O II] 7318.92 F2 2p3 2D* 2p3 2P* 6 2 7329.62 -0.420 3.309c + 00 [O II] 7329.67 F2 2p3 2D* 2p3 2P* 4 2

c(H^)(Bal) = 1.25

Table A.19: NGC 3576 ESO linelist.

Aobs(^) f(A) /(A) Ion Ao Mult Lower term Upper term gl g2

4008.73 0.202 2.526c - 01 He I 4009.26 V55 2p IP* 7d ID 3 5 4025.77 0.198 1.891c+ 00 He I 4026.21 V18 2p 3P* 5d 3D 9 15 4068.12 0.189 1.540c + 00 [S II] 4068.60 FI 2p3 4S* 2p3 2P* 4 4 0.187 * O II 4075.86 VIO 3p 4D* 3d 4F 8 10 4075.73 0.187 5.658c - 01 [S II] 4076.35 FI 2p3 4S* 2p3 2P* 2 4 4101.16 0.182 2.532c+ 01 H 6 4101.74 H6 2p+ 2P* 6d+ 2D 8 72 0.178 * O II 4119.22 V20 3p 4P* 3d 4D 6 8 A.3. Galactic H ll Regions 262

^obs (Â) f(A) 7(A) Ion Ao Mult Lower term Upper term gl g2 4120.36 0.177 2.104e - 01 He I 4120.84 V16 2p 3P* 5s 3S 9 3 4143.17 0.172 1.715e - 01 He I 4143.76 ** 4266.58 0.144 3.122e - 01 C II 4267.15 V6 3d 2D 4f 2F* 10 14 0.142 * O II 4275.55 V67a 3d 4D 4f F4* 8 10 4276.29 0.142 7.218e - 02 O II 4275.99 V67b 3d 4D 4f F3* 4 6 4286.98 0.140 1.155e - 01 O II 4285.69 V78b 3d 2F 4f F3* 6 8 4316.94 0.132 1.807e - 01 O II 4317.14 V2 3s 4P 3p 4P* 2 4 4339.87 0.127 4.680e -1- 01 H 5 4340.47 H5 2p+ 2P* 5d-|- 2D 8 50 0.125 * O II 4345.56 V2 3s 4P 3p 4P* 4 2 4349.02 0.125 4.847e - 02 O II 4349.43 V2 3s 4P 3p 4P* 6 6 4362.64 0.121 1.221e-h00 [O HI] 4363.21 F2 2p2 ID 2p2 IS 5 1 4367.45 0.120 1.214e - 01 O II 4366.89 V2 3s 4P 3p 4P* 6 4 4387.39 0.115 4.441e - 01 He I 4387.93 V51 2p IP* 5d ID 3 5 4412.28 0.109 7.132e - 02 Ne II 4413.22 * * 4413.96 0.109 6.229e - 02 O II 4414.90 V5 3s 2P 3p 2D* 4 6 4416.03 0.109 7.943e - 02 O II 4416.97 V5 3s 2P 3p 2D* 2 4 4437.45 0.103 7.845e - 02 He I 4437.55 V50 2p IP* 5s IS 3 1 0.101 * O II 4448.19 V35 3p’ 2F* 3d' 2F 8 8 4449.91 0.101 3.336e - 02 ? 4449.91 * * 4451.61 0.100 5.400e - 02 O II 4452.37 ** 4470.93 0.095 4.271e -f 00 He I 4471.50 V14 2p 3P* 4d 3D 9 15 0.056 * N HI 4634.14 V2 3p 2P* 3d 2D 2 4 4638.79 0.055 5.086e - 02 O II 4638.49 ** 4640.57 0.054 4.221e - 02 N III 4640.64 V2 3p 2P* 3d 2D 4 6 4641.74 0.054 9.279e - 02 O II 4641.44 ** 4648.70 0.052 9.293e - 02 O II 4649.13 VI 3s 4P 3p 4D* 6 8 4650.41 0.052 5.530e - 02 O II 4650.84 VI 3s 4P 3p 4D* 2 2 4657.68 0.050 5.891e - 01 [Fe HI] 4658.10 ** 4661.21 0.049 9.155e - 02 O II 4661.63 VI 3s 4P 3p 4D* 4 4 4675.90 0.046 5.604e - 02 O II 4676.24 VI 3s 4P 3p 4D* 6 6 4701.14 0.039 1.793e - 01 [Fe III] 4701.62 * * 4710.98 0.037 8.705e - 02 [Ar IV] 4711.37 FI 3p3 4S* 3p3 2D* 4 6 4712.78 0.036 4.919e - 01 He I 4713.17 V12 2P 3p* 4S 3s 9 3 4733.55 0.031 4.522e - 02 [Fe HI] 4733.93 * * 4739.65 0.030 7.545e - 02 [Ar IV] 4740.17 F I 3p3 4S* 3p3 2D* 4 4 4754.43 0.026 1.094e - 01 [Fe HI] 4754.83 * * 4860.93 0.000 l.OOOe-f-02 H 4 4861.33 H4 2p+ 2P* 4d+ 2D 8 32 4880.75 -0.005 2.234e - 01 [Fe HI] 4881.11 * * A.4- Magellanic Cloud H II Regions 263

-^obs (-^) f(A) /(A ) Ion Ao Mult Lower term Upper term gl g2 4921.62 -0.015 1.073e4-00 He I 4921.93 * * 4924.22 -0.016 6.393e - 02 O II 4924.53 * 4930.97 -0.017 7.829e - 02 [O III] 4931.80 * 4958.53 -0.024 1.170e-t-02 [0 III] 4958.91 *

A.4 Magellanic Cloud H ii Regions

c(H/3)(gal) = 0.087

Table A.20: SMC N66 linelist.

-^obs (Â) f(A) /(A) Ion Aq Mult Lower term Upper term gl g2

3705.43 0.260 2.523e -f 00 H 16 3703.86 * * 3713.54 0.259 1.934e -h 00 H 15 3711.97 * * 0.257 * H 14 3721.94 H14 2p+ 2P* 14d+ 2D 8 * 3723.31 0.257 3.739e 4- 00 [S III] 3721.63 F2 3p2 3P 3p2 IS 3 1 3727.50 0.257 5.026e -t- 01 [O II] 3726.03 FI 2p3 4S* 2p3 2D* 4 4 3730.29 0.256 7.160e + 01 [O II] 3728.82 FI 2p3 4S* 2p3 2D* 4 6 3736.06 0.255 3.096e + 00 H 13 3734.37 H13 2p+ 2P* 13d+ 2D 8 * 3751.70 0.253 2.979e + 00 H 12 3750.15 H12 2p+ 2P* 12d+ 2D 8 * 3772.19 0.249 4.059e + 00 H 11 3770.63 H ll 2p+ 2P* lld+ 2D 8 * 3799.47 0.244 5.430e + 00 H 10 3797.90 HIO 2p+ 2P* 10d+ 2D 8 * 3836.98 0.237 7.627e + 00 H 9 3835.39 H9 2p+ 2P* 9d+ 2D 8 * 3870.25 0.230 3.897e + 01 [Ne III] 3868.75 FI 2p4 3P 2p4 ID 5 5 3890.56 0.226 1.980e + 01 H 8 3889.05 * * 3969.00 0.211 1.078e + 01 [Ne III] 3967.46 FI 2p4 3P 2p4 ID 3 5 3971.61 0.210 1.563e + 01 H 7 3970.07 * 4027.79 0.198 2.476e + 00 He I 4026.21 V18 2p 3P* 5d 3D 9 15 4071.28 0.189 9.736e - 01 [S II] 4068.60 FI 2p3 4S* 2p3 2P* 4 4 4076.17 0.188 1.490e - 01 7 4073.74 * * 0.187 * O II 4075.86 VIO 3p 4D* 3d 4F 8 10 4079.03 0.187 2.968e - 01 [S II] 4076.35 FI 2p3 4S* 2p3 2P* 2 4 4081.53 0.187 6.655e - 02 O II 4078.84 VIO * * 4104.35 0.182 2.676e + 01 H 6 4101.74 H6 2p+ 2P* 6d+ 2D 8 72 4114.66 0.179 5.230e - 02 ? 4112.59 * * 0.178 * O II 4119.22 V20 3p 4P* 3d 4D 6 8 4123.59 0.177 2.602e - 01 He I 4120.84 V16 2p 3P* 5s 38 9 3 A.4. Magellanic Cloud H II Regions 264

Aobs (Â) f(A) /(A) Ion Ao Mult Lower term Upper term gl g2 4128.30 0.176 7.752e - 02 ? 4125.62 * * 4146.25 0.172 3.184e - 01 He I 4143.76 * * 4155.79 0.170 1.116e - 01 O II 4153.30 V19 ** 4171.46 0.167 9.561e - 02 O II 4169.22 V19 * * 4207.77 0.158 1.089c —01 7 4204.93 * * 4342.93 0.127 4.731e + 01 H 5 4340.47 H5 2p+ 2P* 5d+ 2D 8 50 4349.37 0.125 6.083e - 02 S II 4347.20 ** 4358.29 0.123 6.452c —02 ? 4355.74 * * 4361.93 0.122 4.991c - 02 [Fe II] 4359.34 7F * * 4365.70 0.121 6.263c + 00 [O III] 4363.21 F2 2p2 ID 2p2 IS 5 1 4390.47 0.115 4.718c - 01 He I 4387.93 V51 2p IP* 5d ID 3 5 4408.60 0.111 8.526c - 02 [Fe II] 4406.38 * * 4412.07 0.110 1.034c - 01 [Fe II] 4409.85 V55e 3d 4F 4f 2i5i* 8 10 4416.00 0.109 8.522c - 02 [Fe II] 4413.78 * * 4420.38 0.108 8.521c - 02 Fe II] 4418.15 * * 4439.95 0.103 1.136c - 01 He I 4437.55 V50 2p IP* 5s IS 3 1 4474.10 0.095 4.012c+ 00 He I 4471.50 V14 2p 3P* 4d 3D 9 15 4485.62 0.092 6.612c - 02 7 * * * 4538.63 0.080 5.804c —02 7 4535.96 * * 4543.46 0.078 5.136c —02 He II 4541.59 ** 4546.67 0.078 4.595c —02 N III 4544.80 * * 4553.44 0.076 4.624c - 02 7 4550.76 * * 4565.18 0.073 2.093c - 01 Mg I] 4562.60 * * 4573.68 0.071 1.313c - 01 Mg I] 4571.10 * * 4589.44 0.067 5.565c —02 7 4586.75 * * 4593.51 0.066 4.616c —02 O II 4590.97 V15 * * 4597.60 0.065 3.977c - 02 7 4594.90 V15 * * 4598.72 0.065 2.020c - 02 0 II 4596.18 V15 3s’ 2D 3p’ 2F* 4 6 4609.66 0.062 4.290c - 02 [Fe III] 4607.13 3F * * 4633.28 0.057 3.913c - 02 N II 4630.54 V5 * * 4641.96 0.055 3.110c - 02 O II 4638.49 * * 4644.60 0.054 6.372c - 02 N HI 4640.64 V2 3p 2P* 3d 2D 4 6 4650.63 0.052 3.544e - 02 7 4650.63 * * 4652.34 0.052 1.905c - 02 O II 4649.13 VI 3s 4P 3p 4D* 6 8 4653.46 0.051 2.126c - 02 ? 4653.46 * * 4654.05 0.052 2.292c - 02 O II 4650.84 VI 3s 4P 3p 4D* 2 2 4654.68 0.051 7.087c - 03 ? 4654.68 ** 4658.66 0.050 9.368c - 02 C IV WR 4655.92 * * A.4- Magellanic Cloud H ll Regions 265

■^obs (Â) f(A) /(A) Ion Ao Mult Lower term Upper term gl g2 4661.01 0.050 1.655e - 01 [Fe III] 4658.10 3F ** 4664.28 0.049 8.428e - 02 O II 4661.63 VI 3s 4P 3p 4D* 4 4 4672.37 0.047 2.353e - 02 O II 4669.27 * * 4676.84 0.046 2.352e - 02 O II 4673.73 VI 3s 4P 3p 4D* 4 2 0.046 * O II 4676.24 VI 3s 4P 3p 4D* 6 6 4685.79 0.044 3.927e - 02 7 4683.25 * * 4696.16 0.041 3.878e - 02 ? 4693.31 ** 0.041 * O II 4696.35 V I * * 4704.42 0.039 4.407e - 02 [Fe III] 4701.62 3F ** 4714.07 0.037 3.058e - 01 [Ar IV] 4711.37 FI 3p3 48* 3p3 2D* 4 6 4715.87 0.036 4.907e - 01 He I 4713.17 V12 2P 3p* 48 3s 9 3 4724.87 0.034 5.705e - 02 ? 4722.83 ** 4742.99 0.030 2.069e - 01 [Ar IV] 4740.17 FI 3p3 48* 3p3 2D* 4 4 4744.50 0.029 5.164e - 02 O II 4741.71 25 ** 4757.83 0.026 5.805e - 02 [Fe III] 4754.83 3F ** 4760.61 0.025 3.731e - 02 Fe II? 4757.66 * * 4815.82 0.012 4.871e - 02 Si III 4813.20 9 ** 4818.07 0.011 5.866e - 02 S II 4815.45 9 * * 4820.77 0.011 7.900e - 02 7 4818.10 * * 4828.09 0.009 5.953e - 02 7 4825.61 * * 4833.25 0.007 6.606e - 02 7 4830.55 * * 4837.60 0.006 4.089e - 02 7 4834.70 * 4842.50 0.005 5.949e - 02 7 4840.03 ** 4846.88 0.004 5.948e - 02 7 4844.15 * * 0.000 l.OOOe-t- 00 H 4 4861.33 H4 2p-b 2P* 4d-b 2D 8 32 4891.36 -0.007 7.925e - 02 [Fe II] 4889.63 4F ** 4896.47 -0.008 3.248e - 02 7 4893.75 -F ** 4905.43 -0.010 5.579e - 02 [Fe IV] 4903.10 -F ** 4909.17 -0.011 4.436e - 02 O II 4906.83 V28 ** 4914.39 -0.013 4.494e - 02 7 4914.84 * * 4917.57 -0.015 4.492e - 02 ? 4921.93 * * 4924.75 -0.015 1.069e -f 00 He I 4921.93 * * 4927.35 -0.016 7.021e - 02 O II 4924.53 V28 * * 4933.84 -0.017 4.810e - 02 [O III] 4931.80 ** 4937.97 -0.018 3.687e - 02 7 4934.52 * * 4943.12 -0.020 7.024e - 02 O II 4941.07 V33 * * 4965.36 -0.024 1.696e-b 02 [0 HI] 4958.91 ** 5013.46 -0.036 5.088e -b 02 [O III] 5006.84 FI 2p2 3P 2p2 ID 5 5 A.4- Magellanic Cloud H ll Regions 266

-^obs (^) f(A) /(A) Ion Ao Mult Lower term Upper terra gl g2 5525.05 -0.154 3.468e - 01 [Cl III] 5517.66 FI 2p3 4S* 2p3 2D* 4 6 5545.02 -0.158 4.280e - 01 [Cl III] 5537.60 FI 2p3 4S* 2p3 2D* 4 4 5760.62 -0.195 1.261e - 01 [N II] 5754.60 F3 2p2 ID 2p2 IS 5 1 6142.57 -0.256 5.672e - 01 ? 6136.36 * * 6172.14 -0.261 5.667e - 01 [Mn V] 6166.00 * * 6204.09 -0.266 2.440e - 01 O II 6197.92 * * 6237.22 -0.271 2.047e 4- 00 [Ni III] 6231.09 ** 6260.46 -0.275 1.568e-f 00 [Fe II] 6254.30 **

6301.05 -0.282 1.482e 4- 01 [0 I] 6300.34 FI 2p4 3P 2p4 ID 5 5 6312.81 -0.283 1.642e 4- 00 [S III] 6312.10 F3 2p2 ID 2p2 IS 5 1 6467.49 -0.306 1.936e 4- 00 C II 6461.95 ** 6499.73 -0.310 1.934e 4- 00 ? 6492.97 * * 6529.90 -0.313 2.118e + 00 ?WR 6516.12 ** 6543.22 -0.315 1.918e 4- 00 7 6529.41 * *

6561.94 -0.318 2.768e 4- 00 [N II] 6548.10 FI 2p2 3P 2p2 ID 3 5 6574.48 -0.320 2.861e4-02 H 3 6562.77 H3 2p4- 2P* 3d4- 2D 8 18 6593.61 -0.323 3.824e 4- 00 [N II] 6583.50 FI 2p2 3P 2p2 ID 5 5 6683.76 -0.336 2.970e 4- 00 He I 6678.16 V46 2p IP* 3d ID 3 5 6720.72 -0.342 7.896e 4- 00 [S II] 6716.44 F2 2p3 4S* 2p3 2D* 4 6 6735.11 -0.344 5.784e 4- 00 [S II] 6730.82 F2 2p3 4S* 2p3 2D* 4 4 7072.20 -0.387 2.281e 4- 00 He I 7065.25 VIO 2p 3P* 3s 3S 9 3 7144.83 -0.396 8.945e 4- 00 [Ar III] 7135.80 FI 3p4 3P 3p4 ID 5 5 7166.84 -0.399 1.285e 4- 00 ? 7159.07 * * 7178.54 -0.388 3.344e - 01 [Ar IV] 7070.62 * * 7186.58 -0.401 2.307e - 01 He II 7177.50 * * 7247.76 -0.408 7.428e 4- 00 C II 7231.32 * * 7252.87 -0.408 7.427e 4- 00 C II 7236.42 * * 7293.51 -0.414 5.275e 4- 00 He I 7281.35 V45 2p IP* 3s IS 3 1 7333.56 -0.418 1.772e 4- 00 [0 II] 7318.92 F2 2p3 2D* 2p3 2P* 6 2 7344.32 -0.420 1.553e 4- 00 [O II] 7329.67 F2 2p3 2D* 2p3 2P* 4 2 7540.02 -0.443 1.861e4-00 [Cl IV] 7529.90 * * * *

c(H^)(gal) = 0.073

Table A.21: LMC N llB linelist.

Aobs (^) f(A) /(A) Ion Aq Mult Lower term Upper term gl g2

3729.14 0.257 8.273e -hOl [O II] 3726.03 FI 2p3 4S* 2p3 2D* 4 4 3731.94 0.256 1.102e 4- 02 [O II] 3728.82 FI 2p3 4S* 2p3 2D* 4 6 0.255 * H 13 3734.37 H13 2p+ 2P* 13d+ 2D 8 * 3872.02 0.230 2.056e 4- 01 [Ne III] 3868.75 FI 2p4 3P 2p4 ID 5 5 A.J^. Magellanic Cloud H ll Regions 267

-^obs (À) f(A) /(A) Ion Aq Mult Lower term Upper term gl g2 3892.03 0.226 2.234e + 01 He I 3888.65 V2 2s 3S 3p 3P* 3 9 3970.38 0.211 3.502e + 00 [Ne III] 3967.46 FI 2p4 3P 2p4 ID 3 5 3973.23 0.210 1.400e + 01 H 7 3970.07 * * 4072.97 0.189 8.647e - 01 [S II] 4068.60 FI 2p3 4S* 2p3 2P* 4 4 4074.13 0.189 3.171e - 0 1 O II 4069.62 VIO 3p 4D* 3d 4F 2 4 0.189 * O II 4069.89 VIO 3p 4D* 3d 4F 4 6 4076.53 0.188 1.027e - 01 O II 4072.16 VIO 3p 4D* 3d 4F 6 8 0.187 * 0 II 4075.86 VIO 3p 4D* 3d 4F 8 10 4080.73 0.187 5.055e — 01 [S II] 4076.35 FI 2p3 4S* 2p3 2P* 2 4 4083.22 0.187 1.670e - 01 O II 4078.84 VIO * * 4086.71 0.185 1.412e - 01 ? 4086.71 ** 4088.71 0.186 1.545e - 01 O II 4083.90 V48b ** 4089.92 0.185 1.092e - 01 O II 4085.11 VIO 3p 4D* 3d 4F 6 6 4091.97 0.185 1.138e - 01 O II 4087.15 V48c 3d 4F 4f G3* 4 6 4094.11 0.184 1.305e - 01 O II 4089.29 V48a 3d 4F 4f G5* 10 12 4106.02 0.182 2.771e + 01 H 6 4101.74 H6 2p+ 2P* 6d + 2D 8 72 4114.41 0.180 1.564e - 01 0 II 4110.78 V20 * * 4116.86 0.178 1.492e - 01 7 4116.86 * * 0.178 * O II 4119.22 V20 3p 4P* 3d 4D 6 8 4125.03 0.177 2.827e - 01 He I 4120.84 V16 2p 3P* 5s 3S 9 3 4135.89 0.176 1.112e - 01 O II 4129.32 V19 ** 4139.38 0.175 1.363e - 01 O II 4132.80 V19 * * 4147.96 0.172 3.163e - 01 He I 4143.76 * * 4157.64 0.170 9.819e - 02 O II 4153.30 V19 ** 4160.87 0.169 1.280e - 01 O II 4156.53 V19 ** 4164.34 0.168 2.373e - 01 ? 4164.34 * * 4169.09 0.167 2.357e - 01 ? 4169.09 * * 0.167 * O II 4169.22 V19 ** 4179.89 0.165 8.859e - 02 N II 4176.16 ** 4190.33 0.162 8.877e - 02 C HI 4186.90 V18 ** 4193.23 0.162 1.555e - 01 O II 4189.79 V36 ** 4207.42 0.158 1.328e - 01 ? 4207.42 ** 4212.17 0.157 4.007e - 01 ? 4212.17 ** 4218.52 0.155 1.184e - 01 ? 4218.52 ** 4221.32 0.155 1.798e - 01 ? 4221.32 * * 4241.04 0.151 9.933e - 02 N II 4237.05 * * 4245.50 0.150 2.956e - 01 N II 4241.78 V48a 3d 3D* 4f li3i. 5 7 4252.40 0.148 2.088e - 01 ? 4248.15 * * A.4- Magellanic Cloud H II Regions 268

-^obs (^) f(A) /(A) Ion Ao Mult Lower term Upper term gl g2 4257.04 0.147 2.641e - 01 ? 4252.73 * * 4263.04 0.145 1.996e - 01 ? 4263.04 * * 4265.80 0.144 2.604e - 01 ? 4265.80 * * 4268.90 0.144 1.959e - 01 C II 4267.15 V6 3d 2D 4f 2F* 10 14 4277.87 0.142 4.264e - 01 0 II 4276.28 V67b * 4282.91 0.141 2.678e - 01 O II 4281.32 V53b ** 4287.28 0.140 1.097e - 01 O II 4285.69 V78b 3d 2F 4f F3* 6 8 0.139 * O II 4288.82 V53c 3d 4P 4f Dl* 2 4 4292.84 0.138 2.795e - 01 0 II 4291.25 V55 3d 4P 4f G3* 6 8 4316.78 0.133 3.879e - 01 0 II 4312.11 V78a 3d 2F 4f F4* 8 8 4319.54 0.133 * O II 4313.44 V78a 3d 2F 4f F4* 8 10 4320.07 0.133 8.002e - 01 0 II 4315.40 V63c * * 4318.80 0.132 1.345e - 01 O II 4317.14 V2 3s 4P 3p 4P* 2 4 4321.01 0.132 3.418e - 01 O II 4319.63 V2 3s 4P 3p 4P* 4 6 4329.52 0.130 2.924e - 01 0 II 4325.76 V2 3s 4P 3p 4P* 2 2 4332.30 0.129 1.934e - 01 7 4329.08 * * 4334.89 0.129 2.887e - 01 O II 4331.13 V65b * * 4342.73 0.127 4.679e + 01 H 5 4340.47 H5 2p+ 2P* 5d+ 2D 8 50 0.125 * O II 4345.56 V2 3s 4P 3p 4P* 4 2 4353.52 0.125 1.342e - 01 O II 4349.43 V2 3s 4P 3p 4P* 6 6 4357.68 0.124 1.345e - 01 O II 4353.59 V76c 3d 2F 41 G3* 6 8 4367.56 0.121 1.689e + 00 [O III] 4363.21 F2 2p2 ID 2p2 IS 5 1 4392.35 0.115 4.952e - 01 He I 4387.93 V51 2p IP* 5d ID 3 5 4412.04 0.110 2.045e - 01 ? 4412.00 * * 4419.80 0.109 2.227e - 01 O II 4414.90 V5 3s 2P 3p 2D* 4 6 0.109 * O II 4416.97 V5 3s 2P 3p 2D* 2 4 4475.98 0.095 4.313e + 00 He I 4471.50 V14 2p 3P* 4d 3D 9 15 4532.10 0.079 4.202e - 02 ? 4541.25 * * 4537.53 0.079 8.927e - 02 ? 4537.54 * * 4539.73 0.080 5.857e - 02 N HI 4534.58 V3 * * 4542.95 0.078 7.974e - 02 ? 4542.95 * * 4546.93 0.078 7.873e - 02 He II 4541.59 ** 4560.15 0.075 1.510e - 01 [Fe II] 4555.00 * * 4561.54 0.075 2.631e - 02 Fe II 4556.39 * * 4567.55 0.073 1.952e - 01 Mg I] 4562.60 * * 4576.06 0.071 2.038e - 01 Mg I] 4571.10 * * 4589.35 0.068 5.993e - 02 ? 4584.40 * * 4609.38 0.063 4.693e - 02 N V? 4604.43 ** A.J^. Magellanic Cloud H u Regions 269

■^obs (Â) f(A) /(A) Ion Ao Mult Lower term Upper term gl g2 0.062 * [Fe III] 4607.13 3F ** 4643.65 0.055 1.231e - 01 O II 4638.49 * * 4646.29 0.054 2.758e - 02 N III 4640.64 V2 3p 2P* 3d 2D 4 6 4654.05 0.052 8.046e - 02 O II 4649.13 VI 3s 4P 3p 4D* 6 8 4655.76 0.052 8.045e - 02 O II 4650.84 VI 3s 4P 3p 4D* 2 2 4662.81 0.050 2.721e - 01 [Fe III] 4658.10 3F ** 4665.73 0.049 7.111e - 02 0 II 4661.63 VI 3s 4P 3p 4D* 4 4 4670.96 0.048 2.565e - 02 [Fe III] 4667.00 3F * * 4677.69 0.046 2.446e —02 O II 4673.73 VI 3s 4P 3p 4D* 4 2 0.046 * 0 II 4676.24 VI 3s 4P 3p 4D* 6 6 4686.08 0.044 5.043e - 02 ? 4681.44 * * 4695.95 0.042 8.205e - 02 ? 4691.31 * * 4700.84 0.041 3.461e - 02 O II 4696.35 V I * * 4706.12 0.039 9.087e - 02 [Fe III] 4701.62 3F ** 4711.25 0.038 2.616e - 02 ? 4706.60 ** 0.037 * [Ar IV] 4711.37 FI 3p3 4S* 3p3 2D* 4 6 4717.81 0.036 5.095e - 01 He I 4713.17 V12 2P 3p* 4S 3s 9 3 4724.58 0.035 6.233e - 02 7 4719.94 ** 4732.73 0.033 1.844e - 02 ? 4728.09 ** 4744.92 0.030 1.132e - 01 [Ar IV] 4740.17 FI 3p3 4S* 3p3 2D* 4 4 4749.76 0.029 2.430e - 02 [Fe II] 4745.48 ** 4759.60 0.026 8.774e - 02 [Fe III] 4754.83 3F ** 4774.42 0.022 3.947e - 02 [Fe III] 4769.60 3F ** 4779.56 0.021 2.623e - 02 [Fe II] 4774.74 20F ** 4795.55 0.016 7.198e - 02 7 * * * 4803.51 0.015 2.186e - 02 ? 4798.87 * * 4815.88 0.011 6.538e - 02 ? 4815.88 * * 4819.93 0.011 9.985e - 02 [Fe II] 4814.55 20F ** 1.000 l.OOOe -f 00 H 4 4861.33 H4 2p+ 2P* 4d-t- 2D 8 32 4885.52 -0.005 5.327e - 02 [Fe III] 4881.11 2F ** 4905.47 -0.010 2.917e - 02 [Fe IV]? 4900.50 -F ** 4926.64 -0.015 1.079e4-00 He I 4921.93 ** 4929.24 -0.016 1.260e - 01 O II 4924.53 V28 ** 4935.89 -0.017 7.424e - 02 [0 HI] 4931.80 ** 4968.01 -0.024 1.099e-h02 [O HI] 4958.91 **

5016.07 -0.036 3.301e 4- 02 [O HI] 5006.84 FI 2p2 3P 2p2 ID 5 5 5052.34 -0.047 1.315e - 01 7 5052.34 ** 5060.59 -0.049 4.696e —01 7 5060.59 * * A.4- Magellanic Cloud H ll Regions 270

-^obs (-^) f(A) /(A) Ion Ao Mult Lower term Upper term gl g2 5526.40 -0.154 4.944e - 01 [Cl III] 5517.66 FI 2p3 4S* 2p3 2D* 4 6 5546.37 -0.158 4.811e - 01 [Cl III] 5537.60 FI 2p3 4S* 2p3 2D* 4 4 5752.46 -0.194 1.198e - 01 ? 5752.46 ** 5763.87 -0.195 2.028e - 01 [N II] 5754.60 F3 2p2 ID 2p2 IS 5 1 5883.90 -0.215 1.471e -f 01 He I 5875.66 V l l 2p 3P* 3d 3D 9 15 6564.36 -0.318 5.792e -f 00 [N II] 6548.10 FI 2p2 3P 2p2 ID 3 5 6578.06 -0.320 3.167e + 02 H 3 6562.77 H3 2p+ 2P* 3d+ 2D 8 18 6597.40 -0.323 1.658e-k01 [N II] 6583.50 FI 2p2 3P 2p2 ID 5 5 6687.50 -0.336 3.681e + 00 He I 6678.16 V46 2p IP* 3d ID 3 5 6724.72 -0.342 1.422e-f 01 [S II] 6716.44 F2 2p3 4S* 2p3 2D* 4 6 6738.76 -0.344 1.055e + 01 [S II] 6730.82 F2 2p3 4S* 2p3 2D* 4 4 7076.35 -0.387 2.324e -t- 00 He I 7065.25 VIO 2p 3P* 3s 3S 9 3 7149.04 -0.396 1.159e4-01 [Ar III] 7135.80 FI 3p4 3P 3p4 ID 5 5 7158.08 -0.399 2.498e - 01 ? ** * 7248.81 -0.408 2.101e + 00 C II 7231.32 * * 7253.99 -0.408 2.1596 -f- 00 C II 7236.42 ** 7294.45 -0.414 1.7126 + 00 He I 7281.35 V45 2p IP* 3s IS 3 1 7338.55 -0.418 4.2896 + 00 [O II] 7318.92 F2 2p3 2D* 2p3 2P* 6 2 7350.41 -0.420 5.8406 + 00 [O II] 7329.67 F2 2p3 2D* 2p3 2P* 4 2

c(H/3)(gal) = 0.087, c(H/?)(lmc) = 0.32

Table A.22: 30 Doradus linelist.

-^obs (-^) f(A) /(A) Ion Ao Mult Lower term Upper term gl g2

3729.15 0.375 5.0226 + 01 [O II] 3726.03 FI 2p3 4S* 2p3 2D* 4 4 3731.94 0.374 5.5476 + 01 [O II] 3728.82 FI 2p3 4S* 2p3 2D* 4 6 3753.37 0.365 3.5236 + 00 H 12 3750.15 H12 2p+ 2P* 12d+ 2D 8 * 3773.87 0.356 4.2826 + 00 H 11 3770.63 H ll 2p+ 2P* lld + 2D 8 * 3801.17 0.344 5.6176 + 00 H 10 3797.90 HIO 2p+ 2P* 10d+ 2D 8 * 3822.94 0.335 1.1866 + 00 He I 3819.62 ** 3838.69 0.329 8.2606 + 00 H 9 3835.39 H9 2p+ 2P* 9d+ 2D 8 * 3871.99 0.315 3.8786 + 01 [Ne III] 3868.75 FI 2p4 3P 2p4 ID 5 5 3892.31 0.306 1.9476 + 01 He I 3888.65 V2 2s 3S 3p 3P* 3 9 3970.79 0.276 1.1496 + 01 [Ne III] 3967.46 FI 2p4 3P 2p4 ID 3 5 3973.40 0.274 1.7196 + 01 H 7 3970.07 ** 4072.50 0.237 5.6166 - 01 [S II] 4068.60 FI 2p3 4S* 2p3 2P* 4 4 4073.66 0.237 1.3286 - 01 O II 4069.62 VIO 3p 4D* 3d 4F 2 4

0.237 * O II 4069.89 VIO 3p 4D* 3d 4F 4 6 4076.06 0.236 6.8286 - 02 O II 4072.16 VIO 3p 4D* 3d 4F 6 8 A.^. Magellanic Cloud H II Regions 271

-^obs (Â) f(A) /(A) Ion Ao Mult Lower term Upper term gl g2 0.235 * O II 4075.86 VIO 3p 4D* 3d 4F 8 10 4080.26 0.235 2.270e - 01 [S II] 4076.35 FI 2p3 4S* 2p3 2P* 2 4 4082.75 0.234 3.243e - 02 O II 4078.84 VIO * * 4087.98 0.232 1.081e - 02 O II 4083.90 V48b * * 4089.20 0.232 2.576e - 02 0 II 4085.11 VIO 3p 4D* 3d 4F 6 6 4091.24 0.231 1.852e - 02 O II 4087.15 V48c 3d 4F 4f G3* 4 6 4093.38 0.230 2.103e - 02 O II 4089.29 V48a 3d 4F 4f G5* 10 12 4105.56 0.225 2.602e + 01 H 6 4101.74 H6 2p+ 2P* 6d + 2D 8 72 4114.32 0.222 2.404e - 02 O II 4110.78 V20 * * 0.219 * O II 4119.22 V20 3p 4P* 3d 4D 6 8 4124.62 0.219 2.089e - 01 He I 4120.84 V16 2p 3P* 5s 3S 9 3 4136.45 0.214 5.082e - 02 O II 4132.80 V19 ** 4147.48 0.211 2.942e - 01 He I 4143.76 * * 4156.82 0.207 6.860e - 02 O II 4153.30 V19 ** 4160.05 0.206 2.510e - 02 O II 4156.53 V19 * * 4172.75 0.202 6.664e - 02 0 II 4169.22 V19 * * 4270.73 0.169 9.194e - 02 C II 4267.15 V6 3d 2D 4f 2F* 10 14 0.166 7.849e - 02 O II 4275.55 V67a 3d 4D 4f F4* 8 10 0.161 1.091e - 01 O II 4291.25 V55 3d 4P 4f G3* 6 8 4307.40 0.157 5.188e - 02 O II 4303.61 V65a 3d 4D 4f G5* 8 10 4314.99 0.154 2.967e - 02 O II 4312.11 V78a 3d 2F 4f F4* 8 8 4315.49 0.154 4.326e - 02 O II 4313.44 V78a 3d 2F 4f F4* 8 10 4317.45 0.153 3.833e - 02 O II 4315.40 V63c ** 4319.19 0.153 5.334e - 02 O II 4317.14 V2 3s 4P 3p 4P* 2 4 4321.68 0.152 8.763e - 02 O II 4319.63 V2 3s 4P 3p 4P* 4 6 4344.21 0.145 4.848e + 01 H 5 4340.47 H5 2p+ 2P* 5 d + 2D 8 50 4367.00 0.138 3.3576 -|- 00 [O III] 4363.21 F2 2p2 ID 2p2 IS 5 1 4391.79 0.130 5.366e - 01 He I 4387.93 V51 2p IP* 5d ID 3 5 4412.55 0.124 3.105e - 02 Ne II 4409.30 V55e 3d 4F 4f 2j5i* 8 10 4418.15 0.122 3.316e - 02 O II 4414.90 V5 3s 2P 3p 2D* 4 6 4420.23 0.121 4.359e - 02 O II 4416.97 V5 3s 2P 3p 2D* 2 4 4441.41 0.115 6.223e - 02 He I 4437.55 V50 2p IP* 5s IS 3 1 4475.44 0.105 4.359e + 00 He I 4471.50 V14 2p 3P* 4d 3D 9 15 4566.56 0.078 5.060e - 02 Mg I] 4562.60 3s2 IS 3s3p 3P* 1 5 4575.32 0.076 5.709e - 02 Mg I] 4571.10 3s2 IS 3s3p 3P* 1 3 4611.30 0.066 5.069e - 02 [Fe III] 4607.13 3F * * 4613.96 0.065 2.094e - 02 0 II 4609.44 V92a 3d 2D 4f F4* 6 8 4638.92 0.059 1.014e - 02 N III 4634.14 V2 * * A.4- Magellanic Cloud H II Regions 272

-^obs (Â) f(A) /(A) Ion Ao Mult Lower term Upper term gl g2 4643.03 0.057 6.404e - 02 O II 4638.86 VI 3s 4P 3p 4D* 2 4 4645.99 0.056 8.740e - 02 O II 4641.81 VI 3s 4P 3p 4D* 4 6 4653.38 0.054 6.416e - 02 O II 4649.13 VI 3s 4P 3p 4D* 6 8 4655.09 0.054 6.413e - 02 O II 4650.84 V I 3s 4P 3p 4D* 2 2 4662.27 0.052 6.251e - 01 [Fe III] 4658.10 3F * * 4665.18 0.051 5.652e - 02 O II 4661.63 VI 3s 4P 3p 4D* 4 4 * 4671.08 0.050 2.353e - 02 [Fe III] 4667.00 3F * 4677.91 0.048 1.202e - 02 O II 4673.73 VI 3s 4P 3p 4D* 4 2 4680.42 0.047 2.235e - 02 O II 4676.24 V I 3s 4P 3p 4D* 6 6

* 4689.06 0.045 1.715e - 02 He II 4685.68 3.4 * 4692.90 0.043 1.125e - 02 7 4692.90 * *

4696.27 0.042 1.772e - 02 ? 4696.27 * * 4705.74 0.040 1.885e - 01 [Fe III] 4701.62 3F * * 4715.54 0.038 2.103e - 01 [Ar IV] 4711.37 FI 3p3 4S* 3p3 2D* 4 6 4717.34 0.037 4.805e - 01 He I 4713.17 V12 2P 3p* 4S 3s 9 3

4738.18 0.032 6.316e - 02 [Fe III] 4733.93 3F ** 4744.40 0.030 1.816e - 01 [Ar IV] 4740.17 FI 3p3 4S* 3p3 2D* 4 4 * 4753.57 0.027 1.915e - 02 ? * * * 4758.89 0.026 1 .3 8 4 e -0 1 [Fe HI] 4754.83 3F * * 4773.70 0.022 6.790e - 02 [Fe HI] 4769.60 3F * * 4781.58 0.020 3.210e - 02 [Fe III] 4777.88 3F * ** 4800.53 0.015 1.8786 - 02 ? * * 4819.44 0.011 3.0746 - 02 [Fe II] 4814.55 20F *

4865.50 -0.000 9.9976 + 01 H 4 4861.33 H4 2p+ 2P* 4d-f 2D 8 32 4885.14 -0.005 1.8316 - 01 [Fe III] 4881.11 2F * * 4893.48 -0.008 1.7116 - 02 [Fe II] 4889.63 4F ** * 4907.18 -0.011 3.3526 — 02 [Fe IV] 4903.50 -F * 4910.38 -0.012 3.6946 - 02 O II 4906.83 * * 4926.27 -0.015 1.1476-h 00 He I 4921.93 * * 4935.46 -0.018 6.9436 - 02 [O III] 4931.80 * * 4967.44 -0.024 1.687e -t- 02 [O III] 4958.91 * *

5015.56 -0.035 5.0506 + 02 [0 III] 5006.84 FI 2p2 3P 2p2 ID 5 5 5279.65 -0.094 2.791e - 01 [Fe II] 5273.38 18F * * 5315.59 -0.102 1.5886 - 01 7 * * *

5526.27 -0.140 4.7646 - 01 [C l III] 5517.66 FI 2p3 4S* 2p3 2D* 4 6

5546.24 -0.143 3.6946 - 01 [C l HI] 5537.60 FI 2p3 4S* 2p3 2D* 4 4 5762.29 -0.181 2.0646 - 01 [N II] 5754.60 F3 2p2 ID 2p2 IS 5 1

5883.20 -0.201 1.2796 + 01 He I 5875.66 V l l 2p 3P* 3d 3D 9 15 A.4- Magellanic Cloud H II Regions 273

-^obs (^ ) %A) /(A ) Ion Ao Mult Lower term Upper term gi g2 6238.42 -0.259 1.454e - 01 [Ni III] 6231.09 * * 6261.66 -0.262 1.264e-01 [Fe II] 6254.30 * * 6306.34 -0.269 1.070e 4- 00 [OI] 6300.34 FI 2p4 3P 2p4 ID 5 5 6318.12 -0.271 1.945e -h 00 [S III] 6312.10 F3 2p2 ID 2p2 IS 5 1 6562.58 -0.307 3.226e-t-00 [N II] 6548.10 FI 2p2 3P 2p2 ID 3 5 6568.93 -0.309 3.030e + 02 H 3 6562.77 H3 2p+ 2P* 3d+ 2D 8 18 6595.37 -0.312 8.7356 -|- 00 [N II] 6583.50 FI 2p2 3P 2p2 ID 5 5 6685.50 -0.325 3.6536-1- 00 He I 6678.16 V46 2p IP* 3d ID 3 5 6722.67 -0.330 6.0096 4- 00 [S II] 6716.44 F2 2p3 4S* 2p3 2D* 4 6 6736.74 -0.333 5.4326 4- 00 [S II] 6730.82 F2 2p3 4S* 2p3 2D* 4 4 7074.23 -0.377 3.5556 4- 00 He I 7065.25 VIO 2p 3P* 3s 3S 9 3 7146.92 -0.386 1.3586-KOI [Ar III] 7135.80 FI 3p4 3P 3p4 ID 5 5 7166.65 -0.389 3.4946 - 02 [Fe II] 7155.14 14F * * 7247.38 -0.398 2.4356 - 01 C II 7231.32 * * 7252.50 -0.399 2.5466 - 01 C II 7236.42 * * 7296.23 -0.404 7.3856 - 01 He I 7281.35 V45 2p IP* 3s IS 3 1 7335.88 -0.409 2.0726 4- 00 [O II] 7318.92 F2 2p3 2D* 2p3 2P* 6 2 7346.65 -0.410 1.7016-h 00 [O II] 7329.67 F2 2p3 2D* 2p3 2P* 4 2 B ibliography

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