Extragalactic Near Infrared Spectroscopy C>£ C-A-u-AV.f£C, by
Andrea Heather Prestwich
A Thesis submitted for the Degree of Doctor of Philospphy of the University of London
October 1989
Astrophysics Group Blackett Laboratory Imperial College of Science, Technology and Medicine London SW7 2BZ
1 Abstract
This thesis is devoted to near infrared spectroscopy of starburst and active galaxies. The introductory chapter describes the transitions seen in the H and K windows, concentrating on infrared recombination lines, the rotation-vibration spectrum of molecular hydrogen andforbiddeniron transitions. The excitation mechanisms for these transitions are reviewed. An overview of galactic and extragalactic infrared spectroscopy is given. Chapter 2 gives a detailed description of the data reduction techniques employed, especially spectral calibration.
Chapter 3 is devoted to infrared spectroscopy of merging galaxies. A simple starburst model is used to show that the By recombination line fluxes in these galaxies are consistent with star formation. Multiaperture spectroscopy of the S( 1) and [Fell] lines is presented to show that the spatial extent of these shock excited lines is much larger than t^e infrared emission, suggesting that the shocks may be induced by the merging process. Chapter 4 is devoted to simple line mapping of the interacting system Arp 299, where it is demonstrated that the bulk of the luminosity in this system can be understood in terms of star form ation.
Chapter 5 focuses on two Seyfert galaxies, NGC 1068 and NGC 4151. Multiaperture spectroscopy is presented for both of these galaxies. It is suggested that the By recombination line arises in clouds photoionized by the active nucleus in both of these galaxies, in contrast to the interacting and merging galaxies discussed in Chapters 3 and 4. Comparison of large aperture spectra of NGC 1068 with the smaller aperture spectra by previous workers suggests that there may be spatially extended fluorescently excited molecular hydrogen in this galaxy.
In Chapter 6 non-parametric partial correlation analysis is used to investigate the relationship between the By recombination line flux, the 1 -0S( 1) line flux of molecular hydrogen and theinfrared continuum fluxes measured by the Infrared Astronomical Satellite (IRAS). It was found that the By recombination line is well correlated with the 12 and 25 pm emission in starburst galaxies, but is not significantly correlated at 60 or 100 pm. This result can be understood if the 12 and 25 pm emission arises in a centrally concentrated starburst, whereas the longer wavelength emission has an extended component.
2 In contrast to these results, the By recombination line flux is well correlated with all four IRAS bands in Seyfert galaxies. This correlation implies that the infrared emission is centrally concentrated, and is associated with the active nucleus. Finally, the main results of this thesis are reviewed and suggestions given for further work.
<4
3 Contents
Abstract 2
Figures 8
Tables 13
Acknowledgments 15
Chapter 1: Astronomical Near Infrared Spectroscopy
1.1 Introduction • 16 1.2 Hydrogen Recombination Lines 19 1.2.1 Derivation of Hydrogen Line Intensities 19 1.2.2 Derivation of Nebular Quantities 21 1.3 Rotation-Vibration Spectrum of Molecular Hydrogen 23 1.3.1 Ultraviolet Fluorescence and Heating 27 1.3.2 Thermal Excitation Behind a Shock Front 29 1.3.3 Observational Diagnostics 30 1.4 Forbidden Iron Transitions 31 1.5 Observations of Molecular Hydrogen Emission from Galactic Sources 32 1.5.1 Molecular Hydrogen in the Orion Star Formation Region 35 1.5.2 Molecular Hydrogen Emission from IC 443 36 1.6 Infrared Line and Continuum Emission from HII Regions 37 1.6.1 Radio Continuum Emission 38 1.6.2 Infrared Continuum Emission from HII Regions 39 1.6.3 Infrared Recombination Lines from HII Regions 43 1.7 Observations of [Fell] from Galactic Sources 46 1.8 Extragalactic Near Infrared Spectroscopy 47 1.8.1 Extragalactic Infrared Recombination Lines 49 1.8.2 Excitation Mechanisms for [Fell] and H 2 52 1.8.3 Spatial Extent of Infrared Lines 54 1.8.4 Infrared Linewidths 54 1.8.5 Overview 55
4 1.9 Objectives 56
Chapter 2: Observational and Data Reduction Techniques in Infrared Spectroscopy
2.1 Circular Variable Filter (CVF) Spectroscopy 57 2.2 Observing Techniques 59 2.3 Spectral Calibration 63 2.3.1 Overall Strategy 6 3 2.3.2 Stellar Absorption Features 65 2.3.3 Instrumental Effects 79 2.4 Flux Calibration and Calculation of Errors 83
Chapter 3:Near Infrared Spectroscopy of Merging Galaxies
3.1 Introduction v 88 3.2 Data and Sample Selection 89 3.3 Brackett Line Spectroscopy of Mergers 105 3.3.1 The Starburst Model 105 3.3.2 Extinction Corrections to By Line 112 3.3.3 Results and Discussion 114 3.4 Shocked Gas in Merging Galaxies 121 3.4.1 Excitation Mechanisms 121 3.4.2 Spatial Distribution of Shocked Gas 123 3.4.3 Interaction Induced Shocks 128 3.4.4 The Distribution of Molecular Gas in Arp 220 129 3.5 Summary and Conclusions 131
Chapter 4: Star Formation in Arp 299 - a Case Study
4.1 Arp 299 ' 132 4.2 The Data 133 4.3 Excitation Mechanisms * 140 4.4 Extinction 143 4.5 Star Formation in NGC 3690 144 4.5.1 Thermal Radio Emission 144 4.5.2 Infrared Emission 145 4.6 Star Formation in IC 694 148
5 4.6.1 Thermal Radio Fluxes 148 4.6.2 Spatial Extent of the By Emission 148 4.7 Spatial Distribution of Shocked Gas 151 4.7.1 Comparison with CO Distribution 151 4.7.2 S( 1) Line Distribution at Source A 152 4.8 Summary and Conclusions 153
Chapter 5: Near Infrared Spectroscopy of Seyfert Galaxies
5.1 Introduction 153 5.2 Near Infrared Spectroscopy of NGC 1068 157 " 5.2.1 NGC 1068 -Seyfert and Starburst 157 5.2.2 The Data - 158 5.2.3 Comparison with Previous Work 165 5.2.4 Spatial Extent of the Infrared Lines 166 5.2.5 Excitation Mechanisms 170 5.2.6 Origin of the Extended By Emission 171 5.2.7 Distribution and Excitation of the 1 -0S( 1) and [Fell] Lines 174 5.3 Near Infrared Spectroscopy of NGC 4151- EvidenceforBrackettLineVariability 174 5.3.1 Variability inNGC4151 175 5.3.2TheData 175 5.3.3 Evidence for Variability 180 5.3.4 Implications for Origin and Excitation of Infrared Lines 184 5.4 Summary and Conclusions 184 Chapter 6: A Statistical Study of Infrared Lines in Galaxies
6.1 Introduction 186 6.2 Data and Statistical Analysis 186 6.2.1 The Sample 186 6.2.2 Biases due to Distance Effects 193 6.2.3 The Analysis 201 6.2.4 Comparison of Partial Correlation Analysis with Restricted Redshift Range Analysis 204 6.3 Starburst Galaxies 210 6.3.1 Results 210
6 6.3.2 Brackett line Correlations 211 6.3.3 S(l) Line Correlations 224 6.4 Active Galaxies 224 6.4.1 Results 224 6.4.2 Brackett Line Correlations 226 6.4.3 S( 1) Line Correlations 228 6.5 Summary 228
Chapter 7: Summary, Conclusions and Suggestions for Further Work
7 1" Summary 230 7.2 Matters Arising - UV Fluorescence in Spiral Galaxies? 231 7.3 Suggestions for Further Work 235 7.3.1 Line Imaging 235 7.3.2 Future Spectroscopjc Observations 236 7.4 Concluding Remarks 237
Appendix 1: Starburst Parameters
A 1.1 The Initial Mass Function 238 A1.2 Mass-Luminosity Law 239 A1.3 Mass Main-Sequence Lifetime 240 A1.4 Mass-Ionising Photon Flux 241
References 243
7 Figures
Chapter 1
1.1 Main Transitions in the H and K windows. 18 1.2 Schematic diagram showing potential energy versus intemuclear distance for a diatomic molecule. 25 1.3 Comparison of fluorescent and thermal emission. 31 1.4a [Fell] line ratios as a function of temperature. 3 3 1.4b [Fell] line ratios as a function of electron density. 34 1.5 Continuum Energy distribution of an HE region. 39 1.6 Ionising Photon Flux versus Stellar Spectral Type. 40 1.7 Ionising photon flux versus stellar luminosity HII regions. 42 1.8 Emission measure versus luminosity for ordinary galactic HII regions and line excess objects. 45
Chapter 2
2.1 Design of circular variable filter, taken from Yen (1969). 58 2.2 Schematic diagram of infrared line and continuum. 59 2.3 Schematic diagram of a Cassegrain system. 61 2.4 Atmospheric transmission in the H and K windows. 63 2.5 Schematic diagram showing effect of stellar absorption features. 66 2.6 H-window spectrum of NGC 6240, ratioed with standard SAO 121790. 67 2.7 H-window spectrum of SAO 121790. 69 2.8 H-window spectrum SAO 191840. 70 2.9 H-window spectrum of NGC 6240, ratioed with the "syntheticatmosphere" 72 2.10 [Fell] line flux plotted as a function of aperture for NGC 6240 ratioed with SAO121790 and thesyntheticatmosphere. * 74 2.11 K-window spectrum of NGC 1068, ratioed with standard SAO 129959. 76 2.12 K-window spectrum of SAO 129959 77 2.13 Final reduced K-window spectrum of NGC 1068, with stellar absorption features removed. 7 8
8 2.14 K-window spectrum of NGC 3256 80 2.15 Ratio of two spectra of the standard star BS 4013, takenaboutanhourapart. 81 2.16 Ratio of two observations of BS 4013, with wavelength scale aligned. 82 2.17 K-window spectrum of NGC 6240 with fitted polynomial. 85 2.18 NGC 6240, continuum subtracted. 86 2.19 Gaussian fit to the NGC 6240 S(l) line, with residuals. 87
Chapter 3.
3. la(i) 19.6 arcsecond K-window spectrum of NGC 6240. 92 3. la(ii) 7.8 arcsecond K-window spectrum of NGC 6240. 93 3. la(iii) 19.6 arcsecond H-window spectrum of NGC 6240. 94 3. la(iv) 12.4 arcsecond H-window spectrum of NGC 6240. 95 3. la(v) 7.8 arcsecond H-window ^spectrum of NGC 6240. 96 3. la(vi) 5.5 arcsecond H-window spectrum of NGC 6240. 97 3. lb(i) 19.6 arcsecond K-window spectrum of Arp 220. 98 3. lb(ii) 7.8 arcsecond K-window spectrum of Arp 220. 99 3. lc(i) 19.6 arcsecond K-window spectrum of NGC 1614. 100 3. lc(ii) 19.6 arcsecond H-window spectrum of NGC 1614. 101 3. lc(iii) 12.4 arcsecond H-window spectrum of NGC 1614. 102 3. lc(iv) 7.8 arcsecond H-window spectrum of NGC 1614. 103 3. ld(i) 19.6 arcsecond K-window spectrum of NGC 6052. 104 3.2 By flux plotted against infrared flux for mergers. 115 3.3 By flux plotted against infrared flux for the mergers andSeyferts 118 3.4 IRAS colour-colour diagram for mergers 120 3.5The 1-0S(1) line flux plotted against aperture for Arp 220. 124 3.6a Surface brightness versus area weighted aperture for NGC 6240 S(l) line and [Fell] line . 125 3.6b Surface brightness versus area weighted aperture for Arp 220 S(l) line. * 126 3.6c Surface brightness versus area weighted aperture for NGC 1614 [Fell] line. 127
9 Chapter 4.
4.1 10pm map of Arp 299 from Gehrz, Sramekand Weedman (1983). 135 4.2a 19.6 arcsecond spectrum of IC 694 (source A). 136 4.2b 12.4 arcsecond spectrum of NGC 3690 (source B plus C). 137 4.2c 12.4 arcsecond spectrum of IC 694 (source A). 138 4.2d 7.8 arcsecond spectrum of NGC 3690 (source B alone). 139 4.3 By flux plotted against infrared flux for Arp 299 147 4.4 Surface brightness plotted against area weighted aperture for IC 694. 149 *♦ Chapter 5
5.1 Schematic diagram of an active nucleus 156 5.2 The 10pm emission superimppsed on CO contours of NGC 19068 158 5.3a 19.6 arcsecond K-window spectrum of NGC 1068 160 5.3b 12.4 arcsecond K-window spectrum of NGC 1068. 161 5.3c 7 .8 arcsecond K-window spectrum of NGC 1068. 162 5.3d 12.4 arcsecond H-window spectrum of NGC 1068. 163 5 .3e 7 .8 arcsecond H-window spectrum of NGC 1068. 164 5.4a S(l) line flux plotted against aperture. 167 5.4b By line flux plotted against aperture. 168 5.4c [Fell] line flux plotted against aperture. 169 5.5 [OIII] line map of NGC 1068 172 5.6a 19.6 K-window spectrum of NGC 4151 177 5.6b 7.8 arcsecond K-window spectrum of NGC 4151 178 5.6c 5.4 arcsecond K-window spectrum of NGC 4151 179 5.7a NGC 4151 line fluxes during die period January 1985 to January 1986. 182 5.7b NGC 4151 continuum flux densities during the period January 1985 to Janilfcry 1986. 183
Chapter 6
6.1a K-window spectrum of NGC 4631. 188 6. lb K-window spectrum of NGC 2798. 189
10 6.1c K-window spectrum of NGC 3079 190 6.2a By flux plotted against recession velocity for starburst galaxies. 19 3 6.2b 12p flux plotted against recession velocity for starburst galaxies. 194 6.2c 25pm flux plotted against recession velocity for starburst galaxies. 194 6.2d 60pm flux plotted against recession velocity for starburst galaxies. 195 6.2e 100pm flux plotted against recession velocity for starburst galaxies. 195 6.3 Random Byflux plotted againstrandom 100pm flux. 197 6.4a By luminosity plotted against 100pm luminosity, using observed fluxes. 197 6.4b By luminosity plotted against 100pm luminosity, using random fluxes. ^ 198 6.5 Warm and cold component luminosities plotted againstinfrared luminosity. 199 6.6a Random warm component luminosity plotted againstrandominfraredluminosity. 200 6.6b Random cool component luminosity plotted againstrandominfraredluminosity. 200 6.7a By flux plotted against recession velocity for the restricted redshiftrange sample. 205 6.7b 12pm flux plotted against recession velocity for the restricted redshift range sample. 206 6.7c 25pm flux plotted against recession velocity for the restricted redshift range sample. 206 6.7d 60pm flux plotted against recession velocity for the restricted redshiftrange sample. 207 6.7e 100pm flux plotted against recessionvelocity for the resctricted redshift range sample. 207 6.8 a By flux plotted against recession 12pm flux for the restricted redshift range sample. 208 6.8 b By flux plotted against recession 25 pm flux for the restricted redshiftrange sample. 208 6.8 c By flux plotted against recession 60pm flux for the restricted redshift range sample. 209
1 1 6.8 d By flux plotted against recession 100pm flux for the restricted redshift range sample. 209 6.9 Integrated mid-infrared emission plotted against By flux. 217 6.10 Groundbased 10pm flux plotted against By flux for starburst galaxies. 220 6.11 IRAS colour-colour diagram showing the positions of the starburst galaxies. 223
Chapter 7. Spectrum of NGC 1808, taken from Puxley et al. 1989. 233
Appendix. A 1.1 Initial mass function from Scalo (1986), with power law fits. 243 A1.2 Stellar mass plotted as a function of luminosity. 240 A 1.3 Main sequence lifetime as a function of mass. 24 1 A1.4 Ionising photon flux versus luminosity. 242
12 Tables
Chapter 2.
2.1 List of spectral standard stars. 64 2.2 Stellar absorption features compared to galaxy emission features. 68
Chapter 3
3.1 Line fluxes for merging galaxies 91 3.2 Starburstparameters » 107 3 • 3 F l ir as a function of upper and lower mass cut off for a ZAMS starburst. 109 3.4 Q as a function of upper and lower mass cut off for a ZAMS starburst. 110 3.5 F t m as a function of upper and lower mass cut off for an evolved starburst. Ill 3.6 Q as a function of upper and lower mass cut off for an evolved starburst. 112 3.7 Extinction corrections to By line flux for merging galaxies. 113 3.8 Extinction corrections to By line flux for Seyfert galaxies. 117 3.9 2-lS(l)/l-0S(l) and [FeII]/By ratios for ir Aging galaxies. 122 3.10 Exponential scale sizes for shocked lines and infrared continuum emission. 128 3.11 Differential extinction in Arp 220. 131
Chapter 4
4.1 Line fluxes for Arp 299 134 4.2 -1S(1)/1-0S(1) and By/1-0S(1) line ratios for m erging A rp 299. 141 4.3 Extinction in Arp 299 143 4.4 By/S( 1) ratio in IC 694 as a function of aperture 153
13 Chapter 5
5.1 line fluxes for NGC 1068 159 5.2 Lgy/LBOLfor stars of different spectral type. 171 5.3 Line fluxes for NGC 4151. 176
Chapter 6
6.1a Line and continuum fluxes for starburst galaxies 191 6.1b Line and continuum fluxes for Seyfert galaxies 192 6.2 Mean and standard deviation of randomly generatedfluxes. * 196 6.3 Spearman correlation analysis for restricted redshiftrange sample. 204 6.4 Partial correlation analysis foi;the whole sample of starburst galaxies. 205 6.5 Brackett line correlations for starburst galaxies. 210 6.6 S(l) line correlations for starburst galaxies. 211 6.7 Correlation of D 25 with IRAS fluxes for starburst galaxies. 212 6.8 Infrared compactness and size scales for starburst galaxies. 214 6.9 10 pm fluxes for starburst galaxies. 219 6.10 Brackett line correlations for Seyfert galaxies. 225 6.11 S(l) line correlations for active galaxies 225 6.12 Correlation of D 25 with IRAS fluxes for active galaxies. 226
14 Acknowledgements
First and foremost I would like to thank my supervisor, Bob Joseph. His good natured supervision, knowledge of astrophysics, and enthusiasm for this project have been a constant source of encouragement to me. I would also like to thank Gillian Wright for her encouragement at the start of my postgraduate work and her hospitality in Hawaii.
I would like to extend my thanks to everyone in room 1018 for many interesting conversations about astrophysics and life in general, and for helping to make my time at Imperial College so enjoyable. Special thanks to Rene Doyon for being a cheerful office mate, and fqr many stimulating discussions on spectroscopy. Thanks to Sunil Sidher for his help with computing problems and Jason Spyromilio for providing me with the [Fell] models. My undying gratitude is due to the small army of people who helped me produce this thesis. My partner Jonathan McDowell It read the zeroth version of this thesis and did the diagrams. Andrew Kispal typed some of the references. Natasha Sykes was chief proof reader, and was assisted by Amanda Baker and Dave Clements. I would like to thank Peter Hingley, Librarian of the Royal Astronomical Society, for unearthing a number of obscure references for me. Natasha Sykes and Troy Cooper provided much needed moral support, and Phil Hopkins deserves a medal for staying up 48 hours to finally print out this thesis. Andy Connolly bought me half a cup of coffee.
I would like to take this opportunity to thank the many people who encouraged my interest in Astrophysics at an early age, especially Heather Couper, Nigel Henbest and Bob Browning. I would not have had the courage to embark on a career in Astrophysics without the love and support of Jonathan McDowell. Thanks to Claire Craig for her perceptive friendship. My parents have put up with me being a student far longer than my contemporaries, and have always been generous with their financial support. Finally, I would like to dedicate this thesis to Jonathan McDowell and to my parents, with love and thanks.
I acknowledge a SERC studentship
15 Chapter 1 Astronomical Near Infrared Spectroscopy
1.1 Introduction
T is clear from the history of twentieth century astronomy that I progress inunderstanding astrophysicalphenomena correlates with the advent of new technology that opens up new parts of the electromagnetic spectrum for study. Infrared astronomy is a classic example of this, and has led to advances in all branches of astrophysics.
Astronomical infrared spectroscopy is plagued with difficulties, especially atmospheric absorptionfeatures and (untilrecently) detectors that are far less sensitive than detectors for optical astronomy. In spite of these problems, pioneering work in this field was carried out in the 1940’s by Gerard Kuiper and his collaborators who obtained spectra of the the sun, planets (Mercury, Venus, Jupiter and Saturn) and several bright stars including Betelgeuse, Mira C eti, R Leonis and Regulus (Kuiper et al. 1947, Kuiper 1947). Infrared spectroscopy of stars was developed in the 1960’s (Spinrad and Wing 1969), including detections of Paschen and Brackett lines from early-type stars (Kuiper 1963). Infrared spectroscopy was revolutionised in the 1970’s, with the devolpment of sensitive photo-conductors and photovoltaic detectors (Soifer and Pipher 1978). These new instruments led to advances in planetary astronomy (Larson 1980) and stellar astronomy (Merrill and Ridgway 1979), as well as to the first infrared spectra of interstellar molecules (e.g. Gautier et al. 1976).
The information derived from infrared spectroscopy differs from spectroscopy at shorter wavelengths in two ways. Firstly, molecular transitions with low excitation energies have wavelengths which fall in the infrared part of the electromagnetic spectrum. Secondly, extinction in the near infrared is typically only 10% of that in the optical, and so infrared spectroscopy can be used to investigate obscured sources. The objective of the work presented in this thesis is to explore the ways in which near infrared transitions can be used to understand physical
16 processes in galaxies. The aims of this introductory chapter are threefold; to describe the transitions most frequently seen in the H and K windows (see Figure 1.1), to outline the results of galactic spectroscopy, and to summarise the work to date on extragalactic infraredspectroscopy.
17 He I X2.059 X2.112 S(4) S(3) S(2) S(1 H, 2 -1 ' i 1 S(2) S (l) S(0) H2 1 -0 I I B 7 H 1 I ( 2 2.1 2.2
Wavelength (/mi)
Figure 1.1 Main transitions in the H and K windows
18 1.2 Hydrogen Recombination lines
The emission line spectrum of the hydrogen atom is most frequently observed to be photoionized in astronomical sources. These regions of photoionized gas are known as HII regions, and are often seen in the galaxy associated with hot stars, for example young OB stars or hot white dwarf stars. Photons with an energy E ^ 13.6eV have energies sufficiently high to ionize atomic hydrogen. Electrons and protons then recombine, producing a series of lines with wavelengths given by,
In-1 X=0.0912 pm ( 1.1)
where n’ is the principal quantum number of the upper state and n the principle quantum number of the lower state. Several of these transitions fall in the near infrared, most notably in the Brackett (n=4) and Paschen series (n= 5). In this section I will outline how the relative intensities of recombination lines are calculated, and how their absolute intensities may be related to the source of ionizing photons.
1.2.1 Derivation of Hydrogen Line Intensities
The most common model of photoionized regions assumes that the nebula is in ionization equilibrium, with the number of recombinations per second throughout the nebula equal to the number of Lyman continuum photons (NLyc) emitted per second from the source. The equilibrium state can be expressed as
NLyc» 4®-Unices (1.2) w hereotg is the recombination coefficient to all levels except the ground state, and n^otB gives the number of recombinations per unit volume per 4tt o unit time (Osterbrock 1989). The volume of the ionized gas is -y rs ■*, where rs is the Stromgren radius of the nebula and lie is the electron density.
19 The relative intensities of hydrogen recombination lines have been calculated in detailfor different conditions of temperature and density of astrophysical interest (e.g. see Osterbrock 1989). The goal of all such calculations is to determine the emission coefficients, jnn’> of a line resulting from a radiative transition from an upper level n to a lower level n \ The emission coefficient is given by;
n-1 - hvr *lv nn Jnn’ £ £ NnLA„’L’ (1.3) 4 tt L - 0 L ’ - L + l where hunn’ is the energy of a photon released in an electronic transition from an upper level n to a lower level n \ NnL is the population density of the state with principal quantum number n and orbital angular momentum quantum number L, and AnL^’L’ Is the radiative transition probability between an upper level state withprindpal quantumnumber n and orbital angular momentum quantum number L to a lower level state with principal qu antum number n’ and angular momentum quantum number L \ The emission coefficient is the energy emitted by the line in question per unit time per unit volume per unit solid angle throughout the nebula.. The emission coefficient is closely related to the effective eff recombinationcoefficient, a > given by
4 £ W = 2 eff a (1.4) hvnn.
The emission coefficient (or alternatively the effective recombination coefficient) of any state nL can be calculated from the radiative transition probability coeffidentif NnL is known. This is calculated assuming that the population density of a state is determined only by electron captures and downward radiative transitions, so that collisional exdtation is ignored. The equation of statistical equilibrium may be expressed (Osterbrock 1989)
n -1 (T) + Y N n>L’ An’L’,nL = NnL ^ ^ AnL,n”L’ (1.5) n n” - 1 L”
20 The first term represents the number of electron captures to the nL state, with otnL the recombination coefficient to the state nL. The second term represents radiative transitions into the nL state from higher states, and the final term on the right hand side is the number of radiative transitions from nL to lower states. This equilibrium equation assumes that the nebula is optically thin to all recombination lines (known as Case A recombination).
In most star formation regions the nebula is optically thick to the Lyman lines, in which case any transition n 2P - 12S will be immediately reabsorbed. Equation (1.5) is still used to calculate NnL> but the summation in the final term is from n”=2 rather than n”= 1. These conditions are referred to as Case B recombination, and are generally found to apply in star formation regions. Emission coefficients for infrared recombination lines, assuming Case B, are given by Giles (1977). The coefficients given by Giles are normalised relative to the HP flux, and absolute values of the emission coefficients for HP are given in Brocklehurst (1971). The jnn> are given as a function of density and temperature, butthey are relatively insensitive to both these parameters.
1.2.2 Derivation of Nebular Quantities
Measurement of the flux in a recombination line can be used to estimate the total number of ionizing photons from the source. The ionization equilibrium condition requires that the number of ionizing photons from the source is equal to the number of recombinations throughout the nebula. The number of ionizing photons is given by
( 1.6)
where photons with a frequency >uo have sufficient energy to ionize hydrogen, and LD is the energy of a photon of frequency u. The number
21 L(nn’) _ i 4irjnn’dV = n | Onn’ dV (1.7) hvnn* hvnn>
Therefore the ratio of the number of photons in the recombination line to that in the ionizing continuum is
( 1 .8)
Therefore the number of ionizing photons can easily be calculated once the recombination line flux has been measured. The luminosity in the ionizing continuum can be-calculated by assuming the shape of the continuum .
The derivation of the number of ionizing photons given above does not allow for any decrease in the recombination line flux due to extinction. The extinction can be estimated if the flux of more than one recombination line is measured, assuming that the intrinsic ratio of the recombination lines is given by a model predicting the recombination line intensities (usually Case B in star formation regions), and that the differential extinction (or reddening) between the two lines is known. The ext?«action Ax of any re a r xDination line of wavelength X is given by;
(1.9)
where Mx is the observed magnitude of the recombination line corresponding to a flux density lx- The subscript ’0’ denotes the dereddened values. The difference in extinction between any two recombination lines Xi and X 2 is given by
22 ( 1.10) *2=2-51o s (r 9
in the two recombination lines, and Ro is the intrinsic ratio which is known from case B recombination. Adoption of a reddening curve then gives the extinction at X i and X 2. For example, assuming case B recombination, the ratio of the flux in the By line to that in the B a line is 0.353, and the differential extinction between the two lines is Agy =2.7 5 Afia (using the infrared reddening curve of Rieke and Lebofsky 1985). Therefore the extinction correction to the By line becomes
(1. 11)
23 1.3 Rotation-Vibration Spectrum of Molecular Hydrogen
Hydrogen is the most abundant element in the universe, and molecular hydrogen is the most stable form of hydrogen at low temperatures. Molecular hydrogen is known to exist in the atmospheres of cool stars and planets (Field, Somerville and Dressier 1966) and in the interstellar medium (Shull and Beckwith 1982). Most of the molecular hydrogen in the galaxy lies in molecular cloud complexes that are known to be associated with star formation. In spite of its abundance, the H 2 molecule is very difficult to observe. This is primarily because the H 2 'molecule is homonuclear and so has no dipole moment Hence the the rotation-vibration spectrum has amuch lower transition probability than other diatomic molecules such as CO. Reviews of the physics of diatomic molecules can be found in most elementary spectroscopy texts, e.g. Barrow (1962). Hertz^erg (1950) describes the physics of homonuclear molecules, and the structure of the hydrogen molecule is described in Field, Somerville and Dressier (1966). In this section I will outline the rotation-vibration spectrum of molecular hydrogen, details of which can be found in the texts cited above.
A schematic energy level diagram for the electronic ground state of a diatomic molecule, showing intemuclear distance versus, potential energy , is shown in Figure 1.2. This is known as the Morse curve. It is approximately parabolic for small departures from the equilibrium separation, but atlarger separations the potential energy curve flattens off until the dissociation energy is reached. Each electronic state supports a series of vibrational energy levels (labeled u). There is also a series of rotational levels (denoted J) associated with each vibrational state, but these are not shown because the rotational levels have much smaller energies.
The total energy of&given vibrational-rotational level can be expressed as
£v,j = £v +EJ (1.12)
24 w heree 0 is the vibrational energy of the vibrational state u, and ej rotational energy of the J ^ rotational state. This expression assumes that the vibrational energy of a molecule is independent of the rotational energy (the Bom-Oppenheimer approximation). The rotational energy of a molecule can be calculated by solving the Schrodinger equation for a molecule floating freely in space. It is given by
Figure 1.2 Schematic diagram showing potential energy versus.intemuclear distance for a diatomic molecule, from Barrow (1962).
e] = BJ (J + l) (1.13) w here B = —h— (1.14) 87T 2Ic and I is the moment of inertia of the molecule. This expression ignores centrifugal effects, i.e. it assumes that the rotation has no effect on the
25 intemuclear distance, and thus that I is independent of the rotational state. An expression for the vibrational energy of a molecule can be written down as
(1.15) where we is the equilibrium oscillation frequency. Ib is expression is valid only when the displacement from the equilibrium position is small; that is, in the region of the potential energy curve where the Morse curve is well approximated by a parabola. A term can be added to equation (1.15) to take anharmonic effects into account, but this term is small, and so will be neglected here. Equations (1.13) and (1.15) can be substitutedinto equation ( 1. 12) to the totalrotational-vibrational energy of the molecule;
ev,j = BJ (J + 1) + (v + 1/2) coe (1.16)
The rotation-vibration spectrum of the hydrogen molecule is produced by an oscillating quadrupole moment. These quadrupole lines have transition probabilities that are ~ 10"8-10"9 times smaller than transition probabilities for dipole radiation. The selection rules for changes in the rotational states are J’-J=±2,0, where J’ is the quantum number of the upper rotational state and J the quantum number of the lower rotational state. There are no restrictions on changes in v. The energy associated with a transition from a anupper rotational-vibration state w ith v = v" and J = J" to a lower state v = u\ J = J’ can be easily calculated from equation (4). For example, a transition from the u = 1 state to the vibrational ground state, with AJ = 2 is given by;
AE = we + Bj4J + 6) (1.17)
Arotation-vibrationtransitionis conventionally labelled with the upper andlower vibrational states, the change inrotational quantum number and the lower rotational state. The AJ=0 transitions are referred to as the Q-branch, the AJ=+2 as the S branch and the AJ=-2 as the O
26 branch. For example the u = 1 -0S( 1) line is a transition from the v = 1, J=3rotational-vibrationalstatetothe v = 0, J = 1 state. This is often abbreviated to 1 -0S( 1).
The relative intensities of the different transitions in thermal equilibrium can be calculated by assuming that the intensity of a given component is proportional to the population of the rotational energy from which the component originates. Since the multiplicity of the Jth rotational level is 2J + 1, the population of the Jth rotational level in Boltzmann equilibriumis
Nj = (2J + 1) N0exp(-BJ(J+i)/kT) (1.18)
where No is the population of the^rotational ground state. The first rotational level is 6000 K above the rotational ground state for the hydrogen molecule. This high excitation energy adds to the difficulty of detecting the H 2 rotation -vibration spectrum from quiescent molecular clouds w ith a tem perature of « 100 K.
There are two models which describe the excitation of H 2 in astrophysical conditions. One is ultraviolet fluorescence and heating and the other thermal excitation behind a shock front. In the remainder of this section these two excitation mechanisms will be reviewed, and methods for distinguishing between these processes discussed.
1.3.1 Ultraviolet Fluorescence and Heating
Ultraviolet radiation can excite the rotation vibration spectrum in two ways; it can fluorescently pump the gas which decays via radiative transitions (pure fluorescence, Black and Dalgamo 1976, Black and van Dishoeck 1987) and it can heatlhe gas so that spectrum is thermally excited (UV heating, Sternberg and Dalgamo 1989). Therelative importance of these two mechanisms is dependent on the density of the gas. Both of these mechanisms will be described here.
27 Pure UV fluorescence can occur in a gas with number density ^104 cm*3. The basic physical mechanism is as follows. A photon with A>912 A has sufficient energy to excite the molecule via an electronic transition to the Lyman and W erner bands. This is then followed by a radiative decay to the higher rotation-vibrationlevels of the electronic ground state. From here it will cool via a cascade through the infrared vibration-rotation lines. Models of ultraviolet fluorescence predicting the relative strengths of the infrared lines for different values of the gas density and the ultraviolet flux have been computed by Black and Dalgamo (1976) and Black and van Dishoeck (1987). The models of Black and Dalgamo are only valid if the UV radiation field is weak; i.e. once'a molecule has absorbed a photon it has time to relax to the ground state before it absorbs another. Detailed models have been presented by Black and Dishoeck, which include the strong field case. The absolute intensity of the infrared lines in the case of pure fluorescence depends on the intensity of the UV radiation £Leld and on the H 2 density. The relative intensity of the infrared lines is highly non- thermal, and is independentof these parameters.
At higher densities (^ 104 cm-3), H 2 - H2 and H - H 2 inelastic collisions have important effects on the spectrum. The effect of collisions has been investigated in detail by Sternberg and Dalgamo (1989). The main result of this study is that collisions can de-excite the UV pumped H 2, resulting in the conversion of internal molecular energy into kinetic energy which heats the gas. The relative intensities of the collisionally de-exdtedinfrared lines depend on the transition branching ratios for a collisional cascade which in turn depend upon the H - H 2 and H2 - H 2 collision cross sections. These cross sections are not well determined so the spectrum is somewhat uncertain. However, since the downward transitions depend on collisional rather than radiative de excitation, the spectrum is relatively insensitive to the intensity of the UV radiation field. This mechanism is known as ’collisional fluorescence’.
Collisional de-excitation dominates the downward decays at gas densities^ 10 5 cm-3 However the relative intensities of the infrared lines become sensitive to the intensity of the UV radiation field because athigh temperatures collisional excitation dominates the upward transitions of the molecules. In the weak field limit the temperature is
28 the S(l) then UV heating must be questioned. It is not possible to definitively rule out UV heating line by this method, however, since it is possible to imagine geometrical situations in which the filling factor of the molecular gas is greater than that of the ionized gas.
Figure 1.3 Comparison of fluorescent and thermal emission, taken from Black and Van Dishoeck (1987). The dotted line shows thermal emission from a 2000 K gas, and the solid line fluorescent emission.
31 1.4 Forbidden Iron Transitions
The theory of Quantum Mechanics predicts that certain electronic transitions are ’forbidden i.e. have a very low transition probability. In laboratory conditions these transitions are rarely, if ever, seen. If an atom is excited into a state in which there is no high probability decay route (a metastable state), it will be collisionally de-exdted before the forbidden transition has a chance to occur. In astrophysical conditions the gas density is often so low (< 10^ cm " 3 ) that collision timescales are longer than the timescales involved for the forbidden transitions, and forbidden transitions are frequently observed, especially around active galactic nuclei (Osterbrock 1989).
One of the least frequently observed forbidden transitions is [Fell]. Here the AS=0 selection^ rule is violated. [Fell] can be photoionized, as seen in Orion nebula (Low, Moorhead and Wehlau 1979). The metastable states of [Fell] can be collisionally populated, either thermally or behind a shock front. The best models for shock excitation are those of McKee, Chemoff and Hollenbach (1984) outlined in the previous section. Figure 1.4 shows the variation in the relative intensity of the [Fell] X1.664 and the [Fell] X1.644 pm lines as a function of gas density and temperature(Spyromilio, private communication). From the Figure it is clear that while the ratio of these lines is insensitive to temperature, it is a good density indicator in the range 10^ -10^ c m " 3
1.5 Observations of Molecular Hydrogen Emission from Galactic Sources
The rotation-vibration spectrum of molecular hydrogen has been observed to be excited by ultraviolet fluorescence and shocks in galactic sources. Shock excited emission has been observed when an outflow source impacts onto a surrounding molecular cloud, or outflow sources that have molecular components. Examples of shock excited emission include the interaction of supernova remnants with the interstellar medium (e.g. IC 433,Treffers 1979, Graham, Wright and Longmore 1987), the expansion of planetary nebulae,
32 [Fe II] line ratios Fl]ln aisa fnto ftmperature. tem of function a as ratios line [Fell] eprtr (K)Temperature iue 1.4a Figure 33 [Fe II] line ratios molecular clouds. Fluorescent excitation has been observed a in observed been has excitation Fluorescent clouds. molecular reflection nebula by Gatley etal (1987), and in a planetary nebula and by in a planetary nebula (1987), Gatley etal by nebula reflection and Fink 1981) and in the interaction of stellar winds/outflows with stellar winds/outflows of interaction in the and 1981) Fink and Dinerstein et al.(1988). In this section I will describe the H2 emission H2 the describe I will section In this al.(1988). et Dinerstein (Beckwith, Persson and Gatley 1978, Isaacman 1984, Smith, Larson Smith, 1984, Isaacman 1978, Gatley and Persson (Beckwith, [Fell] line ratios as a function of electron density. electron of a as function ratios line [Fell] lcrn est (m 3) (cm density Electron iue 1.4b Figure 34
from star formation regions and supernova remnants. I will concentrate on the Orion Molecular Cloud and IC443 as these are two of the best studied H 2 emitting regions in the galaxy, and have been extensively studied at other wavelengths.
1.5.1 Molecular Hydrogen in the Orion Star Formation Region.
The region around the Trapezium cluster in Orion is the best studied star formation region in the Galaxy, so it will considered here in some detail. Molecular hydrogen emission was first observed in Orion by Gautier et al. (1976) who measured the ratios of the S and Q branches, and concluded that the ratios were consistent with thermal excitation. Grasdalen and Joyce (1976) then mapped the S( 1) line and discovered it to extend over 1 arcminute. Higher angular measurements by Beckwith, Persson and Neugabauer (1978) showed that the emission is clumped on scales of 5 arcseconds.
There are three plausible excitation mechanisms, discussed by Kwan (1977): thermal excitation, thermal excitation behind a shock front and ultraviolet fluorescence. The first is unlikely as the excitation temperature of * 2000 K is much higher than that of the surrounding molecular cloud (» 40 K). Ultraviolet fluorescence is inconsistent with the low 2- 1S(1)/1-0S( 1) line ratio, and is also unlikely on energetic grounds. The observed flux of » 1-100 L o in the S( 1) line requires an ultraviolet flux of 1 0^-1 O^Lo (See introductory section.). The most plausible source for the ultraviolet flux is the O-star a-Ori, which has a UV luminosity of ® 1 04Lq . Thus nearly all the ultraviolet flux from a- Ori must go into the S(l) line. This is physically unreasonable, especially as the projected distance of the H 2 region from a-Ori implies .>• that geometrical diluting of the ultraviolet flux would make it necessary fo r the H2 region project a large solid angle to the source. The emission is clumped, making this unlikely. The shock model is broadly consistent with the main features of the emission, v iz the high excitation temperature, the large angular extent, and the line ratios. Itis also consistent with evidence that the Orion Molecular Cloud contains regions of highly supersonic flow, for example a bipolar outflow seen in the J= 1-0 CO line, Herbig - Haro objects and masers.
35 There is also clear evidence thatUV- pumped H 2 is seen in the Orion Region. Two arc minutes southeast of the Trapezium cluster is a bright ionization front (known as the Orion Bar), thought to be produced by the interaction of the Orion HII region with a molecular cloud. The expansion of the HII region has produced a region of shock-compressed gas, as seen in CO (Schleorb and Loren 1982) and CS (Omoduka et al. 1984). H ayashietal( 1985) have measured the intensities of the 1-0 S(l) line and the 2- 1S(1) line across the ionization front. Away from the front the relative intensities of the 1 -0S( 1) line and the 2-1 S( 1) line are consistent with low density UV fluorescence. Across the ionization froftt, coincident with the CS gas, the ratio l-0S(l)/2-lS(2) ® 3, indicating a shocked component to the emission. In the light of the results presented by Sternberg and Dalgamo (1989; described in section 1.3.1), it is possible that the thermal emission at the bar is due to high density UV heating. However tjie spatial coincidence of the thermal H 2, compressed CS and high velocity CO supports the original interpretation of Hayashi et al. Thus, there is evidence for both shocked and fluorescently excited gas in Orion. The shocked gas is associated with other indicators of supersonic outflow, such as high velocity CO Herbig - Haro objects and masers.
The Orion region is the best studied star formation region, but H2 has been found in many other star formation regions; e.g.. W75N/DR21, OMC 2 (Fischer et al. 1980), S106,(Longmore, Robson and Jameson 1986), Cepheus A (Doyon and Nadeau 1988). All of these regions display some evidence for outflow. A more detailed examination of the energetics of the flows by Fischer et al. (1985) shows that not only do the CO flows have sufficient energy to power the H 2 regions, but that the efficiency is typically between 10 and 90%. Thus H2 emission is the main cooling agent for the CO outflows.
1.5.2 Molecular Hydrogen Emission from IC 443
IC 443 is a supernova remnant located in a molecular cloud complex which includes an HII region (S249) and an OB association, and is thus clearly associated with a region of active star formation. It is visible on sky survey plates as an arc of filamentary emission, approximately a
36 degree across. The molecular hydrogen emission from IC 443 has been studied by Treffers (1979), Graham, Wright and Longmore (1987) and Burton et al. (1988). The most complete map of the S( 1) line emission was obtained by Burton et al. and will be described here.
The main features of the H2 emission from IC 443 are described below. It is extended M x 10 pc along an S-shaped arc, and is spatially coincident with high velocity CO and HI. The emission is clumpy, and the peaks in the H2 emission lie at the peaks of the high velocity CO. The relative intensities of the H 2 lines in the K-window are consistent with thermal excitation. The By line was not detected in any of the apertures where the H2 emission was detected. This, coupled with the spatial coincidence of the H 2 emission with the high velocity gas, argues for thermal excitation behind a shock front rather than UV heating.
The observations described above indicate that the H 2 emission is produced by the interaction of the remnant with a molecular cloud. In this model the high velocity CO originates in the swept-up ambient material, the HI emission from molecular hydrogen that is dissociated in the shock. The mechanical energy per unit area in the shock (estimated from the X ray surface brightness) is comparable to the luminosity in the H 2 lines, indicating that the H 2 emission is an important coolant. The total H 2 luminosity, allowing for extinction is »1000 L o , making it one of the most luminous sources in the galaxy.
1.6 Infrared Line and Continuum Emission from HII Regions
The most common source of ionizing radiation in the Galaxy is young massive OB stars. Ionizing radiation from the newly formed stars produces an HII region, and eventually most of the energy from the star is absorbed by dust grains in and around the HII region. The dust grains radiate in thermal equilibrium at a range of temperatures between 50 and 200 K, and so HII regions are conspicuous in the infrared. In this section I will describe the infrared and radio continuum emission from HII regions, and then summarise the contribution that infrared spectroscopy has made to our understanding of star formation.
37 1.6.1 Radio Continuum Emission
An HII region is formed when an O or B star (luminosity > 10^ Lo ionizes an area of hydrogen gas around it Figure 1.6 shows the continuum energy distribution of a typical HII region from the submillimetre to 1mm, taken from Fazio (1978). The spectrum is flat at radio wavelengths, characteristic of optically thin free- free (or Bremsstrahlung) radiation. Bremsstrahlung radiation is emitted when free electrons are accelerated in the electric field of the positive ions. Measurement of the free-free radio flux can be used to determine the number of ionizing photons in the same way as the flux in a recombination line. This can be demonstrated by noting thatthe number of ionizing photons given in equation ( 1.6) is proportional to f lie 2 dV.
This quantity (also known as the emission measure, E) is proportional to the radio flux density (S0) , since for an optically thin plasma (Bowers andDeeming 1984):
‘tatfSv = Iv s TV (1.19) c2 w here I0 is the intensity at radio frequency u, T is the electron temperature, and r 0 is the optical depth at a frequency u given by (Osterbrock 1989);
tv = 8.24 x 10-2T 35 v 21 E (1.20)
Thus the free-free radio flux density is proportional to E, and the free- free emission coefficient at radio wavelengths can be used to determine NLyc using equation 1.8. The radio emission coefficients have been tablulated by Osterbrock (1989) and Joy and Lester (1987). *1
38 Figure 1.5 Continuum Energy distribution of an HU region
The ionizing photon flux is a sensitive indicator of the spectral type of the star that ionizes the nebula, since early type stars emit more ionizing photons than late type stars. This is illustrated in Figure 1. 6, which shows the ionizing photon flux as a function of surface temperature for O andB stars, takenfromPanagia (1973). Therefore, measuring NLyc provides a method for estimating the spectral type of the star that ionizes the nebula.
1.6.2 Infrared Continuum Emission from HII regions
At shorter wavelengths, the spectrum rises sharply, with intensity in excess of that expected from the ionized gas. This excess infrared flux is characteristic of thermal emission, and the obvious source of thermal
39 Log (Ionising photon flux) (photons s in frared lum inosity inosity lum frared in prim ary heating source w as L a photons, in w hich case the integrated integrated the case hich w the t in a th photons, a suggested L as w odels m source first e h J heating ary 1974). prim Becklin d an illiams W htn s( nzltl 1982); zeletal. en (G as photons em ission is heated dust grains in an d around the H II region (W ynn- ynn- (W region II H the around d an in grains dust heated is ission em 50 48 46 44 i l 2x10 4 Ionizing photon flu x vs. stellar spectral type spectral stellar vs. x flu photon Ionizing 2 1 O5 B 0 0 0 0 - 05 06 07 09 BO ‘ BO.5 B1 B2 ______L S ( t o t ) ) can be estim ated from th e num ber o f ionizing f ionizing o ber num e th from ated estim be can iue 1.6 Figure T.« (ZAMS) 3x10 4 4x10 4 4x10 4 3x10 40 I I i I l I I ______L 5x10 4 5x10 _ FhvLyKNLyc / p i 2 Wm 2 (1.21) ^ 4jr(3xlO‘»)2 W*) where F is the fraction of La photons absorbed by the dust, huLya is the energy of a La photon, and D the distance of the object in kpc. The infrared flux density at any wavelength X can also be estimated as a function of the free-free radio flux density at a frequency u by assuming that the dust grains radiate as a grey body at a temperature Tg> (Genzel et al. 1982)
SirU)_2.28x 1Q16f ( 1.22) S v x 3 T |
where Te is the electron temperature and X is measured in microns. This relationship has been used to test the validity of the L a heating model, and some HII regions have infrared fluxes consistent with L a heating (e.g. Genzel 1982, Thronson, Harvey and Campbell 1978).
The integrated infrared flux of a number of HII regions are not consistent with La heating. Figure 1.7 shows a plot of ionizing photon flux against infrared luminosity for a large sample of HII regions, taken from Panagia (1978). The data were obtained from Emerson et. al. (1973) and Furness et al. (1975). It is clear while there is a good correlation of ionizing photon flux with infrared luminosity, the observed infrared luminosity is an order of magnitude larger than that predicted on the basis of the ionizing photon flux. This effect is parameterised in terms of the ’infrared excess’,
41 Figure 1.7 Ionizing photon flux vs. stellar luminosity for galactic HE regions.
which is the ratio of the infrared luminosity to the L a luminosity. Most HE regions have infrared excesses that are in the region of 10.
There are two reasons for this discrepancy; the dust is absorbing photons with wavelengths longer than the Lyman limit and/or it is competing with the gas for Lyman continuum photons. The infrared luminosity can be expressed as LlR = Ll/x + (1 - f ) (hv Lyt^Lyc) + Lv
42 from the spectral type of the star, determined from the ionizing photon flux. For most HII regions f*0.3-0.5, and the optical depth in the Lyman continuum is «1. Under these conditions essentially all the energy emitted by the young stars is absorbed by the dust and the infrared luminosity is a good estimate of the bolometric luminosity. Since the infrared luminosity is proportional to NLyc, it can be used to determine the spectral type of the stars that ionize th e nebula.
1.6.3 Infrared Recombination Lines from HII Regions
Observations of infrared recombination lines in HII regions have shown that many HII regions have recombination line fluxes that are equal to that predicted on the basis of the ionizing photonflux determined from the radio continuum emission, assuming Case B recombination and allowing for extinction. There js, however, a class of objects where the ionizing flux estimated from the recombination line flux and case B recombination is larg er than that which would be predicted on the basis of the spectral type of the star (as determined from the infrared luminosity), and larger than that estimated from the free-free radio flux, (Thompson 1982, Simon etal. 1983). Note that extinction of the infrared lines would tend to make the recombination line intensities smaller than those predicted by the radio continuum flux. These HII regions are known as ’line excess’ objects.
This is best demonstrated in Figure 1.8 from Thompson (1984). Here ne^V is plotted against luminosity for a selection of line excess objects. n ^V is a measure of the recombination rate inside the nebula and is thus a measure of the ionizing luminosity. It is derived from a measure of the By line flux. The solid line in the diagram shows the luminosity in the Lyman continuum expected on the basis of stellar models (taken from Panagia 1973) for ZAMS stars. All the objects plotted have ne^ V values larger than ZAMS values.
There are three popular solutions to the line excess problem. In the first the excess flux is interpreted as luminosity from accretion onto the young stellar object, probably from a disk (Thompson 1982). Secondly, an outflowing wind has been proposed, rather than the static
43 nebula of Case B. Here the lines arise in the extended atmosphere, and the excess is due to the geometrically larger emission region (Simon et al. 1981). The most popular idea is that there is significant population of the n=2 level, leading to ionization from the Balmer continuum. (Simon et al. 1983, Thompson 1984). The dashed line in Figure 1.8 shows the number of ionizing photons available as a function of luminosity if Balmer continuum photons are allowed. It is obvious that most sources havene^V values smaller than the Balmer ZAMS line. It is interesting that low luminosity objects (< 10^ ) also show line excess (Thompson 1987). If the n=2 level population is due to Lyman trapping this is unexpected since the Lyman flux from these objects is small. It remains to determined whether low luminosity objects have the same line excess mechanism as the high luminosity ones.
44 Log (L/Lo) Em ission m easure vs. lum inosity fo r ordinary galactic H E regions and regions E H galactic ordinary r fo inosity lum vs. easure m ission Em line excess objects. excess line iue 1.8 figure 45 1.7 Observations of [Fell] from Galactic Sources.
Transitions due to [Fell] have been observed from a variety of galactic sources, for example Herbig-Haro objects (Canto 1985) and the Trapezium cluster in Orion (Lowe, Moorehead and Wehlau 1979). The near infrared lines at 1.6pm, however, have been observed in only a handful of sources, either supemovae or supernova remnants: a supernova in M83 (Graham et al. 1986) and the supernova remnant MSH 15-52(Sewartetal. 1983). They have also been detected from s- Carina (Allen, Jones and Hyland 1985).
“ The supernova remnant IC 433 is the only source that has been observed in detail, by Graham, Wright and Longmore (1987) who have detected several lines in the H and K windows which they attribute to [Fell] at 1.644, 2.121, 2.165 pm. The most distinctive feature of the [Fell] 1 644 emission is that 1— ^ ----- * *$0-70 compared to the value of 0.06 predicted on the basis of case B recombination in the Orion region. This very high ratio can be explained by a combination of two processes: that the emission in IC433 is shock excited rather than photoionized, and that iron is less depleted in the supernova remnant relative to the Orion region since dust grains can form in the post-shock cooling region. The line ratios observed in IC 433 are consistent with the shock model of McKee, Chemoff and Hollenbach (1984), which would predict that the iron depletion was » 0.2-0.6, compared to a value of * 0.04 in the Orion region. Thus not all the enhancement of [Fell] relative to By can be explained by a reduced Fell depletion in IC 433.
There is other evidence that the [Fell] is collisionally excited. Ratios of the optical to infrared lines indicate an excitation temperature of 5100k. The scatter is only 1600K, consistent with the emission being from gas with a narrow range of excitation temperatures. Also, none of the observed optical [Fell] trqjnsitions feed the^ Instate as would be expected from photoionized transitions. Thus a high [FellJ/By ratio can be considered to be characteristic of shock excited [Fell] (Graham, Wright and Longmore 1987)
46 1.8 Extragalactic Near Infrared Spectroscopy
The first near infrared spectrum of an extragalactic object was obtained a decade ago by Thompson, Lebofsky and Rieke (1978). Their K- window spectrum of the nucleus of NGC 1068 showed clear detections of the 1 -0(S 1) line, the By line and absorption bands of CO. Pioneering work in this area was also carried out by Fischer etal. (1983) who obtained detections of the By and 1-0S(1) lines from the interacting system NGC 3690/IC 694 in two apertures. This paper showed that the characteristics of the infrared lines in this interacting galaxy were similar to those in galactic star formation regions, and demonstrated the potential of infrared spectroscopy for investigating star formation in other galaxies. Early detectionsof Brackett line fluxes also include NGC 1614 (Aitken, Roche and Phillips 1981) NGC 1808, and He 2-10 (Phillips, Aitken and Roche (1984). Infrared spectroscopy became established with the detection by Joseph, Wright and Wade (1984). (see also DePoy, Becklin and Wynn- Williams 1986) of luminous molecular hydrogen emission in the merging galaxies NGC 6240 and Arp 220. Since then there have been several in-depth studies of individual galaxies (e.g. Joy and Lester 1988, Lester Harvey and Carr 1987, Rieke' etal. 1985) and a few surveys of many galaxies (Kawara, Nishida and Gregory 1987, Moorwood and Oliva 1988).
Infrared spectroscopy of astronomical sources is plagued with difficulties not encountered the radio and optical parts of the electromagnetic spectrum, especially the removal of atmospheric absorption features and problems with sky noise. Extragalactic infrared spectroscopy has the additional problem that the line strength is typically only a few percent of the continuum, requiring long integration times to obtain a good detection. These difficulties are reflected in the number of papers written almost wholly on the basis of upper limits to line fluxes. For example Heckman et al. (1986) obtained upper limits to the 1 -0S( 1) line flux in 11 galaxies, but detected this line only in NGC 7469 and Tuner, Ho and Beck (1987) searched for the Ba and Byrecombination lines in five positions in M83, and detected the By line from only one position and the Ba from only two.
47 The long integration times required to obtain good detections may explain the poor quality of some of the data in the literature. The observations of the By line in NGC 1097, NGC 2903, and NGC 4151 reported by Beck, Beckwith and Gatley (1984), for example, are only 2 and 3a detections. These formal errors do not take into account the uncertainty involved in positioning the continuum level in these spectra. Since only one or two points on either side of the line define the continuum, this uncertainty is considerable. This problem - poor spectral coverage leading to a badly defined continuum level - occurs in many papers, for example Kawara Nishida and Gregory (1987), Turner, Ho and Beck (1987) and to a lesser extent in spectra of Seyfert galaxies presented by Fischer et al. (1987). It is sometimes exacerbated by only sampling the spectrum onee per resolution element. This can lead to the situation where detections are claimed for one point above the continuum level, as in the detection of the 1-0S(1) line in NGC 4038 in Kawara, Nishida and Gregory (1987). *
Approximately 50 detections of the By line and 40 detections of the 1-0S(1) line have been reported in the literature. It is remarkable that, given the difficulties outlined above, most measurements of infrared lines have been shown to be repeatable. Allowing for differences in aperture size, measurements of the 1-0S(1) line in NGC 6240 from Joseph Wright and Wade (1984), Lester, Harvey and Carr (1987) and data presented in this thesis (discussed in Chapter 3) are consistent. There is good consistency between Thompson, Lebofsky andRieke (1978) and Hall et al. (1981) in their measurements of the fluxes of infrared lines in NGC 1068, and the various measurements of By and 1 -0S(1) line fluxes in Arp 299 (described in Chapter 4) are in good agreement There are some exceptions to this, the most notable being the controversy over the By line flux in NGC 6240 (Rieke et al. 1985, DePoy, Becklin and Wynn-Williams 1986 and Lester, Harvey and Carr 1987). In conclusion, near infrared lines have been detected from about 50 galaxies, and although some of the data are of dubious quality, most measurements are consistent when allowances are made for differences in aperture size. The remainder of this section describes the astrophysical insights gained from infrared spectroscopy.
48 1.8.1 Extragalactic Infrared Recombination Lines
One of the essential parameters for any modelling of star formation regions is the luminosity in ionizing photons. In galactic HII regions it is possible to map the thermal radio flux in order to measure NLyc- In external galaxies, however, a larger area is covered by the radio beam and the radio flux is often dominated by non-thermal sources. Therefore measuring the flux in a recombination line is the only way of directly estimating NLyc hi external galaxies, except for nearby galaxies where spectral index mapping is possible, e.g. NGC 253 (Turner and Ho 1983). Infrared recombination lines are less sensitive to extinction than optical recombination lines and therefore provide a good method of estimating NLyc hi dusty galaxies.
The objective of all observations of extragalactic infrared recombination lines has been tq determine what fraction of the bolometric luminosity can be attributed to current star formation, and how much to other energy sources, such as an active nucleus. The basic approach to the problem is as follows. It is assumed that infrared recombination lines arise in HII regions ionized by young stars. The number of ionizing photons, NLyc> from the star formation regions within the aperture is then calculated from the recombination line flux as shown in section 1.2.2. The luminosity expected from the young stars on the basis of NLyc can be calculated by adopting an initial mass function for the burst and assuming a starburst model. This can then be compared to the infrared luminosity derived from ground based or IRAS measurements.
This approach has been taken by Beck, Beckwith and Gatley (1984), Beck, Turner and Ho (1986), and Ho, Beck and Turner (1989) who obtained measurements of the Ba and By fluxes for a number of ordinary, interacting and a couple of Seyfertgalaxies. They calculated the bolometric luminosity expected from stars from under the assumptions that the upper mass cut off is » 30Mq , the lower mass cut off » 5 Mo, and that the burst is unevolved. They found that the luminosity derived from the ionizing flux was within a factor of 5 for the galaxies they observed, with the exception of the Seyfert galaxies NGC
49 4151 and NGC 1068. The uncertainties in this calculation will be discussed below.
There are two major uncertainties involved when deriving the stellar luminosity from the infrared luminosity; the upper and lower mass cut off adopted for the burst, and the age assumed for the burst. The following examples illustrate the limitations of these assumptions. Increasing the upper mass cut off to 60 M o for an unevolved starburst will decrease the stellar luminosity deduced from the number of ionizing photons by a factor of two. The assumption of a high (5 Mo) lower mass cut off becomes questionable if the star formation has proceeded at a constant rate fo r^l 07years - under these circumstances low mass stars can contribute significantly (~ 80%) to the deduced stellar luminosity. The intrinsic uncertainty in the derivation of the stellar luminosity from the ionizing flux using the assumptions outlined above is a factor of 5. * The most extensive survey of the Brackett lines in galaxies has been carried out by Depoy (1987). He measured the Ba and By fluxes in a sample of luminous IRAS-selected galaxies, and found that the ratio of the number of ionizing photons to the bolometric luminosity in most of the galaxies is similar to the empirical value observed in HII regions, with the exception of Arp 220 and NGC 6240. He made no attempt, however, to model the observed recombination line and infrared continuum fluxes in terms of a starburst
As a self-consistency check on starburst models it is useful to compare the number of ionizing photons derived from infrared recombination lines to that inferred from the thermal radio flux. This can only be done in a couple of nearby galaxies where radio spectral index mapping can be obtained with sufficient resolution to deconvolve the thermal from the nonthermal radio flux. In the nearby starburst galaxy NGC 253 it was found that the ionizing photon flux derived from the radio free-free emission is consistent with that derived from the Brackett lines (Beck and Beckwith 1984). The nuclear regions of M83 is somewhat different in character; here the free-free radio flux predicted on the basis of the recombination line fluxes is larger (by a factor 6) th an that inferred from radio spectral index maps of the same area (Turner, Ho and Beck 1987). There are two plausible explanations for this
50 phenomenon; that the star formation is dominated by compact HII regions which are optically thick at 5 GHz (thus suppressing the thermal radio flux) or that there is a large population of the ’Line Excess’ objects described in section 1.2.2. Either of these propositions implies that the starburst is very young. Compact HII regions have a lifetime «10 3- 104 years, and the starburst models of Thompson (1987) require the burst to be ~ 106 years old.
One paper which demonstrates the value of Brackett line spectroscopy is the study of II Zw40 , a dwarf galaxy with a strong HII region spectrum, by Joy and Lester (1988). They obtained a measurement of the By flux, and, in conjunction with a previous observation of the H P flux r derived a value for the extinction at K and thus an extinction-corrected By flux. They used the extinction-corrected By flux to determine the contribution from the free-free emission to the K continuum and the radio fluxv They concluded that the radio flux was dominated by thermal emission from HII regions, rather than non- thermal emission from supemovae, and that the free-free emission contributed about half of the continuum flux at K. Multiaperture optical spectroscopy indicates that young stars contribute 30% of the light at K. Joy and Lester concluded that II Zw40 is a canonical starburst galaxy; the optical and infrared continua are dominated by young stars and free- free emission, and there is little evidence for an evolved stellar population or nonthermal radio emission from supemovae.
There is now considerable evidence, especially from mid infrared spectrophotometry (Phillips, Aitken and Roche 1984, Roche and Aitken 1985) that small grains are an important part of the interstellar medium in non-active galaxies. Brackett line spectroscopy has been used to investigate this phenomenon. Assuming that the 10 pm emission in galaxies arises from dust heated by L a photons, the 10 pm flux density can be related to the free-free radio flux density by equation 1.22. The ratio between the 10pm flux density and radio flux density is dependent on the dust emissivity and dust temperature adopted. A maximum value for this ratio can be obtained by assuming that the dust is at a temperature of 200K and has an emssivity ® u 2 (Ho, Beck and Turner 1989). The emissivity is unlikely to be steeper than this, and so this value is an upper limit. The free-free radio flux can then be related to the
51 By line flux via equation (1.19) to give a relationship between the By flux and 10pm flux density: 1 1 £ F10 ------___ (1 .2 4 ) _mJy -x 10'17 Wm-2-
The relationship between the 10 pm flux density and the By line flux has been investigated for a number of galaxies by Beck, Turner and Ho (1986) and Ho, Beck and Turner (1989), who found that this ratio was larger (by an order of magnitude) than that seen in galactic HE regions. The authors concluded that the excess 10 pm flux arises from non-equilibrium emission from small grains, which emit strongly at 10 pm when they are heated to a temperature of 300K by the absorption of a single photon.
To sum up, the studies described above attempt to compare the luminosity derived from the ionizing flux with the observed infrared luminosity. With a few exceptions (Arp 220, NGC 6240) the bulk of the galaxies have ionizing fluxes that can be understood in terms of star formation. The 10 pm fluxes predicted on the basis of the By fluxes are larger than those seen in galactic HII regions. The excess emission may come from small grains. The number of ionizing photons derived from the thermal radio flux is consistent with that derived from the infrared recombination lines in NGC 253, however in M83 there is a line excess that can be explained if the burst is very young. In II Zw40 a continuum decomposition by Joy and Lester (1987) on the basis of the By flux shows that the old stellar population contributes very little to the K flux, and the radio emission is dominated by thermal emission from HII regions. II Zw 40 is therefore a young starburst.
1.8.2 Excitation Mechanisms for [Fell] and H2
The relative intensities of the U 2 lines of the first few K-window spectra of galaxies were consistent with thermal excitation behind a shock front (Thompson, Lebofsky and Rieke 1978, Fisher et al. 1983, Joseph, Wright and Wade 1984. There is evidence that the emission originates from a wide variety of physical processes including outflow from star formationregions and possibly supemovae, andinteraction-induced shocks. For example, the H 2 emission in NGC 253 is predominantly
52 excited by the starburst activity, both outflows from young stars and supernova remnants (Rieke, Lebofsky and Walker 1988, Walker, Lebofsky and Rieke 1988), the H 2 in NGC 6240 is probably associated with the merger process (Joseph,Wright and Wade 1984, Depoy, Becklin and Wynn-Williams 1986, Lester, Harvey and Carr 1987). In fact, Heckman et al.(1986) felt justified in asserting that the 1-0S(1) line they had observed from the nuclear regions of the Seyfert galaxy NGC 7469 was shock-excited, solely on the basis that every other extragalactic detection of the 1 -0S( 1) line was shock-excited.
Such confidence may be misplaced, since there is clear evidence forfluorescent excitation in galaxies. For example, Fischer et al. (1987) have obtained detections erf the 1 -0S( 1) line from the Seyfert galaxies NGC 1275, NGC 4151 andNGC 3227, and a detection of the 2-1 S(l) line in NGC 3227. They found that the relative intensities of the 1 -0S( 1) line and the 2-1 S( 1) line in NG£ 3227 were consistent with pure fluorescence. More recently spectra of several nearby spirals were obtained by Puxley, Hawarden and Mountain (1989), which show clear detections of the 2-1 S( 1) line and the 1 -0S(0) line with relative intensities consistent with fluorescent excitation.
In order to determine whether UV pumping is energetically feasible, Fischer et al. and Puxley et al. have calculated the ratio of the number of photons available for pumping H 2 molecules to the number required to excite the observed S( 1) line. The number of photons available for pumping H 2 molecules was calculated from the observed By flux. Fischer et al. (1987) found that pure fluorescence was energetically viable only in the case of NGC 4151 ( assuming that the UV continuum in NGC 4151 had a power law dependence) whereas Puxley et al. found that young stars could produce the required ionizing flux in all of their sample.
The forbidden line due to [Fell] XI .644 has been observed from a number of galaxies, for example Kawara, Nishida and Taniguchi (1988) and Moorwood and Oliva (1988). The ratio of the flux in the [Fell] XI.644 line to that in the By line is between 1 and 10 in all galaxies observed. This value should be compared to that predicted for photoionization of [Fell] (® 0.06) and shock excitation of [Fell] (*20-
53 50). In order to produce the observed ratios a mixture of shock excitation and HII regions are required; probably most of the By line is photoionized in HII regions and most of the [Fell] from shocks. The enhancement of [Fell] relative to By compared to that observed in HII regions points to a substantial shocked component to the [Fell].
1.8.3 Spatial Extent of Infrared Lines
Very little is known aboutthe spatial distribution of infrared line emission. The size of the emission region in both the S( 1) line and the By line in NGC 1068 is » 200 pc (Thompson, Lebofsky and Rieke 1978 and Hall et al. 1981). Lester, Harvey and Carr (1987) have measured the Gaussian scale size of the l-S(l) line in NGC 6240 and found it to be « 1.7 kpc, and Beck and Beckwith (1984) have obtained crude By line maps of NGC 253 and found that the By line is extended * 500 pc, and is coincident with the 1 Ojinj emission. Indirect estimates of the spatial extent of the S( 1) line in galaxies have been made by Kawara Nishida and Gregory (1987), who plotted the S(l) flux (normalised to distance by the bolometric luminosity) versus aperture/diameter for a sample of high luminosity IRAS galaxies. There is a clear trend for the flux to increase with aperture/diameter indicating that the emission is extended. They also suggest that Seyfert galaxies have emission that is more extended than other galaxy types. This suggestion should be treated with caution for a number of reasons; the S( 1) line flux in NGC 1068 is largely confined to ~ 200 pc compared to a much larger scale size of 1.7 pc for NGC 6240, and several of active galaxies observed by Kawara et al. have optical spectra characteristic of composite objects, and the extended emission may be due to a circumnuclear starburst
In summary, the spatial extent of the S( 1) and By lines has been estimated for a few individual galaxies, and as yet no clear trends have emerged. Kawara, Nishida and Gregory (1987) have suggested that the S( 1) line flux is more extended in Seyfert galaxies than in starburst galaxies, although this has yet to be confirmed by multiaperture spectroscopy.
54 1.8.4 Linewidths of Infrared Lines
Most of the detections of infrared lines have been obtained at relatively low resolution i.e. X/AX * 120, and the lines remain unresolved. The width of the 1 -0S( 1) line has been measured in a number of galaxies; for example in NGC 1068 it is *450 km s -1 (Hall et al. 1981), in NGC 6240 * 600 km s -1 (Joseph, Wright and Wade, 1984) and in NGC 4151 < 900 km s_1 (Fischer et al. 1987). Moorwood and Oliva have presented higher resolution detections of the 1 -0S( 1) line in a number of galaxies and they find it to be * 300 km s_1, except in NGC 6240 where it is rather wider. Thus the 1-0S(1) line is » 300 km_1 s in m ost galaxies, irrespective of the nature of the nuclear activity.
The widths of the infrared recombination lines have also been measured in a number of galaxies. In non-Seyfert galaxies the widths of the recombination lines are comparable to the width of the 1 -0S(1) line (e.g. Moorwood and Oliva 1988). However, in Seyfert galaxies there is a tendency for the B a and By lines to be broader than the H 2 lines; for example Hall et al. (1981) have measured the width of the By line in NGC 1068 to be * 1000 km s_1, and Fischer etal. (1987) estimate the width of this line in NGC 4151 to be * 1500 km s_1. DePoy, Becklin and Wynn-Williams have measured the width of the P a line in Arp 220 to be * 1300 km s_1. The fact that Seyfert galaxies have systematically broader recombination line widths than do non-Seyferls suggests that the ionization mechanism may be different for the two galaxy types.
1.8.5 Overview
The results described in sections 1. 8 .1 -1 .8 .3 show that infrared recombination lines have been used with considerable success to probe the energetics and dynamics of star formation in galaxies, even though the data presented is sometimes of poor quality. This is largely because the theoretical basis of recombination line spectroscopy is well understood, as is the relationship between recombination line spectra and infrared and radio continua in HII regions. This is in contrast to molecular hydrogen line spectroscopy, which is much less well developed as a tool for extragalactic research. This is partly because detailed theoretical models of H 2 emission spectra have only just become
55 available, but also because detailed modelling requires good detections (preferably at high resolution) of several lines, especially the 2-lS(l), 1- 0(0), and the Q-branch. Some attempts at such modelling have been made by Lester, Harvey and Carr (1987), and Puxley, Hawarden and Mountain (1988) but most papers present only detections of the 1 -0S( 1) line (e.g. Moorwood and Oliva 1988, Kawara, Nishida and Gregory 1987). Detailed theoretical models of infrared [Fell] emission, giving details of line strengths under different physical conditions, are non existent. Therefore, while it is possible to use infrared recombination line spectroscopy to model extragalactic star formation in a quantitative way, H 2 and [Fell] infrared spectroscopy is still largely exploratory.
1.9 Objectives
In this thesis I present H and K window spectra of a variety of galaxies - interacting and merging galaxies, as well as a couple of Seyfert galaxies and ordinary spirals. This work' has three major objectives. The first is to determine the spatial extent of the infrared lines. Multiaperture spectroscopy has therefore been obtained for a number of galaxies, especially luminous merging and interacting galaxies and two Seyfert galaxies. This work is largely exploratory. The objective is to compare the spatial distribution of the infrared lines in galaxies of differing morphology, and to contrast them with other line and continuum maps at other wavelengths (for example CO and 10pm emission). The second is to use measurements of the Byline flux to investigate extragalactic star formation, especially in very dusty galaxies where recombination lines are heavily obscured. Finally, the origin and excitation of the infrared lines of a large number of galaxies will be investigated using statistical techniques to lookfor correlations between the infrared line and continuum fluxes.
56 Chapter 2 Observational and Data Reduction Techniques in Infrared Spectroscopy
HE data presented in this thesis were obtained using a CVF spectrometer on the 3 .8 m United Kingdom Infrared Telescope (UKIRT) situated on Mauna Kea in Hawaii. These observations were made with two points in mind. The first was to obtain good spectral coverage (especially at K) to enable an accurate continuum to be fitted, and to get an estimate of the 2-1 S( 1) line flux to determine the H 2 excitation mechanism. The second was to acquire multiaperture spectra to measure the spatial extent of the infrared lines. In this Chapter I will explain why the CVF spectrometer was chosen for these objectives, and describe the observational and data reduction techniques employed.
2.1 Circular Variable Filter (CVF) Spectroscopy
Figure 2.1 shows a typical layout for a CVF. The CVF is a thin film optical interference filter deposited on a circular substrate. The thickness of the film varies linearly with angular position. Thus the CVF acts as a series of filters, with bandpasses approximately equal to the resolution of the interference filter. There are three major advantages to be gained in using the CVF for extragalactic spectroscopy. The first is that large (up to 19.6 arcseconds) apertures are available. Secondly, the UKT 9 CVF spectrometer employs reimaging optics, with the CVF itself in a collimated beam which means that the resolving power (X/5X) is independant of aperture. This is ideal for obtaining multiaperture observations of low surface extended sources. The second is that the CVF is rather more sensitive than other spectrometers, for example grating spectrometers. This is illustrated in the example below, where the performance of the UKIRT UKT9 CVF is compared to that of the UKIRT Cooled Grating Spectrometer (CGS) 2. The values for the sensitivities for these instruments were taken from the Starlink on-line users manual in July 1986.
57 Figure 2.1 Design of circular variable filter, taken from Yen (1969).
Consider a line in the K-band (* 2 Jim) of flux density 5xl0"17 Wnr2, and a continuum of 5x10-13 Wnr2 unr1 (See Figure 2.2). These values are typical for an extragalactic source. It is further assumed that the line has a width of« 300 km s_1, which is typical of stellar rotational velocities in galaxies, and corresponds to * 2x10-3 pm at this wavelength. The resolution of the CVF at 2 tim is 0.018 nm, and so this line is unresolved by the CVF. A S/N of 3 on the line therefore requires the noise in an integration to be » 1.7x10-1*7 W nr2 or a noise level of 1.7xl0-l7/0.018«9.2xl0-16 Wnr2 pm‘l. The CVF has a la ls sensitivity of 4.7x10-15 Wnr2 pm-1, and so an integration time of - 25 s is required to reduce the noise to the required level. This assumes that die CVF is background limited (see below), so that the signal to noise ratio increases as the square root of the integration time. The CGS 2 300 lines/mm grating has a resolution of 9x10-3 11m, so the line will still be unresolved. The lo ls sensitivity is 3x10_ 14 Wm*2 pm’l so a 30 detection of the line requires an integration time of M 220 s.
58 Intensity
- / W avelength
Figure 2.2 Schematic diagram of infrared line and continuum.
It is clear from the above comparison that the increased sensitivity of the CVF cuts down the integration time required to get reasonable detections. This is especially true when a large spectral range is required. To cover the whole K-window with the CVF with the sensitivity considered above, oversampling by a factor of 3, requires 40 minutes integration time. With CGS 2 the whole K window, oversampling by a factor 3,would take * 90 minutes, even allowing for the 7- element array. This increased sensitivity with the CVF is at the expense of the higher spectral resolution available with CGS 2.
2.2 Observing Techniques
Observing techniques in infrared astronomy are dominated by thefaetthatthe surroundings, such as the atmosphere and telescope, emit thermal radiation at infrared wavelengths. Atmospheric emission introduces noise into the detector in two ways; there are statistical fluctuations in the number of thermal background photons and there is ’sky noise’. The term sky noise is used to mean temporal changes in the background flux (associated with atmospheric turbulence) and spatial gradients in the background flux across the sky.
59 Detector
Prim ary m irror, Emisivity = ep Reflectivity = Rp Transmission = Tp Secondary m irror, Emisivity = Ej R eflectivity = Rs Transmission = Ts Mirror, Emisivity = t m R eflectivity = Rm Transm ission = Tm
Dewar Window Emisivity = ew Reflectivity = Rw Transm ission = Tw
Figure 2.3 Schematic diagram of a Cassegrain system
61 P s k y "" ^sky IxdXdQdA
2*2 where Ix is the intensity of a K blackbody at awavelength X, £sky the sky emissivity, dX the bandpass, and
dOdA is the throughput of the telescope. Assuming that the throughput of the detector is matched to that of the telescope, the power incident on the detector from the mirrors and dewar window becomes
Ptel - [Sp RsRmTw+ £b R mTw + EmTw + ejj I IxdXdOdA /A.X.Q where the symbols e, R and T represent the emissiviy, reflectivity and transmission of the mirrors and dewar window, and are defined in Figure 2.3. Adopting emlssivites for the sky and mirrors of5 %, sa l transmission for the dewar window of 95%, then the total power incident on the detectors becomes
Ptot Ptel + Psky = (0.34 + 0 . IxdXdOdA
0 ' 0 l % evaluating the integral for a 19.6 arcsecond aperture and abandwidth of C - pm (the _ fe£5 Urn of the CVF), Ptot becomes K^xlO'1 W. The r.m.s noise due to this background after an integration time of t seconds is given by:
AP = - 1/2 P
An integration time of 30 seconds gives AP = 6 x 10-18W. Therefore the flux incident at the top of the atmosphere which is measurable with a 1 astandard deviation is
O ' l * -A x IQ *8 s x l Q w 11.3x 3.6x 10*9x 0 'Ip m2 m and 0-2 lha ej^cieny cj. Vv where 0.6 is the transmission of the telescope and atmosphere^This can be compared dedzdto C. with the measured noise of « 2 x 10-11 W m ^ 1 in the 19.6 arcsecond spectrum of , u or ajoc^l' 3os point . NGC 6240, for which the integration time was **30-40 minutes, The agreement h 62 between the calculated and measured noise shows that the CVF is background limited even in the largest apertures.
2.3 Spectral Calibration
2.3.1 Overall Strategy
The main problem involved in infrared spectroscopy is die removal of atmospheric absorption features. These features are mainly due to water vapour, and are particuladyprominentatl.8-2pmandthenat2.3-2.5pnL This is illustrated in Figure 2.4, taken from Low and Rieke (1974).
Figure 2.4 Atmospheric transmission in the H and K windows.
The standard approach to this problem involves observing a source (usually a star) .which does not have spectral features at the wavelengths of interest, and which is at a similar airmass to the galaxy. The galaxy spectrum is then ratioed with this spectral standard, thus removing the atmospheric contribution. The standard star is assumed to have a continuum that is a blackbody with a temperature equal to the effective temperature of the star. This stellar continuum is removed from the ratioed galaxy spectrum by multiplying by a blackbody of the appropriate temperature.
Problems arise when the standard has intrinsic features, or when the atmospheric features vary in strength between the times of observing the galaxy and standard. The latter problem crops up when observing through thin cirrus, and can cause imperfect atmospheric cancellation. A standard star was observed before and after each galaxy integration to avoid this. Emission or absorption features in the standard star will produce spurious features in the galaxy spectrum; this effect and the method used to eliminate stellar features will be more fully described in the next section.
63 Spectral standard stars were selected to have the minimum number of intrinsic features in the 1-2.5pm region. Hot stars (A and earlier) have strong Brackett line absorption in their spectra. This is particularly inconvenient when observing the By line in nearby galaxies. There are also Brackett transitions in the H-window, although these have equivalent widths less than that of By. Cooler stars have more molecular absorption bands, in particular the CO rotation transitions at 2-2.5 pm and again in the H-window. The main aim of the work in this thesis is to determine the fluxes and spatial extents of the By recombination line, the 1-0S(1) line of H2 and the [Fell] transition. Therefore cooler G stars were chosen. A list of spectral standard stars observed is given in Table 2.1.
Table 2.1
Star Spectral Type Effective Temperature (SAO) (Kelvin)
129959 K0 5240 131873 G5 5610 84070 K5 4410 166647 K0 5240 28017 G5 5610 27438 G5 5610 158836 F4 6720 155739 K2 4780 82333 F5 6540 80511 G5 5610 84019 G5 5610 42876 G5 5610 81260 GO 5920 80511 G5 5610 130704 K2 4780 63462 K0 5240 121186 GO 5920 44230 GO 5920 121790 K0 5420 191840 G5 5610
64 2.3.2 Stellar Absorption Features
Consider the situation shown in Figure 2.5, a simple case chosen to illustrate the effects of stellar absorption features. Here the galaxy has no spectral features, but the standard has an absorption feature. The intensity as a function of wavelength for the standard is ^(X), the continuum level of the standard is C 2 , and the continuum level of the galaxy C j. The equivalent width of the standard feature is therefore
When the galaxy is ratioed with the standard the continuum level will be at a level Ci/Oj, but in the region of the absorption line the ratio will be C j/^. which will be larger than C j/C^. Thus a spurious emission line is produced in the ratioed spectra. The equivalent width of this spurious feature will be
M ^ L _ d X C2l AI(X)
65 Intensity
W avelength
Figure 2.5 Schematic diagram showing the effects of stellar absorption features.
For an unresolved stellar feature A I« C2 and so Wsp * W. The equivalent width of the spurious emission line will be die same as that of the standard absorption line but of opposite sign. Since it is the equivalent width of the line which remains die same, the flu x in the spurious line scales with the continuum level of the galaxy. Thus ratioing several spectra of the same galaxy taken in increasing apertures produces spurious emission lines with fluxes that increase with aperture, since the continuum level of the galaxy increases with aperture, if the galaxy fills the largest aperture. This effect can be mistaken for extended emission.
The following examples have been chosen to illustrate both of the effects described above; spurious emission features due to stellar absorption features and spurious extended emission. Figure 2.6 shows the 19.6 arcsecond H-window spectra of NGC 6240. The galaxy spectrum has been ratioed with the star SAO 121790. There are several prominent emission lines, with wavelengths (in microns) shown in Table 2.2. These lines may originate in four ways; they may be intrinsic galaxy features, standard absorption features, imperfect cancellation of atmospheric features, or instrumental features.
6 6 Relative intensity H-window spectrum of NGC 6240, ratioed with standard SAO 121790. SAO standard with ratioed 6240, NGC of spectrum H-window aeegh (^m) Wavelength iue 2.6 Figure 67 Table 2.2
NGC 6240 SAO 121790 CO Bands Brackett Lines (E m ission) (Absorption)
1.621 1.618 1.612 1.652 1.652 1.639 1.641 1.693 1.693 1.661 1.681 1.742 1.744 1.683 1.737
Figure 2.7 shows the spectrum of the standard SAO 121790. There is only one scan, and so no errors are shown. No atmospheric correction has been applied. Note the absorption dips, for which wavelengths are shown in Table 2.1. They correspond to the wavelengths of the emission lines in Figure 2.5 to within one CVF step. It is unlikely that these features are atmospheric; inspection of the Atlas Of The Solar Spectrum (Debouille et al. 1988) shows no sign of features in this wavelength range with equivalent widths large enough to account for these features. Instrumental features are unlikely since these features have been seen in a large number of standard stars, and the equivalent width depends on the spectral type of the standard star. This is illustrated by comparing the spectrum of SAO 121790 (a K0 star) with SAO 191840 (a G5 star). The spectrum of SAO 191840 is shown in Figure 2.8. There are hints of absorption features at 1.61 and 1.64pm, butnot as prominent as in SAO 121790. Thus these features are intrinsic to the star.
6 8 Flux (10< widw pcrm 121790 O A S f o spectrum indow -w H Wavelength Wavelength iue 2.7 Figure 69 (fim) Relative intensity widw pcrm A 191840 SAO spectrum indow -w H iue 2.8 Figure aeegh (yum) Wavelength 70 Identification of these stellar features has been problematic. CO bands are frequently seen in the spectra of late-type stars; the wavelengths of the CO transitions that fall in the H-window are shown in Table 2.2. They clearly do not match up with the observed stellar transitions, and neither do the Brackett lines. Closer inspection shows that the wavelength difference in successive absorption dips in the standard matches the wavelength difference between successive Brackett line transitions. This is tentative evidence that the stellar features may be Brackett line absorption. It is generally accepted that hydrogen absorption is very weak in K stars and only becomes noticeable in stars earlier than G, therefore this suggestion must be treated with caution. Higher resolution spectroscopy must be used to identify these features unambiguously.
The stellar features have been removed by fitting a low order polynomial to the standard spectrum, effectively smoothing out higher order spectral frequencies. This ’synthetic atmosphere’ was then used to ratio the galaxy spectra. The results for the NGC 6240 H-window spectrum are shown in Figure 2.9. The emission line at 1.652pm has gone, and was thus entirely spurious. The peak of the ’emission line’ at 1.744pm seen when the galaxy was ratioed with SAO 121790 has shifted to 1.734pm. It is likely that this point is on the continuum since there are redshifted CO bands at 1.725 and 1.747pm. The final identifications are shown in Figure 2.9.
71 Relative intensity widw pcrm G 20 aie t h s hei t phere' h sp o atm etic th n 'sy the ith w ratioed 6240, NGC f o spectrum indow -w H . (- 7.5 8.5 - 8.5 1.6 i ' ------1 i ------______1------1 ! ______'''111 ------1.65 1 ------aeegh (/mm) Wavelength iue 2.9 Figure r ------72 1------1 ------r 8-6 CO8-6 1.7 t ------1 ------r 1.75 Figure 2.10 illustrates the spurious flux increase with aperture that results from stellar absorption features. This Figure shows the flux in the [Fell] X 1.644 line asa function of aperture. Four aperture sizes are shown, 5.4,7.8,12.4 and 19.6 arcseconds. These spectra will be described in more detail in Chapter 3, and are not shown here. The dotted line connects the fluxes measured from spectra where SAO 121790 is used as a standard, and the solid line connects the points where the ’synthetic atmosphere’ has been used. The flux in the smallest aperture is the same for the synthetic atmosphere and the star. However the difference becomes apparent in larger apertures since the spurious flux introduced by the stellar features scales with the continuum level of the galaxy.
73 Line Flux (10 [Fell] line flu x plotted as a function o f aperture fo r NGC 6240 ratioed w ith SAO SAO ith w ratioed 6240 NGC r fo f aperture o function a as plotted x flu line [Fell] 270a di ytei t osphere atm synthetic ie d d an 121790 prue (arcsec) Aperture iue 2.10 Figure 1A Stellar absorption features may be responsible for a discrepancy in the [Fell] X1.644 line flux in NGC 6240 measured by Lester Harvey and Carr.(1988) and the value which will be presented in in Chapter 4 of this thesis. Lester Harvey and Carr.(1988) measured aline flux of 10.5xl0-17 Wm-2 in a 3 arcsecond aperture, much larger than the value of 5.2xl0-17 Wm_2in a 5 arcsecond aperture I present in in Chapter 4. A calibration error is unlikely to account for this difference; the continuum level in the Lester etal. spectrum is « 2.7x10*14 Wm*2 pm-1, which is entirely consistent with the continuum level of 3.3xl0"l4 Wm-2 pm-1 in the spectra shown in Chapter 4. Lester Harvey and Carr report thatthe standard used as a standard to remove atmospheric features is a AO star, with strong By absorption. It is possible thatthe larger flux reported by Lester et al. is due to the presence of the N=11-4 Brackett line. This line has a wavelength of 1.681pm, which is coincident with the redshifted [FellJX 1.644 line at 1.685pm.
Fig 2.11 shows the K-window spectrum of NGC 1068, another example of spurious emission lines due to standard absorption. Here there is a series of emission bands from 2.3-2.5 pm. The wavelengths of these bands match up with the rest wavelengths of the 2nd overtone CO bandheads. These bands can be seen in absorption in the standard SAO 129959, shown in Figure 2.12. Once again using a synthetic atmosphere to ratio the galaxy spectrum successfully removes the emission bands, except for the 2-0 CO low R J-branch, which has also been seen in emission by Hall et al.(1981) and Thompson et al.(1978). The final spectrum of NGC 1068, complete with identifications, is shown in Figure 2.13.
75 Flux (Relative intensity) widw pcrm G 16,rtoe wihsadr A 129959 SAO standard ith w ed ratio 1068, NGC f o spectrum indow -w K aeegh (/xm) Wavelength iue 2.11 Figure 76 Relative intensity widw pcrm A 129959 SAO f o spectrum indow -w K aeegh (Mm) Wavelength iue 2.12 Figure 77 Flux (10 F inal reduced K window spectrum o f N G C 1068, w ith stellar absorption features absorption stellar ith w 1068, C G N f o spectrum window K reduced inal F . 22 . 24 2.5 2.4 2.3 2.2 2.1 aeegh (/xm) Wavelength iue 2.13 Figure e oved. rem 78 2.3.3 Instrumental Effects
Figure 2.14 shows a K-window spectrum of NGC 3256, which was observed at the Anglo-Australian Telescope using FIGS. Figure 2.15 shows the ratio of two observations the standard star BS 4013, one of which was taken before the observation of NGC 3256 and the other afterwards. It is clear there are a number of bands in this ratio which should be entirely flat This banding is caused by the spectrometer shifting in wavelength over timescales of approximately one hour, so that the peaks of the atmospheric features do not fall in the same pixels in both spectra. Shifting tiie the wavelength scale so they match considerably reduces the banding as can be seen in Figure 2.16. It is obvious that a shift in the wavelength calibration between the galaxy and standard will produce similar banding in the galaxy spectra. The banding is worst at the edge of the window, between 2.3-2.5 pm, which unfortunately coincides with the CO bandheads. The positions of the CO bandheads at the redshift of NGC 3256 are shown in Figure 2.15 along with the positions of the instrumental bands (marked A, B, and C). The obvious approach to this problem is to use atmospheric features as a wavelength calibrator to align the individual galaxy scans with the standard. The galaxy scans are too noisy to permit them to be aligned with sufficient accuracy to eliminate the instrumental features.
79 T T r T i— .— .— <—i— r Wavelength Microns T 1—'—'—r —■— t r
Figure 2.14 K-window spectrum of NGC 3256 ^rm c>i
Csi
oin o in — CX' &' Os! (U 1' o> 5 CN Csj m CM i i ] (s*nm jtreimqB) Xjisuajui Figure 2.15 Ratio of two spectra of the standard star BS 4013, taken about an hour apart 81 / C\* r ~ TT -/r ~ (N h 7 ) ■ l - \ ) J L ; \ ) \ ) ) 1 ti \ L r ! J?i ►i l i j _ ( S i ■ j rv ) ) \ l rI !' o - - l w (siiun jCreimqE) jfysuajui Figure 2.16 Ratio of two observations of BS 4013, with wavelength scale aligned. 82 2.4 Flux Calibration and Calculation of Errors. The scans obtained after each nod cycle were inspected for spikes, the mean levels adjusted to the same level and then coaddecLand smoothed. The error on each point is the standard error obtained after coadding. The absolute flux level was set with reference to the spectrum of a standard star at a similar airmass to die galaxy. The spectrum of the spectral standard star was ratioed with that of the flux standard, and the absolute flux level of this ratio set by multiplying it by the absolute continuum level of the flux standard. Table 2.3 shows the standards used and the magnitudes adopted. Table 2.3 Star Kmagnitude H-K HD 40335 6.45 +0.02 BD +0°1964 4.585 +0.255 HD 22686 7.185 +0.005 HD 84800 7.53 +0.00 BD+2°2957 4.245 +0.18 GL 390 6.045 +0.205 HD 136754 7.135 -0.005 HD 105601 6.685 +0.03 HD 162208 7.11 0.035 Magnitudes were converted to flux densities assuming that the K band centre is at 2.22 pm, has a width of 0.52pm and a zero point flux of 667Jy (Campkins, Rieke, and Lebofsky 1985). The value assumed for the H-band centre is 1.6pm, the bandwidth 0.36pm, and the zero point flux 1075Jy (Campkins, Rieke and Lebofsky 1985) In order to measure line fluxes a continuum was removed from the reduced galaxy spectrum. Typically this was done by fitting a low order polynomial. An example of this procedure is given in Figure 2.17, which shows the K-window spectrum of NGC 6240 with the fitted polynomial. Figure 2.18 shows the K-window of NGC 6240 with the continuum subtracted. It should be emphasised that this is necessarily a subjective process. The line fluxes were then calculated by a least 83 squares fit of a Gaussian profile to the points lying above the continuum. In no case was a line found to be resolved (ie Av>2000 km"l) so the width of the instrumental profile was used as the FWHM of the fitted Gaussian. The line centre was initially derived from the galaxy redshift, and values for the flux and peak flux of the Gaussian were calculated. This calculation was repeated shifting the line centre to either side of the redshifted line, to find a Gaussian with the smallest residuals. The flux in the Gaussian with the smallest residuals was adopted as the line flux. The smallest residuals were usually obtained when the line centre was within one CVF step of the redshifted line. The errors quoted on the fluxes are statistical errors resulting from the fit Figure 2.19 shows the Gaussian fitted to the 1-0S(1) line in NGC 6240 and the residuals. 84 Relative intensity K-window spectrum of NGC 6240 with fitted polynomial fitted with 6240 NGC of spectrum K-window aeegh Gum) Wavelength iue 2.17 Figure 85 Flux (10 NGC 6240, continuum subtracted. continuum 6240, NGC vlnt (jum) avelength W iue 2.18 Figure 86 Flux (10 G aussinan fit to th e NGC 6240 S<1) line, w ith residuals. ith w line, S<1) 6240 NGC e th to fit aussinan G aeegh (/^m)Wavelength iue 2.19 Figure 87 Chapter 3 Near Infrared Spectroscopy of Merging Galaxies 3.1 Introduction OST galaxies can be classed as one of two basic morphological types; either Mspiral or elliptical. However, as demonstrated by Arp in his 'A tla s o f Peculiar Galaxies'j a small fraction of galaxies have morphologies that can be quite bizarre. Some of these irregular systems have features such as tidal tails (for example NGC 3256) or ripples (NGC 3310) which have been identified as characteristic of two merging galaxies. (Toomre and Toomre 1972, Schweizer 1983). The effect of such collisions became apparent when several studies showed that interacting and merging galaxies have higher luminosities than isolated systems, with the luminosities in mergers as high as 1012 L©. (Joseph e t al 1984, Joseph and Wright 1985, Lonsdale, Persson and Matthews 1984, Cutri and McAlary 1985). The study of merger remnants became very popular in the late 1980’s after the IRAS mission, when it was realised that two of the most luminous galaxies detected by IRAS, NGC 6240 and Arp 220, were well known mergers (Wright, Joseph and Meikle 1984, Soiferetal 1984a). Since then, optical imaging of IRAS galaxies has shown that many are distorted, have faint tails or double nuclei - characteristics usually identified with interactions and mergers (Arums, Heckman and Miley, 1987, Sanders et al. 1987, Sanders et al. 1988, Hutchings and Neff 1987, Lawrence et al. 1989). Studies of the space density of luminous (1011 - 1012 Lo) IRAS galaxies in the local universe (z < 0.081) show that at the high end of the luminosity function (> 2.5 x l0 12Lo) IRAS galaxies are as common as quasars (Soifer et al. 1986). This study also suggests th at« 0.05 % of ’normal’ galaxies may have infrared luminosities an order of magnitude larger than their visual luminosities. If this is the case, and the high luminosity phase lasts * 108 years, then about half of all normal galaxies must go through a high luminosity phase. The origin of the luminosity found in mergers is controversial, with three main competing theories. The first is the starburst scenario, discussed by Joseph and Wright (1985). As two gas rich galaxies start to merge, cloud-cloud collisions induce fast shocks in the molecular gas. This gas cools, collapses and fragments, producing a burst of star formation. The discovery of shock excited molecular hydrogen in the two merging galaxies NGC 6240 and Arp 220 (Joseph, Wright and Wade 1984, 8 8 DePoy, Becklin and Wynn-Williams 1986) is consistent with this scenario and further observations will be discussed in this chapter. The main rival to this theory is that the infrared luminosity is produced by a dust embedded active nucleus, the merger of two gas rich galaxies providing the ’fuel to feed the monster’ (DePoy, Becklin and Geballe 1987, Scoville et al. 1986). Finally, there is the suggestion that the interaction is the source of the luminosity (Harwit et al. 1987). There has been speculation that there is an evolutionary link between starbursts and active nuclei, and that possibly even high luminosity AGN’s and quasars were formed from a starburst (Norman and Scoville 1988, Sanders et al. 1988). Optical imaging of the host galaxies of quasars indicates that they are often distorted and disturbed, lending some credibility to this idea (Hutchings and Neff 1988). Galaxy mergers may be instrumental in the formation and evolution of active galactic nuclei and mergers may dominate the space density of high luminosity objects. This chapter presents H and K window spectra of a number of merging galaxies. The objectives of this study are to determine whether there are sufficient ionising photons (as measured by the By line flux) to produce the observed infrared luminosity through star formation, and to study the spatial distribution of the shocked gas and its relation to the luminosity source. 3.2 Data and Sample Selection The galaxies discussed here have coalesced into a single object and show morphological signs of merging-for example tidal tails or shells. The data have been reduced as described in chapter 2. Figure 3.1 shows the spectra discussed in this chapter. Table 3.1 shows the line fluxes from the data presented here. Data for NGC 3256, NGC 4194, NGC 520 and the 5 arcsecond K-window data for NGC 6240 have been taken from the literature. 89 The flux adopted for the By recombination line in NGC 6240 requires some explanation. The 2 pm spectrum of Lester, Harvey and Carr (1988) shows a detection of the 2- 1S(2) line, which has a rest wavelength of 2.154 pm. This line has a flux of 1.2 xlO-17 W m -2 in their 3 arcsecond aperture, compared to the By line strength of 0.6x10-17 W n r2 These two lines are blended in the lower resolution spectra presented here. The combined flux in these two lines in the 19.6 arcsecond aperture is 2.6xl0-17 W n r2. The flux in the unblended 1-0S(1) line has doubled in going from the 3 arcsecond aperture of Lester, Harvey andCarr.(1988) to the 19.6 arcsecond aperture measurementpresented in this thesis. Assuming that the excitation conditions across the disk of NGC 6240 are constant, the flux in the 2-lS(2) line should also be double that of the Lester, Harvey and Carr (1988) measurement This assumption implies that the 19.6 arcsecond 2-1S(2) line flux should be “2.4x10-l7W m-2. This is equal to the total flux seen in the By plus 2-1 S(2) blend. Thus there is little evidence that the By flux contributes significantly to this blend in the larger aperture. I have therefore adopted the By flux in the small (3 arcsecond) aperture of Lester, Harvey and Carr (1988). 90 Table 3.1 Galaxy Aperture Line Flux N ote (arcseconds) (xl0-17Wm-2) NGC6240 19.6 S (l) 26.810.7 1 By+2-lS(2) 2.610.4 1,2 [Fell] 7.8±0.7 1 12.4 [Fell] 7.010.7 1 7.8 S (l) 23.210.4 1 [Fell] 5.610.4 1 5.5 S (l) 1510.1 3 [FeD] 5.210.5 1 Arp 220 19.6 S (l) 5.210.5 1 By 1.310.5 1 7.8 S (l) 3.110.2 1 5.5 S (l) 2.510.5 3 NGC 1614 19.6 S (l) 4.510.5 1 By 4.610.8 1 [Fell] 12.711 1 12.4 [Fell] 8.510.8 1 7.8 [Fen] 7.910.1 1 NGC 6052 19.6 S (l) 1.510.2 1 By 3.010.2 1 NGC 2623 5.5 By 1.2i0.3 4 NGC 3256 13.5 By 14.81.2.2 5 NGC 520 5.5 By 1.910.4 4 NGC 4194 5.5 By 4.310.4 4 Notes 1 Data from this thesis 2 See text 3 Data from Joseph, Wright and Wade (1984) 4 Data from DePoy (1987). 5 Data from Moorwood and Oliva (1988) 6. Kawara, Nishida and Gregory (1987) 91 Wavelength (/xm) Figure 3.1a (i) 19.6 arcsecond K-window spectrum of NGC 6240 92 Flux (10 7 .8 arcsecond K -w indow spectrum o f NGC 6240 NGC f o spectrum indow -w K arcsecond .8 7 iue .a (ii) 3.1a Figure vlnt (i ) (/im avelength W 93 Flux (lO"'9 Wm“2//m 19.6 arcsecond H-window spectrum of 6240 NGC of spectrum H-window arcsecond 19.6 iue .a (iii) 3.1a Figure aeegh (^m) Wavelength 94 Flux (10“,s W m-Vn 12.4 arcsecond H-windaw spectrum of NGC 6240 NGC of spectrum H-windaw arcsecond 12.4 aeegh (/um) Wavelength iue .a (iv) 3.1a Figure 95 Flux (KT19 Wm 7.8 arcsecond H-window spectrum of NGC 6240 NGC of spectrum H-window arcsecond 7.8 aeegh (/um) Wavelength iue .a (v) 3.1a Figure 96 Flux (10 5.5 arcsecond H-window spectrum of NGC 6240 NGC of spectrum H-window arcsecond 5.5 aeegh (jxm) Wavelength iue .a (vi) 3.1a Figure 97 Flux (1015 Wm'2/im 19.6 arcsecond arcsecond 19.6 aeegh (/xm) Wavelength iue .b (i) 3.1b Figure widw indow -w K 98 spectrum spectrum o f f o Arp 220 Arp Flux (10 7.8 arcsecond K-window spectrum of 220 Arp spectrum of K-window arcsecond 7.8 iue .b (ii) 3.1b Figure vlnt (mi) (/m avelength W 99 Flux (10 96 rscn -idwsetu fNC 1614 NGC of spectrum K-window arcsecond 19.6 aeegh (//m) Wavelength iue .c (i) 3.1c Figure 100 Flux (10 19.6 arcsecond H-window spectrum of NGC 1614 NGC of spectrum H-window arcsecond 19.6 aeegh (/^m) Wavelength iue .c (ii) 3.1c Figure 101 Flux (10 12.4 arcsecond H-window spectrum of NGC 1614 NGC of spectrum H-window arcsecond 12.4 iue .c (iii) 3.1c Figure aeegh (/xm) Wavelength 102 CM 1 I Flux (10 B a. B 4 - -2 6 ------1.6 «- . aceodHwno pcrmo G 1614 NGC of spectrum H-window arcsecond 7.8 aeegh (/urn) Wavelength iue .c (iv) 3.1c Figure O O CO CO CO I 1 I I 1.65 103 F I] F H] [Fe II] [Fe 1.7 B Flux (10 B a. .5 . 21 22 .5 2.3 2.25 2.2 2.15 2.1 2.05 —i —i —i I— i— i— i— i— — t — 19.6 arcsecond K -w indow spectrum o f NGC 6052 NGC f o spectrum indow -w K arcsecond 19.6 —i i— i— i— j —i —i —i —i —| —i —i —i —i r i— i— i— |— i— i— i— i— |— i— i— i— i— ]— i— i— i— i— — ------1 ____ aeegh (^m) Wavelength i ____ iue .I (i) 3. Id Figure i ____ I ____ 104 i ____ i ____ i ____ I i I i ____ i ____ i ____ i ____ I ____ i'i' 3.3 Brackett Line Spectroscopy of Mergers The aim of this section is to determine whether the number of OB stars inferred from the flux of ionising photons can produce the observed far infrared luminosity in a sample of merging galaxies. (It is assumed that the infrared luminosity originates from dust heated by stars in the starburst, as outlined in Chapter 1). This is done by using a simple starburst model to compare the star formation rate implied by the infrared luminosity measured by IRAS, to the star formation rate inferred from the number of ionising photons measured by the By flux. 3.3.1 The Starburst Model The model chosen was developed by Telesco and Galley (1984) and further used by Wright et al. (1988) in their study of interacting and merging galaxies. It is assumed that the starburst started To years ago. The initial mass function (the number of stars formed per year in the mass range between M and M+dM) of the burst is W (M )dM . The starburst has proceeded with a uniform star formation rate since the burst began. The luminosity of the burst at time To is given by: /■MO -M2 LfTo) Ls(M)^(M)T0dM + Ls (M) ¥(M) T^ (M) dM (3.1) /Ml J Mo where Mi and M2 are the lower and upper mass mass cut offs of the stars in the burst, and Mo is the mass of a star with a main sequence lifetime To. Ls(M) is the luminosity of a star of mass M, and Tms(M) the lifetime of a star of mass M. The first term of equation (1) represents the luminosity from stars whose main sequence lifetime is greater than the age of the burst The second term represents the contribution to the luminosity from stars whose main sequence lifetimes are less than the age of the burst and have thus come into equilibrium, with the stellar birthrate equal to the number that have evolved off the main sequence. A similar equation can be written for the number of ionising photons from the burst at a time To- *M0 -M2 N(T0) = I Ns (M) V(M) T0 dM + I Ns (M) 'F(M) (M) dM (3.2) J M i /M o where NS(M) is the number of ionising photons from a star of mass M. 105 For simplicity the stellar quantities are represented as a function of mass by power laws: V (M ) = KMV (3.3a) Ttns(M) = BMP (3.3b) Ls(M) = AM° (3.3c) Ns(M) = DM5 (3.3d) The values of the constants D, 5, y, B, p, A and a adopted for the the mass range 0.1- 60 M q . are shown in Table 3.2. The constant K sets the scale of the starburst, and is derived from either from the ionising photon flux, or the infrared luminosity. Details of the derivation of these constants are given in Appendix A 1. 106 Table 3.2 Mass Range A a B P D 5 Y (Mo) (Lo) (xlO6 years) (photons s '1) M <0.7 0.5 2.7 10* -3 .0 0.0 0.7 Mass-luminosity 5 Mass-main sequence allM Scalo 1986 lifetim e IMF allM Scalo 1986 Luminosity- allM Panagial973 ionising flux 107 The expressions above for the ionising flux and luminosity of a starburst can be simplified by assuming either that the starburst is very young (“ 106 years) or that the starburst is sufficiently old that the formation of OBA stars has come into equilibrium (i.e. Mo *»2Mo) and the stellar birthrate for these stars equals the deathrate. The star formation rate inferred from the infrared luminosity will be compared to that derived from the number of ionising photons in the limit of an old starburst and a young one. (i) Unevolved Starburst If the starburst is very young (the Zero Age Main Sequence approximation) then the second integral in both equation (1) and equation (2) can be neglected and the luminosity and ionising flux are given by: [MO L = Lg (M) Y(M) T0 dM = KT0 FUR(M, ,Mo) (3.4) J'Mi M l w here m M2 FUr(M 1,M2)= Ls(M)V(M)dM (3.5) J*Mi M i and .M o N(T0) = Ns (M) W(M) T0 dM = KT0 F ion (M i ,M 2) (3.6) J M i with -MO F ion (M,,M 2)“ I Ns(M)^(M)dM (3.7) / M i Note that Mo = M2 for a very young starburst. The values of the integrals FioNafld F t ik depend on the upper and lower mass cut off of the burst The integral Fjo n is essentially independent of the lower mass cut off because stars below 5 M o do not contribute to the ionising flux. For a ZAMS burst with M2 = 40 M o, F jo n “ 4.2 x 1051 s-1, and with M2 = 60 M o, Fjo n = 8*8 x 1045 s'1. The integral Fu r is w eakly dependent on the lower mass cut off, and much more strongly dependent on the upper mass cut off. Table 3.3 shows values of F t.tk fo r M i = 0 .1, 2 and 5 M o and M2 = 40 and 60 M o 108 Table 3.3 F u r as a function of upper and lower mass cut off for a ZAMS starburst M^= 40 Mo 60 Mo Uvw’-fe '• i-*0 Mi = 0.1 M o 263.1 348.9 2.0M © 234.4 346.8 5.0M© 231.0 329.2 The star formation rate can then be derived independently from both the infrared luminosity and the number of ionising photons as M u r - ¥(M )M dM = K u RMY+1dM (3.8) M ion “ ¥ (M )M dM « KioNMY+1dM (3.9) H ere Mu r is the star formation rate deduced from the infrared luminosity and Mio n is the star formation rate inferred from the number of ionising photons. The constant K that sets the size of the burst has been writtenK t.tr if it has been calculated from the infrared luminosity and K io n if it has been calculated from the number of ionising photons. I can now define a parameter Q, with which to compare the star formation rate inferred from the infrared luminosity with that inferred from the ionising photon flux, as 3 .5 x 1 0 *48 [ h 1 ’N ion ' (3.10) M ion K ion F u r . .F io n . In equation (10) the luminosity, L, of the starburst is expressed in Lo> N ion is the number of ionising photons per second from the burst derived from the B y flux: 109 4 ttD2F By N ion ** 70 (3.11) hVB, where D is the distance to the galaxy, Ffiy is the observed By flux, and huBy is the energy of a By photon, and it has been assumed that **70 ionising photons are emitted for each By line photon (Brocklehurst 1971, Giles 1977). The value ofJ~. £as Seen calculated from the valu^ofthejntegrals F ion an d Fu r fo r M i = 0.1, 2 and 5 Mo, andM 2 “ 40 and 6OM0 J These are shown in Table 3.4. Ll/ w,ch] Table 3.4 as a function of upper and lower mass cut off A for a ZAMS starburst, u/i'tH Q ^ • M1 =40M © 60 M o Mi = 0.1 M o 5.. g . 2.0M© 6. S vhcT" 8. ;Ui0. 5.0M© 6. ic ~ 5 9„ From this table it can be seen that Q for a ZAMS starburst is insensitive to the lower mass cut off, and depends weakly on the upper mass cut off. This reflects the fact that young starbursts are dominated by the massive stars, with the low mass stars contributing very little to either the luminosity or ionising flux. (ii) Evolved starbursts If the starburst has evolved to the point where OBA star formation is in equilibrium then the luminosity of the starburst is given by: #M2 U J o ) - I Ls (M)W(M)Tms (M)dM = KURFUR (3.12) J Mo w ith J »M2 Ls (M)W(M)Tms (M) dM (3.13) Mo and the number of ionising photons is given by; 110 M2 f N s(M )^(M )T ms (M) dM - K iqn (3.14) where N s(M )^(M )T ms (M) dM (3.15) The integral F’ion Is insensitive to the lower mass cutoff of the burst, because stars below 5 M o do not contribute to the ionising flux. For an evolved starburstF’ion = 4.2 x 1046 for M2 = 40 M q and F’ion = 6.4 x IQ46 for M2 = 60 M o The integral F’tjr is very sensitive to the lower mass cut off of the burst, because lower mass stars have accumulated over the lifetime of the burst and contribute to the total luminosity. It is , however, insensitive to the upper mass cut off. The values of F’t ir are shown C 1—/ NlON> J inTable 3.5, and the values of for an evolved burst (defined in equation (10)) are shown in Table 3.6. A Table 3.5 L’t jr as a function of upper and lower mass cut off for an 2 evolved starburst ^ ur\ils , l— * M^ = 40MO 60 MO M i-0.1 M o 21949.1 25104.0 2.0MO 7029.0 10184.0 5.0MO 5141.4 5543.0 111 Table 3.6 | J as a function of upper and lower mass cut off for an evolvedstarburst } with ^ ~ I M,=40Mo 60 Mo Mi = 0.1 M o 6.7x10-6 8.8x10-6 2.0M© 2.0 xlO-5 2.2x10-5 5.0Mo 3.0 x 10-5 4.o x 10-5 It can be seen from Table 3.6 that for an evolved starburst is very sensitive to the lower mass cut off of the burst, reflecting the contribution of low mass stars to the integrated luminosity. (iii) Summary The starburst model of Telesco and Gatley (1984) has been developed to define a parameter Q, which is the ratio of the star formation rate inferred from the infrared luminosity and that inferred from the ionising fluk. ^ 'iiasieen calculated for ZAMS starbursts, and starbursts where sufficient time has elapsed for OBA star formation to come into equilibrium. The value of nofparticularly sensitive to either the upper mass cut off or the lower mass cut off for a ZAMS starburst, but is very sensitive to the lower mass cut off in an evolved burst 3.3.2 Extinction Corrections to By line In order for the number of ionising photons to be derived from the By flux, some correction must be made for the extinction to the By line. The extinction for most of the mergers discussed here has been estimated by comparing the observed recombination line ratios with that expected from case B recombination (see Chapter 1). Line fluxes for recombination lines other than the By line have been taken from the literature. Most of the Brackett line fluxes were corrected with reference to the Ba flux, except for NGC 3256 where the Hp flux was used, because there is no value for the Ba flux in the literature. The extinction at 2pm and the dereddened By fluxes for the mergers in this sample are shown in Table 3.7. The extinction in the galaxies NGC 6240 and Arp 220 is controversial, and will be discussed below. Table 3.7 Galaxy A2.2 Dereddened flux Ref (mag) (xlO‘17Wm-2) NGC 520 0.98 4.6 1 NGC 1614 0.42 7.9 1 NGC 4194 0.5 6.8 1 NGC 2623 0.5 5.2 1 NGC 3256 0.1 15.4 2 NGC 6240 0.3 0.8 See tex t A rp 220 0.5 1.3 tr 1 DePoy,1989 2. Veron- Cetty and Veron(1987) Measurements of the extinction in the nuclear regions of NGC 6240 vary considerably. Measurements of the 1 0 pm absorption feature by Rieke et al. (1985) indicates that the visual extinction is » 25 mag. Near infrared (J, H and K) colours suggest a visual extinction of »3 mag (Rieke et al. 1985). Finally, the By line flux of Lester, Harvey and Carr (1988) and the P a line flux of DePoy, Becklin and Wynn- Williams (1986) may be used to estimate the extinction assuming Case B recombination; this procedure suggests that the extinction at 2 pm is essentially zero. These estimates of the extinction suggest that the obscuring material is patchy. The estimate of Av -25 is almost certainly an overestimate of the extinction to the line emitting gas, as indicated by the lower estimates of the extinction from the recombination lines and infrared colours. There is also considerable uncertainty in the calibration of the 10 pm feature. An extinction of Av = 3 mag has been adopted, implying an extinction at 2.2pm of 0.3 mag. The extinction in Arp 220 is also controversial. Measurements of the 10pm feature in a 5 arcsecond aperture and the CO column density in the central 4 arcseconds give values of Av=50 and 130 respectively (Becklin and Wynn-Williams 1986, Scoville et al. 1986). These are likely to be over estimates of the extinction to the line emitting gas, as discussed above. The Ba Auxin a 5 arcsecond aperture of Depoy, Becklin and Geballe (1987) can be combined with the upper limit to the By flux in an 8 arcsecond aperture presented here, to estimate the extinction assuming 113 case B recombination.. The Ba flux of DePoy et al. implies that the By flux in a 5 arcsecond aperture is* 1.6x10"17 Wm-2 for zero extinction. The upper limit to the By flux in the 8 arcsecond aperture is *lxl0*17 Win*2, giving a nuclear extinction of Av » 5 mag. This is consistent with the estimate of the nuclear extinction of Av * 5 mag given by Norris (1985) from the near infrared colours. The question of the extinction is further addressed in section 3.4.4, where multiaperture JHK photometry is presented to calculate the differential extinction between the 19 arcsecond and 5 arcsecond apertures. The average extinction in the 5 arcsecond aperture (as derived from JHK photometry) is Av = 5 mag. This is entirely consistent with the estimate from the near infrared recombination lines and the photometry of Norris (1985), and will be adopted here. 3.3.3 Results and Discussion Figure 3.2 shows the By flux plotted againstthe integrated infrared flux for all of the mergers in this sample. Also shown are lines of Q= 1 for a ZAMS starburst, and evolved starbursts with lower mass cut offs of 5Mo> and 0.1 Mo. It is clear that most of the mergers have values of Q that are consistent with a ZAMS starburst If the starbursts are evolved, then the lower mass cut off of the burst must be *5 Mo, a conclusion reached by Wright etal. (1988) in their study of interacting and merging galaxies. The two exceptions are NGC 6240 and Arp 220. The deficit of ionising photons (relative to the infrared luminosity) in these galaxies has been noted by previous workers, in particular by DePoy, Becklin and Wynn -Williams (1986) and Lester, Harvey and Carr (1988) for NGC 6240, and DePoy, Becklin and Geballe (1986) for Arp 220. These two galaxies will be discussed below. 114 Figure 3.2 By flux plotted against infrared flux for the mergers discussed in this sample. Also shown are lines of Q* 1 for a ZAMS starburst and two evolved starbursts, one with Mi* 5 Mo. and the other with Mi* 0.1 Mq . 115 The location of NGC 6240 in Figure 3.2 suggests that if the infrared luminosity is powered by a starburst then there is a large contribution to the integrated luminosity from low mass stars, perhaps indicating that die burst is more evolved than the other mergers in this sample. The position of Arp 220 is difficult to reconcile with the starburst models discussed here. It is also possible that Arp 220 and NGC 6240 are exceptionally dusty objects in which the dust successfully competes with the gas for Lyman continuum photons, so that the By flux seriously underestimates the ionising flux. This is observed in galactic HE regions (as described in Chapter 1), and is usually characterised in terms of the infrared excess, defined by (3.16) Llq hVLccNiON where Ljr is the infrared luminosity of the HE region, Lixt is the luminosity in the La line, and Nion is the number of ionising photons. The extinction corrected By flux and the IRAS FIR parameter for Arp 220 and NGC 6240 indicate that the infrared excess in these galaxies is -300 and 110 respectively. This is an order of magnitude larger than the value seen in HE regions. While it is not impossible to rule out dust absorption as the cause of the large Q values in NGC 6240 and Arp 220, the extreme values of the infrared excess make this scenario unlikely. The deficit of NGC 6240 and Arp 220 may well indicate that some energy source other than a starburst dominates the far infrared luminosity, f or example an active nucleus. There is independent evidence that there is an active nucleus in Arp 220; DePoy, Becklinand Geballe (1987) found that the Ba recombination line was broader (-1000 kms-1) than is usually found in HE regions, and Norris (1985) found a compact radio source. It is therefore of interest to see whether there is any systematic difference in the values of Q between starbursts and Seyferts. Table 3.8 gives the By line flux for a selection of known Seyfert galaxies. The Brackett lines in Seyfert 2 galaxies have been corrected for extinction by comparing the By flux with the Hpflux given by Veron-Cetty and Veron(1986), and assuming case B recombination. The extinction for the Seyfert 1 galaxy NGC 3227 has been taken from Ward etal. (1987). The integrated infrared flux has been plotted against the By flux for these Seyfert galaxies in Figure 3.3, along with the mergers discussed above. It is clear that, with the exception of Arp 220, the Seyferts tend to have a smaller By flux for a given infrared flux than do mergers. Table 3.8 Galaxy A2.2 Dereddened flux Ref (mag) (x l0‘17Wm-2) NGC 1672 0.1 2.4 1 NGC 1808 0.2 6.0 2 NGC 1068 0.05 25 3 NGC 7552 0.2 10 2 NGC 5643 0.18 2.8 1 NGC 7582 0.14 3.9 2 NGC 3227 0.23 0.81 3 References 1 Kawaraetal. (1988) 2 Moorwood and Oliva (1988) 3 This thesis 117 some Seyfert galaxies. The Seyfert galaxies are denoted by filled squares, and the and squares, filled by are denoted galaxies Seyfert The galaxies. Seyfert some By flux plotted against infrared flux for the mergers discussed in this sample, and sample, this in discussed mergers the for flux infrared against plotted By flux lo g (By/ W m mergers by (Tosses. by mergers iue 3.3 Figure 118 It is possible that the smaller By fluxes in Seyferts reflect the different redshift distributions for the Seyferts and mergers; the Seyferts shown in Table 3.3 have recession velocities of around 1000 km s_l, while the mergers have recession velocities of around 4000-7000 km s_ 1. Thus there maybe extended Brackett line flux outside the aperture for the Seyfert galaxies. While it is impossible to rule out the possibility that the lower By flux in Seyferts is due to selection effects, there are good reasons to suggest that these differences may be due to intrinsically different luminosity sources. The smallest aperture used on the Seyfert galaxies was 5 arcseconds, which covers a physical size of 500 pc diameter at a distance of 18-20 Mpc. This is large enough to encompass the Broad lin e Region (**1-3 pc) and the Narrow Line Region (100-200pc), which are likely to dominate the recombination line flux (see Chapter 4). The one Seyfert (NGC 1068) for which multiaperture spectroscopy is available shows that the By flux increases by at most 30% between a 5 arcsecond aperture and a 19 arcsecond aperture (see Chapter 4). Thus a small value of the ratio of the By flux to the integrated infrared flux may indicate thatthe luminosity is dominated by an active nucleus rather than a starburst To summarize, with the exception of Arp 220 and possibly NGC 6240, the Brackett line fluxes suggest that the infrared luminosity is dominated by star formation. There is other evidence, discussed in Wright et al. (1984) and Wright et al. (1988), which indicates that merging galaxies are powered by star formation. This will only be briefly recounted here; the nuclear optical emission lines are characteristic of either HII regions or shock excitation, and the optical linewidths are * 600 km s_1, similar to those found in HII regions. The radio emission is non-thermal and extended, characteristic of supemovae. Finally, Figure 3.4 shows the IRAS 60/25 vs. 100/60 colour-colour diagram for the galaxies discussed here. The regions occupied by starburst, disk and Seyfert galaxies have been outlined (Rowan-Robinson and Crawford 1985). While there is some overlap between starburst and Seyfert galaxies, the merging galaxies studied in this thesis do not have the small 60/100 ratio characteristic of Seyfert galaxies (DeGrijp et al. 1985). The far infrared colours of these mergers are more characteristic of starbursts. Figure 3.4 IRAS colour-colour diagram for the mergers discussed in this chapter. 120 3.4 Shocked Gas In Merging Galaxies The sample spectra in Figure 3.1 clearly show detections of the 2.122 pm 1-0S(1) line of molecular hydrogen and the 1.644pm [Fell] line. The aim of this section is to investigate the origin of these lines by determining their excitation mechanism and spatialextent 3.4.1 Excitation Mechanisms In Chapter 1 it was noted thatthe rotation-vibration spectrum of molecular hydrogen has been observed to be excited by two processes; thermal excitation behind a shock front and fluorescent emission following the absorption of a UV photon. Models of shock excitation indicate thatthe u=2 -1 S(l) line should have an intensity *0.1 times that of the u=l -0S(1) line. Similarly, the 1.644[FeII] line can be shock excited or photoionized. If photoionized the [Fell] line flux should be ** 1 % that of the By recombination line. Table 3.9 shows the ratios of the intensities of the o=2 -1 S( 1) line flux to the u=l - 0S(1) and the [Fell] line flux to the By line flux for the galaxies to be studied in this section. The fluxes were taken from Table 3.1. 121 Table 3.9 2- 1S(1) JFeD Galaxy Aperture 1-0S(1) By (arcseconds) NGC 6240 19.6 <0.2 >13 7.8 <0.2 >11 5.5 <0.1 >8 NGC 1614 19.6 0.1 3 12.4 - >1.8 7.8 - >1.7 Arp 220 19.6 <0.2 - 7.8 <0.2 - NGC 6052 19.6 <0.2 - 122 Thus in all these mergers both the molecular hydrogen emission and the [Fell] emission is probably dominated by shock excitation. 3.4.2 Spatial Distribution of Shocked Gas Table 3.1 shows that the fluxes in the shocked lines increase with aperture, indicating that the emission is spatially extended. This is illustrated in Figure 3 S, which shows the 1 -0S( 1) line flux plotted against aperture for Arp 220. In order to obtain a more quantitative estimate of the scale size, it is assumed that the distribution is exponential with the intensity I(r) at a radius r given by; I(r) = I o e x p ( - (3.17) Here Io is peak flux at r = 0, and ro is the exponential scale size of the distribution. The exponential scale size ro is thus the inverse of the slope of the plot of the logaritham of the surface brightness against radius. This plot is shown for the multiaperture spectra of NGC 6240, NGC 1614 and Arp 220 in Figure 3.4. The values of the surface brightness, Z, are derived from the fluxes in Table 3.1 where V - f M - f (ri) (3.18) -2 J l The flux in a within an aperture radius T2 is given by f(r2>. The abscissa in Figure 3.4 is the ’ area weighted aperture’ given by 3 3 2 £ i i l i (3.19) D3 tr2 2 " r r2 l The error bars on the points in Figure 3.4 reflect the statistical errors in the fluxes, and arrows on error bars show that the errors go off the scale of the plot. 123 Flux (10- The 1 aperture Arp 220 The for against -0S{ plotted 1)flux line —i —i —i —| r |— i— i— i— i— |— i— i— — i prue (arcsec) Aperture iue 3.5 Figure 124 t — —r i— T —i r i— i— — i Figure 3.6a Surface brightness vs area weighted aperture forNGC 6240 S(l) line (squares) and the [Fell] line (triangles) 125 Figure 3.6b Surface brightness vs area weighted aperture for Arp 220 S(l) line. 126 Figure 3.6c Surface brightness vs area weighted aperture forNGC 1614 [Fell] line. 127 Table 3.10 shows rough exponential scale sizes for the plots in Figure 3.4. No attempt was made to fit a line through the 3 or 4 points in each plot; the scale size was derived from a straight line drawn by eye through the points. All of the scale sizes derived here are in the region 1-2 kpc. It is striking that, within the errors, the scale sizes of the [Fell] emission region and 1-0S(1) line emission region in NGC 6240 are the same. The scale size of the 1 -0S(1) line emission for NGC 6240 of * 1.7kpc is identical to that derived by Lester, Harvey and C arr.(l 988). Table 3.10 Galaxy Line Scale Size ER Scale Size Reference (kpc) (kpc) (for IR scale size) NGC 1614 S(l) 2 0.2-0.5 D eP o y 1988 NGC 6240 S(l) 1.7 0.2 0 rt [Fell] 1.7 0 0 A rp 220 S (l) 2.0 <0.2 0 3.4.3 Interaction Induced Shocks? The fact that the scale size of shocked lines is similar in all three galaxies suggests a common excitation mechanism, the most plausible being shocks induced in cloud- cloud collisions due to the merging process. This can be further investigated by comparing the scale size of the shocked lines with the scale size of the infrared emission. The spatial distribution of the infrared emission in these mergers will be described in Chapter 5 (section 5.4.1) and will not be repeated here. However the exponential scale for the 10 pm emission (derived from the infrared compactness, R io - see 5.4.1) is shown in Table 3; JOJt is clear that the scale size of the shocked lines is five to ten times larger than that of the infrared continuum. Both of these observations discussed here - the similarity of the size scales of the shocked lines in the different galaxies and the extent of the infrared source 128 compared to that of the shocked gas - point to the merging process as the source of the shocks. It is therefore important to demonstrate that the kinetic energy released as a result of the merger of two disk galaxies is sufficient to power the observed shock luminosity. Harwit et al. (1987) have shown that the energy input resulting from the merger of two disk galaxies, LKin> with a combined gas mass mn, that approach each other with a velocity v* is given by: -1 , Ljan 1.5 x 10^ ergs s"1 (3.20) 2 x 10% .5 x 107an s_1. .2.7 x lO^cm. where Z is the scale height of the molecular gas. Assuming that va = 500 km-1 and Z = 90 pc, the energy available due to the interaction is *1012 L© for both NGC 6240 and Arp 220, since both of them have gas masses *3 x 1010 M©- This is an upper limit to L since the above expression assumes the galaxies collide head-on. An oblique collision will lower LKin by at most a factor of 10. This energy input is to be compared to the luminosity in the shocked gas. For NGC 6240 the luminosity in the 1-0S(1) line is * 108 L© so the luminosity in all the shocked lines will be * 109 L© (Kwan 1977). Thus the luminosity in the shocked lines is ^ 10-2 L©. Therefore merger induced shocks are an energetically favourable mechanism, however they do not dominate dissipation of the kinetic energy. This is in contrast to molecular outflows around star formation regions where * 10-90% of the kinetic energy of the outflow is radiated in H2 line emission (Fischer et al. 1985). 3.4.4 The Distribution Of Molecular Gas in Arp 220 The spatial distribution of the shocked molecular gas can be compared to the distribution of the cool molecular material, as seen in CO. Both NGC 6240 (Sargent etal. 1987) and Arp220(Scovilleetal. 1986) .have been mapped with the Owens Valley interferometer. The cool molecular gas in NGC 6240 does not show any clumping on scales of 4 arcsecond, and yet it is one of the most gas - rich galaxies known. This is qualitatively consistent with our result that the shocked molecular gas is extended. The aperture synthesis maps of Arp 220 suggest that a large fraction (and perhaps all) of the molecular material is centrally concentrated. In contrast, single dish mapping ofthe J = 1-0 CO line by Cassoli et al. (1988) shows the molecular gas to be extended along the dust lane. These observations of the cold molecular material will be compared to the spatial distribution of the shocked gas here. 129 Scoville et al. (1986) argue that at least 70%, and perhaps all, of the J=l-0 CO line flux in Arp 220 comes from an unresolved 4x6 arcseconds source. At first sight, the result that the 1 -0S( 1) line flux doubles between the smallest and largest apertures (section 3.4.3) appears to be inconsistent with this. However, an estimate of the fraction of shocked molecular gas that is within the 5 arcsecond aperture is dependent on the differential extinction between the 5 and 19 arcsecond apertures. In order for » 70% of the flux to be concentrated within the smallest aperture, the intrinsic (i.e. unobscured) nuclear flux must be * 6.25 , implying a differential extinction at 2pm of - 1 mag, or Av * 10 mag. Therefore a differential extinction of -10 magnitudes in the visual between the regions covered by the the 5 and the 19.6 arcsecond apertures is required to make this estimate of the concentration of the molecular material consistent with that of Scoville et al. In order to obtain a crude estimate of the differential extinction between 5 and 19.6 arcseconds, multiaperture (5,7.8 and 19.6 arcseconds) JHK photometry was undertaken atUKIRTin May 1987. The near infrared colours and the average extinction in each aperture are shown in Table 3.11. To estimate the differential extinction it is assumed that the underlying stellar population has intrinsic colours J- H=0.7 and H-K=0.3, and that the reddening in Arp 220 follows van de Hulst curve No 15 (in which Ax - X-1*9, where Ax is extinction in magnitudes at a wavelength X). The extinction at 2.2pm in the 5 arcsecond aperture is 0.54 magnitudes, consistent with the estimate of Norris (1985) of 0.4 mags. The differential extinction between the smallest and largest apertures is 0.39 mag at 2.2pm. Uncertainties in the reddening law affect this estimate of the extinction; a X-1 curve will increase the differential extinction between the 5 and 19 arcs apertures to - 10 mags in the visual. However, if there is an active nucleus it will contribute more to the flux at K than J (assuming a F0- power law, with a - 1-. 15). Therefore Av - 10 is an upper limit to the differential extinction. 130 Table 3.11 Aperture J-H H-K 5.0 1.26 0.88 0.54 7.8 1.06 0.74 0.32 19.6 0.87 0.51 0.15 Adopting Av “ 10 as an upper limit to the differential extinction, the 5 arcsecond aperture flux increases by a factor of 2.5, and so the fr a c tio n of shocked molecular gas that is within 5 arcseconds is consistent with the CO measurements of Scoville et. al., however, these results rule out the tentative suggestion made by Scoville et al. that a ll of the molecular material is concentrated within 4x6 arcseconds, and indicate that 70% is an upper lim it This is qualitatively consistent with the results of Cassoli et al. (1988). 3.5 Summary and Conclusions This chapter has focussed on using infrared spectroscopy to investigate the physical processes in merging galaxies. A simple starburst model was used to show that the Brackett line fluxes and the infrared continuum fluxes are consistent with the bulk of the luminosity being produced in a Zero Age Main Sequence starburst, or an evolved starburst with a lower mass cut off « 5 M o Multiaperture spectroscopy of the 1 - 0S(1) line and the [Fell] line for three merging galaxies (NGC 6240, NGC 1614 and Arp 220) was presented. The lines were found to be shock excited, with exponential scale sizes of « 2000 pc. This is larger than the scale size for the infrared emission, which is typically 200 pc. This suggests that the luminous shocked emission observed from these galaxies is not directly associated with the source of the infrared luminosity. The most likely alternative is that the shocks are induced in cloud-cloud collisions due to the merging process. Finally, the spatial distribution of the 1-0S(1) line of molecular hydrogen was compared to the distribution of molecular gas as traced out by the J= 1 -0 line of CO in Arp 220. The extended H2 emission clearly rules out the suggestion made by Scoville et al. (1986) that all of the molecular gas is concentrated within 4 arcseconds of the nucleus in this galaxy. 131 Chapter 4 Star Formation in Arp 299 - a Case Study 4.1 Arp 299 rp 299 (alternatively Markarian 171) is a system comprising two A distorted, interacting spiral galaxies, NGC 3690 and IC 694. This system is one of the most luminous interacting pairs known, with a bolometric luminosity of » 1012 L o (assuming that Ho=50 kms-l Mpcrl). There is evidence from infrared, radio and optical observations that large scale star formation is occurring in Arp 299, as well as evidence for an active nucleus (Gehrz, Sramek and Weedman 1983). Arp 299 is closer (60 Mpc) than many of the ultraluminous merging galaxies discussed in the previous chapter, making it an ideal candidate for simple line mapping. Interest focused on Arp 299 after the publication of a paper by Gehrz, Sramek and Weedman (1983), who presented optical spectra as well as infrared and radio maps of this system. The radio and infrared maps showed one peak associated with the nucleus of IC 694 (named source A), and two peaks in NGC 3690. Of the latter two peaks one is associated with the nucleus of NGC 3690 and is known as source B, and the second (source C) is situated » 10 arcseconds north of B and is spatially coincident with an extra-nuclear HD region. Gehrz, Sramek and Weedman were able to show, using simple starburst models, that the radio, infrared and optical observations they presented for sources B and C were consistent with star formation. The major results of their starburst modelling were that there are a relative deficiency of low mass stars and that several supemovae per year are required to produce the observed radio flux. In contrast to this, source A has a strong, compact, flat spectrum radio source that cannot be explained in terms of star form ation. 132 Arp 299 has been studied very intensively since the publication of thispaper. Telesco, Decher and Gatley (1985) obtained near infrared maps and multicolour near infrared photometry of this system, and found that the strongest sources were identifiable with the nuclei of IC 694 and NGC 3690 (sources A and B in the notation of Gehrz, Sramek and Weedman ). Sargent et al. (1986) used the Owens Valley Millimeter Wave Interferometer to map Arp 299 in the 1=1-0 line of CO, in order to determine the distribution of molecular gas. They found regions of compact molecular gas situated at components A and C, but found no evidence for a source at B. More recently, Joy et al. (1989) obtained scans across Arp 299 at 50 and 100pm in order to determine the relative contribution of these two sources to the far infrared luminosity. They found that * 60% of the 50 and 100pm emission was attributable to source A. Near infrared spectra of Arp 299 have been obtained by Fischer et al. (1983), who detected the S( 1) line and the By line in an 8 arcsecond aperture centred on A and a 21 arcsecond aperture encompassing sources B and C. These spectra showed that Arp 299 is a luminous (1.5xl07 Lo) source of vibrationally excited molecular hydrogen, and that both the molecular hydrogen emission and the recombination line emission were extended on scales of kiloparsecs. Detections of the Ba and By lines have been reported from sources A and B by DePoy (1987). Nakagawa et al. (1989) obtained detections of the By line in a 3 arcsecond aperture from sources A, B and C, and a detection of the S(l) line from source A. These observations show that there is a vast body of data on Arp 299. This, coupled with its proximity, makes Arp 299 an ideal candidate for studying the detailed distribution of infrared lines, and relating them to star formation activity. 4.2 The Data The data presented in this thesis consist of 19 and 12 arcsecond apertures centred on source A, a 19 arcsecond aperture centred on sources B and C, and an 8 arcsecond aperture on source B alone. These apertures are shown superimposed on the 10pm map of Gehrz, Sramek and Weedman in Figure 4.1. Figure 4.2 shows the spectra,with the redshifted positions 133 of the major infrared lines indicated. Table 4.1 shows the measured fluxes of the S(l) and By lines, along with the fluxes of these lines taken from the literature. Table 4.1 Position Flux Aperture Reference xlO-17 Wm-2 (arcseconds) S(l) By IC 694 7.3±0.5 9.7±0.3 19.6 1 (A) 6.1±0.5 5.8±0.9 12.4 1 6. 1± 1.6 8 2 4.4±0.4 5 3 4.5±1.2 5 4 3.5±0.7 1.9±0.7 3 4 NGC 3690 3.9±0.5 5.4±0.6 7.8 1 (B) 2.6±0.4 5 2 <3.1 5 4 < 1.8 <1.7 3 4 NGC 3690 4.8±0.3 7.9±0.5 19.6 1 (B+C) 6.4±1.3 10.1±1.3 21 2 NGC 3690 <1.8 4.1±1.2 3 4 (C alone) <2.4 5 4 References 1 This thesis 2 Fischer etal. (1983) 3 DePoy (1987) 4 Nakagawaetal.(1989) 134 ♦ W SO'OO * Figure 4.1 10 pm map of Arp 299, reproduced from Gehrz, Sramek and Weedman (1983). The drcles show the apertures used in this thesis. 135 Flux (10 19.6 arcsecond K w indow spectrum o f IC 694 (source A). (source 694 IC f o spectrum indow w K arcsecond 19.6 iue 4.2a Figure T vlnt (yum) avelength W T T 136 T T Flux (10 19.6 arcsecond K w indow spectrum o f NGC 3690 (sources B plus C). plus B (sources 3690 NGC f o spectrum indow w K arcsecond 19.6 iue 4.2b Figure vlnt (,um) avelength W 137 Flux (10 12 4 arcsecond K w indow spectrum o f IC 694 (source A). (source 694 IC f o spectrum indow w K arcsecond 4 12 iue 4.2c Figure vlnt (i ) (/im avelength W 138 Flux (10 7.8 arcsecond K w indow spectrum o f NGC 3690 (source B). (source 3690 NGC f o spectrum indow w K arcsecond 7.8 vlnt (^im) avelength W iue 4.2d Figure 139 The line fluxes presented here are consistent with the detections and upper limits in the literature. An estimate of the By flux at position C can be made from the data in this thesis by subtracting the 8 arcsecond flux at B from the flux in the 19 arcsecond aperture centred on Band C. This implies a By line flux at C of 2.5±0.8 Wm-2, compared to 4.1±1.2 Wm-2 measured by Nakagawaetal. (1989). These two estimates are formally consistent, and suggest that the source of the Brackett line emission at position C is concentrated within the smaller (3 arcsecond)aperture. 4.3 Excitation Mechanisms The vibrational spectrum of molecular hydrogen can be excited both by fluorescence and shocks, as outlined in Chapter 1. The ratio of the flux in the 2- 1S(1) line to the flux in the 1-0S(1) can be used to distinguish between pure fluorescence and thermal excitation since the 2-1 S( 1) line flux is enhanced when the gas is fluorescently excited. Measurements of the flux in the 2- 1S(1) line were only obtained for the 12 arcsecond aperture centred on source A and the 8 arcsecond aperture centred onB. The ratio of the flux in the 2- 1S(1) line to the flux in the 1-0S(1) is <0.15 at source A and <0.35 at source B (see Table 4.2) These values are more consistent with that expected for shock excitation («0.1) than pure fluorescence («0.5). While the relative intensities of the 2-lS(l) and the 1-0S(1) are consistent with thermal excitation, the 1-0S(0) line is enhanced relative to the l-O S(l)lineinthe 12 arcsecond spectrum centred on A. There is no obvious reason for this, and it is not seen in any other apertures. 140 Table 4.2 2-lS(l) position aperture by S(l) 1-0S(1) (arcseconds) source A 19.6 :)±0*Z - 12.4 Q - q s ± 0 ’ib <0.15 5 0.97±0.27 2.7 0.54±0.22 source B 7.8 1-3$ ± 0-2 <0.35 source B and C 19.6 {■oh -i r ;V source C 2.7 *1.5 If the molecular hydrogen emission is dominated by fluorescence it is important to demonstrate that there are sufficient ionizing photons to produce the observed S(l) line flux. Models of fluorescent emission indicate that « 50 UV photons are absorbed for each S(l) line photon emitted, as discussed in Chapter 1. The number of ionizing photons is « 70 times the number of By photons, assuming case B recombination. Therefore for fluorescent excitation to be viable the ratio /Fsi >1, where F si is the flux in the 1-0S(1) line, and Fgy if the By line flux. This ratio is shown in table 4.2 for all apertures and positions for which S(l) line fluxes and By line fluxes are available, j^j /$ ' kaslbt*- Cct Scp-i^ u t) j 141 It should be noted that it is impossible to rule out high density fluorescence on die basis of the data available. However, all the data are consistent with shock excitation, and the luminous [Fell] emission independently shows that large scale shocks are present I therefore conclude that the S(l) line flux is dominated by shock excitation, subject to confirmation by further observations (see chapter 7). 142 4.4 Extinction In order quantitatively to compare the By flux to the radio and infrared fluxes, it is necessary to determine the extinction correction at 2pm. Table 4.3 shows estimates of the visual extinction (taken from the literature) at positions A, B and C in afive arcsecond aperture. Table 4.3 Source Method Av Ref (M ag) Source A N ear IR 4.6 1 colours B a/B y 4.9 4 silicate absorption 24 2 Nco 130 3 Source B Bot/By 10.2 4 Ho/Hp 1.8 2 silicate absorption 14.0 2 Nco <28 3 Source C silicate absorption 5 2 Ho/HP 1.5 2 Nco 130 3 References 1 Telesco, Decher and Gatley (1985) 2 Gehrz SramekandWeedman(1983) 3 Sargent etal.( 1987) 4 DePoy(1987) 143 Table 4.3 shows the results of four different methods to estimate the extinction. These results are wildly discrepant in places; for example the value of the visual extinction at source A varies by two orders of magnitude. There is a clear trend for the value of the extinction to increase with the wavelength at which the measurement was taken. This is primarily due to two effects; longer wavelength measurements probe further into the obscured gas, and the extinction to the emitting region is patchy. The extinction to the line emitting gas is probably best estimated from the ratio of the infrared recombination lines. Measurements at longer wavelengths preferentially probe dense molecular cloud cores, and thus over-estimate the extinction to the emitting line gas, whereas the opacity of the optical recombination lines is probably greater than unity. The values adopted for the extinction at sources A and Bare A v =4.9 and Av= 10.2 respectively, estimated from the B a to By ratio and assuming case B recombination. The extinction at C is derived from the silicate absorption feature, and is given by Av = 5. The By recombination lines were corrected for reddening by assuming that Agy = 0.1 Av. 4.5 Star Formation in NGC 3690 The optical, infrared and radio properties of NGC 3690 all suggest that the infrared luminosity is powered by star formation. The Brackett emission then arises in star formation regions, and it should be possible to relate the By flux in a self consistent way to the infrared and radio properties of this galaxy. The aim of this section is compare the infrared recombination line fluxes at the nuclear source B and source C with the thermal radio flux at these positions, and to use the simple starburst models described in Chapter 3 to describe large scale star formation in this system. 4.5.1 Thermal Radio Emission. The thermal flux at centimetre wavelengths for the sources B and C in a 5 arcsecond aperture can be estimated from the 5 cm and 20 cm radio maps presented by Gehrz, Sramek and Weedman (1983), and can be compared to that inferred from the extinction corrected By flux in a 5 arcsecond aperture given in section 4.3. The radio emission is due to a mixture of 144 thermal emission from HE regions and non-thermal emission form supernova remnants. No-thermal emission from supernova remnants has a characteristic ir 1 power law spectrum; it is therefore assumed that this power-law component dominates the 20 cm flux. The flux at 6 cm due to supernova remnants can be calculated from the power law and subtracted from the observed 6 cm flux to give the thermal contribution. The thermal 6 cm fluxes estimated using these assumptions are 5.6 mJy for source B and 2.0 mJy for source C. The centimetre free-free flux calculated on the basis of the By line flux and assuming case B recombination is given by (see Chapter 1): The extinction-corrected By fluxes given in section 4.3 imply radio thermal fluxes of 4.5 mJy for source B and 6.5 mJy at source C. These values are consistent with the 6 cm thermal radio fluxes inferred from the radio maps of Gehrz, Sramek and Weedman (1983). This agreement is good evidence that the sources at B and C are powered by star formation. 4.5.2 Infrared Emission The results from the previous section demonstrate that the By and radio fluxes are consistent with star formation at sources B and C. The starburst models discussed in Chapter 3 can be used to investigate further the characteristics of the star formation in NGC 3690, by comparing the star formation rate inferred from the infrared flux with that determined from the extinction corrected By flux. The infrared flux from NGC 3690 has been estimated by Joy et al. (1989). The components B and C have a small angular separation and are not resolved by Joy et al., so it is not possible to determine which component dominates the far infrared luminosity from NGC 3690. Therefore star formation in the NGC 3690 system as a whole will be investigated in this section. 145 The integrated infrared flux measured by ERAS for the NGC 3690/IC 694 system is 4.9 x 10"12 Wm'2. Joy et al. (1989) have estimated that« 40 % of the far infrared flux is associated with NGC 3690, so it will be assumed that the integrated infrared flux from B and C is 1.9 x 10"12 Wm-2. The By flux from a 19 arcsecond aperture encompassing B and C is 7.9 x 10-17 Wnr2 . The extinction in the larger aperture is unknown, however the extinction at sources B and C is in the region A^ 0.5 -1.0 mag, and a value of 0.5 will be adopted. This may overestimate the extinction in the larger aperture, since B and C are bright infrared sources and therefore are likely to be dusty regions of the galaxy. The extinction corrected By flux of 12.5 x IO-I7 Wm’2 is therefore an upper limit Figure 4.3 is a reproduction of Figure 3.2, which shows the By flux plotted against the ERAS infrared flux for the mergers discussed in Chapter 3. The By flux and the infrared flux derived above for NGC 3690 are have been plotted for comparison. The ratio of these fluxes is strongly suggestive of a ZAMS starburst, or an evolved starburst with a high (** 5M ) lower mass cut off. These are exactly the characteristics associated with the star formation in the mergers described in Chapter 3, and indicates that the same physical processes underlie the star formation in NGC 3690. 146 log (By/ W By flux versus infrared flux for NGC 3690, with the mergers discussed mergers the with 3690, NGC for flux infrared versus flux By in Chapter 3 plotted for comparison. 3 for Chapter plotted in iue 4.3 Figure 147 4.6 Star Formation in IC 694 The results of the previous section show that the infrared, radio and recombination line fluxes from NGC 3690 can be understood in terms of a burst of high mass star formation. There is some evidence, outlined in the introduction, that IC 694 contains an active nucleus. The purpose of this section is to determine what fraction of the infrared luminosity from source A can be attributed to star formation, by examining the radio and recombination line fluxes and the spatial extent of the By emission. 4.6.1 Thermal Radio Flaxes The radio flux at source A in a 5 arcsecond aperture can be estimated from the radio maps of Gehrz, Sramek and Weedman (1983) to be « 132 mJy. The radio emission from source A has a flat spectral index with very little sign of a steep spectrum power law component; therefore if the radio flux at A is associated with a starburst then it arises in HU regions. The extinction - corrected By line flux in a 5 arcsecond aperture is « 6.9x10' 17 Wm-2, from which it is possible to infer a centimetre radio flux of 4.7 mJy. This estimate of the thermal radio flux is completely inconsistent with the measured value of 132 mJy, and indicates that only a small fraction of the radio flux is associated with star formation. The radio flux is almost certainly dominated by an active nucleus. 4.6.2 Spatial Extent of the By Emission Itis clear from Table 4.1 that the Byfluxin a 19 arcsecond aperture is approximately double thatin a three arcsecond aperture, indicating that the emission is spatially extended on scales of 5 kpc (assuming that Ho = 50 km s'lMpcrl, and cz=3100 km s_1). The scale size of the emitting region can be determined approximately by assuming that the spatial distribution is exponential, as outlined in section 3.4.2. A plot of the logarithm of the surface brightness versus the area weighted aperture is shown in Figure 4.4, from which a scale size of approximately 1.5 kpc can be derived. 148 (arcsec) Figure 4.4 Surface brightness plotted against area weighted aperture for IC 694 149 The By emission may be associated with an extended starburst region, or it may originate in gas photoionized by the active nucleus. In many active nuclei gas is seen photoionized by the active nucleus several hundred parsecs outside the nuclear region; a classic example of this is the Seyfert galaxy NGC 1068 where "broad’ recombination lines ( Av « 1000 km s-1 ) are seen * 200 - 500 pc from the nucleus (see Chapter 5). The Brackett line emission in IC 694 is extended on scales * 1 kpc, strongly suggestive of an extended starburst region surrounding the active nucleus. A lower limit to the By flux associated with the starburst can be obtained by subtracting the small (5 arcsecond) measurement of the By fluxfromDePoy (1987) from the 19 arcsecond aperture measurement presented in this thesis; this extended flux is * 5.3 x 10-17 Wm*2. The 5 arcsecond aperture encompasses a diameter of ” 1.5 kpc, and thus emission outside this aperture is almost certainly due to star formation. The star formation in IC 694 is likely to be characterised by a ZAMS starburst similar to that seen in the other interacting and merging galaxies discussed in this thesis. Under these circumstances Q ”3-6x1 O'3,where Q is the ratio of the star formation rate derived from the infrared luminosity to that inferred from the By luminosity, defined in section 3.3.1. This value of Q implies that the integrated infrared flux associated with the starburst is »0.8-1.8 x 10_12 Wnr2, which is 30- 50% of the observed integrated infrared flux. It should be emphasised that this estimate is a lower limit since no correction to the extended By flux has been made for extinction. In conclusion, assuming that the star formation in IC 694 is similar to the mergers discussed in the previous chapter, the extended starburst contributes at least 30% of the observed infrared luminosity at A. This situation is similar to that seen NGC 1068, where a drcumnuclear starburst contributes 50% of theintegrated infrared luminosity, and is extended on scales ^ lkpc. It also supports the conclusion of Joy et al.( 1989) that the active nudeus does not dominate the far infrared emission since the spectrum of the far infrared emission is cooler than is normally seen in active nuclei. 150 4.7 Spatial Distribution of Shocked Gas. Arp 299 is an exceptionally luminous source of vibrationally excited molecular hydrogen. The luminosity in the S(l) line is « 1.3 x 107 Lo (compared to * 2.5 L o from the Klemmann - Low nebula in Orion), and the emission is extended over several kiloparsecs in both NGC 3690 and IC 694. The S(l) line emission is most likely to be associated with the starburst activity in Arp 299, powered by shocks in molecular clouds generated by outflows from young stars and supernova remnants, as described in Chapter 1. In this section the spatial distribution of the shocked gas is compared to the cold molecular gas traced out by mapping the J= 1 -0 CO line, and the By line. 4.7.1 Comparison with CO Distribution The aperture synthesis maps of Sargent etal. (1987) show that the CO emission from the Arp 299 system is concentrated in two compact masses, one coincident with Source A in IC 694, and one coincident with source C in NGC 3690. Sargent et al. do not report any compact molecular gas at source B, the nucleus of NGC 3690. Table 4.1 shows that the S( 1) line flux measured in a 19 arcsecond aperture centred on sources B and C is 4.8 ± 0.3 compared to 3.9 ± 0.5 in an 8 arcsecond aperture centred on B alone. These results suggest that essentially all of the shocked molecular gas is centred at B, in marked contrast to the results of Sargent et al. The most obvious reason for this discrepancy is that the interferometer observations of Sargent et al. are not sensitive to low surface brightness extended CO emission. This conclusion is reinforced by the fact that only 40% of the single-dish CO flux was recovered in the interferometer measurements, implying that most of the molecular gas is not associated with the high density clumps observed by Sargent etal... 151 There is growing awareness that estimates of the mass of molecular gas in a galaxy from CXD observation can lead to misleading results. The mass of molecular gas derived from CXD is dependent on the ratio of the integrated CXD intensity (Ico) to the molecular hydrogen column density (Young and Scoville 1982). This ratio has been measured in many galactic molecular douds, and it is assumed to be the same in molecular douds in other galaxies. This assumption is on ly valid where the physical conditions (such as temperature, density and metallicity) in molecular clouds are similar to those found in galactic molecular clouds (Israel 1982). For example, the mass of the compact molecular source at C may be overestimated if the temperature of the douds is significantly higher than galactic sources. The discrepancy between the S(l) line and CXD distributions in both Arp 299 and Arp 220 (described in section 3.4.4) further highlights the limitations of using observations of CXD to characterise extragalactic molecular gas. 4.7.2 S(l) Line Distribution at Source A Inspection of the data in Table 4.1 shows that the flux in the By line has clearly increased in going from the 12 arcsecond to the 19 arcsecond aperture, whereas the S(l) line flux shows no such increase. This suggests that the S( 1) line is less spatially extended than the By line. This is further demonstrated in Table 4.4 which shows the ratio of the flux in the By line to that in the S (l) line as a function of aperture. Thisratiois * 1 in the 19 arcsecond aperture, whereas it is « 0.5 in the smallest aperture. It should be emphasised that this is only a 2o result, and further measurements are needed to confirm it This result if confirmed, is in contrast to those of Chapter 3, where it was found that the shocked gas is more extended than the ionized gas and infrared emission in mergers. It is possible that the small spatial extent of the S(l) line relative to the By line is due to a nuclear component of the S(l) line flux associated with the active nucleus. 152 Table 4.4 Aperture By/S(1) (Arcseconds) 19.6 1.05±0.2 12.4 0.75±0.05 5 0.97±0.27 2.7 0.54±0.22 4.8 Summary and Conclusions The main conclusion of this study is that the radio and infrared line and continuum data show that the bolometric luminosity in Arp 299 is dominated by star formation. The small aperture radio and By fluxes centred on components B and C are characteristic of HII regions, and the ratio of the By flux to the integrated infrared flux from NGC 3690 is consistent with ZAMS/high mass evolved star formation. The only evidence for an active nucleus comes from the flat spectrum radio source at source A, where the free-free radio flux predicted from the By flux is a factor of 25 less than the observed free-free radio flux. There is, however, an extended region of By emission around source A , which contributes at least 30% of the total infrared luminosity from IC 694. The spatial distribution of the shocked gas (traced out by the 1 - 0S(1) line) has two unexpected features. The first is that the S(l) line in IC 694 may be less extended than the By line. This is in contrast to the merging galaxies discussed in the previous chapter, where the shocked gas had a size scale 5-10 times that of the infrared emission. The second is that the S(l) line in NGC 3690 is concentrated at component B, whereas the J= 1-0 line of CO is concentrated at C (Sargent et al. 1987). This discrepancy is probably due to the fact that the interferometer used by Sargent et al. is not sensitive to low surface brightness emission, and suggests that CO interferometer maps cannot be used reliably to infer the spatial distribution of molecular gas. 153 Chapter 5 Near Infrared Spectroscopy of Seyfert Galaxies 5.1 Introduction PTICAL astronomers have known since the second decade of this century that the nuclear regions of some ’spiral nebulae’ have bright, broad emission lines (Slipher 1917), even though at the time it was not generally understood that spiral nebulae were external galaxies. These galaxies are named Seyfert galaxies after the astronomer Carl Seyfert who identified a number of ’emission line nebulae’ (Seyfert 1943). A brief summary of the main properties of Seyferts will be given in this section. Seyfert galaxies are characterised by a bright, point-like nucleus, which often has a luminosity comparable to the whole of the rest of the galaxy. These nuclei have broad (» 103 km s_1) emission lines, which are considerably wider than the stellar velocity dispersion (** 300 km s_1) (Weedman 1977). A few percent of all galaxies have Seyfert characteristics, and all Seyferts are spiral galaxies. Some elliptical galaxies do display characteristics of Seyferts (bright nuclei and broad lines) but this is usually accompanied by strong radio emission, and these ellipticals are not classified as Seyferts. Finally, the continuum and line fluxes of some Seyfert galaxies exhibit remarkable time variability (from minutes to years), and this property has been used to measure the size scale of the emitting regions (Peterson 1988). These observations indicate that there is some ’central engine’ responsible for producing the ionising flux required to illuminate the surrounding material and so produce the emission lines. Seyferts are traditionally divided into two categories based on the width of the emission lines (Osterbrock 1977). Seyfert 1 galaxies have very broad wings (FWHM « 5 x 103 km s*1) on the Balmer lines and other permitted lines, and also have narrow (*s 103kms -1) f orbidden lines. Seyfert 2 galaxies typically have permitted lines lines with the same width as the forbidden lines. Seyfert galaxies can also be categorised on the basis of the ratio between the permitted lines and the forbidden lines. Seyfert 1 galaxies have Hp/[ODI]« 1 and Seyfert 2 galaxies have HP/[OIII]*0.1 (Weedman 1977, Veron-Cetty and Veron 1986). The usual interpretation of these observations is that the nuclear regions of these galaxies can be roughly divided into two sections; the wide wings originate in the Broad Line Region 153 (BLR), whichis absent or hidden in Seyfert2’s, and the narrow forbidden lines and the cores of the permitted lines originate in the Narrow Line Region (NLR). The material in the BLR is thought to be photoionised by the central source. There is evidence for a wide range of densities in the BLR, but the prevalence of high excitation lines such as CIV indicates that the radiation field is harder than is seen in star formation regions, as does the presence of hard (> 5 keV) X-ray sources in many AGN. The spatial extent of die BLR is too small to be measured directly, but the approximate size of the region producing the continuum emission has been estimated from variability studies to be a few light days. The size of the emitting region is wavelength dependent, however, and can be as small as light minutes minutes in X-rays (e.g. Barr and Mushotsky 1986, Pounds and Turner 1987). The very short timescale for variability in hard X-rays is generally taken to mean that hard X-rays probe the inner portion of the BLR. It is thought that the Broad Line clouds are opaque to soft X-rays, so variability here may be associated with Broad Line clouds moving across the line of sight to the central source (Wachter, Strauss and Filippenko 1988). Finally, the size of the BLR scales with the luminosity of the central source and in quasars can be up to a light year across (Osterbrock 1985). The density of the material in the BLR has been estimated to be «109-1010 cm-3, and the temperature in the region of 104 K (Kwan and Krolik 1981, Krolik, McKee and Tarter 1981). These are rough order of magnitude estimates only; there is considerable evidence that there is a wide range of densities in the BLR. This is apparent from variability studies which show that changes in the broad line fluxes lag behind continuum changes, suggesting that the line emitting gas is situated further from the central source than the material producing the continuum emission (e.g. Osterbrock 1985). Finally, the broad lines are frequently (but not always) asymmetric, allowing some inferences about the dynamics to be made. The observed asymmetries in the broad lines can be modelled either by asymmetric expansion, or rotation under the gravitational field of a central blackhole (Osterbrock 1989). Itis not possible to differentiate between these two models with current observations (Osterbrock 1989) The NLR encompasses a much larger volume than the BLR, extending out to 10 -1 OOpc. The electron density in this regioni v * 102- 106 cm-3 is sufficiently low that collisional de-excitation is rare, hence the forbidden line emission. The NLR are large enough for the spatial extent and other properties (such as ionization state, density and velocity field) to be mapped directly in nearby Seyferts ( e.g. Heckman et 154 al. 1981, Baldwin, Wilson and Whittle 1987). Once again, a picture is emerging where there is considerable spatial stratification of these properties. For example, the dynamics of the gas in the outer parts of the NLR is probably dominated by orbital motion around the centre of the galaxy (as evidenced by the correlation between the stellar velocity dispersion and the width of the Narrow Lines (Meurs and Wilson 1984)) whereas the gas in the inner part has a velocity component suggestive of outflow (Osterbrock 1989). Radio maps of Seyferts frequently show elongated features, not unlike the much larger jets found in radio galaxies (e.g Stone, Wilson and Ward, 1988, Neff and Ulvestad 1988). It is becoming apparent that the interaction of the radio emitting plasma with the NLR gas may play an important part in the dynamics of this region (Wilson and Heckman 1985) From these characteristics a ’standard model’ of Seyfert activity has been formulated. This model (Lynden-Bell 1969, Rees 1984) suggests that the primary luminosity source is gravitational energy; a central massive (107-109 Mo) black hole is accreting surrounding material. The gravitational potential energy released as material falls onto the black hole is predominantly in the form of X-rays and y - rays. This radiation is responsible for ionising material in the Broad Line Region. Seyfert 1 and Seyfert 2 galaxies differ in the amount of dense broad line gas near the nucleus - in Seyfert 2’s the covering factor of the broad line clouds is very small or the broad line component is nonexistent, hence photons can ionize the circumnuclear gas producing the NLR. A schematic diagram of a 'slice' through an active nucleus (adapted from Rees 1986) is shown in Figure 5.1. This model also explains the similarity between Seyfert galaxies and quasars, which display extreme Seyfert 1 characteristics. However, the exact relationship between Seyferts, quasars and other forms of Active Galactic Nuclei is still a matter for debate. 155 - 10“ c;r, A 10locm A 10,9cm A 10,e cm A 1017 cm A 10,6cm 1015 cm Figure 5.1 Schematic diagram of an active nucleus, taken from Rees (1986). 156 5.2 Near Infrared Spectroscopy of NGC 1068 NGC 1068 is a classic Seyfert 2 galaxy. It is one of the closest known Seyferts (18 Mpc), and therefore has been intensively studied at all wavelengths. In this section, multiaperture H and K band spectroscopy of NGC 1068 is presented. The aim is to discover the origin and excitation mechanism of the By recombination line, [Fell] line and the rotation vibration spectrum of molecular hydrogen, and their relation to the active nucleus. This is achieved by comparing the relative extents of these lines with other line and continuum maps. 5.2.1 NGC 1068 -Seyfert and Starbnrst. There is strong evidence that the Seyfert nucleus at the centre of NGC 1068 is surrounded by an extended starburst region. Optical photographs showthatthe circumnuclear region is more luminous than most Sb spirals, and the UBV colours are characteristic of B-stars (Smith, Weedman and Spinrad 1972). Radio maps show that collimated radio emission extends several arcseconds from the nucleus, reminiscent of the radio jets seen in other active nuclei. (Wilson and Ulvestad 1983, Pedlar Dyson and Unger 1985, Ulvestad, Neff and Wilson 1987). Apart from this jet, which is associated with the active nucleus, there is also evidence for a more extended plateau of radio emission interpreted by Wynn-Williams, Becklin and Scoville (1985) as supernova emission. Finally, scans across the nucleus at 10pm show that only half of the infrared luminosity can be attributed to the active nucleus; the other half comes from a more extended (about lOkpc) circumnuclear region (Telesco, Becklin and Wynn-Williams 1984). More recently, high resolution mapping at all wavelengths has helped to clarify the morphology of the starburst region. Aperture synthesis maps in the 1 -0 CO line by Myers and Scoville (1987) show that the molecular hydrogen is concentrated in an incomplete ring around the nucleus, with a radius of approximately 10 arcseconds. There is also evidence for some CO within 4 arcseconds of the active nucleus. The infrared map by Telesco and Decher(1988) shows that the 10pm emission is concentrated in the vicinity of the molecular gas. These two maps are shown in Figure 5.2. 157 Figure 5.2 The 10 Jim emission (hatched region) mapped by Telesco and Decher (1988), superimposed on the CO contours of Myers and Scoville (1987). Taken from Telesco andDecher(1988). There have been two papers presenting near infrared spectra of the nuclear regions of NGC 1068, the first by Thompson, Lebofsky and Rieke (1978) in an 8 arcsecond aperture, and the second by Hall etal. (1981) in a 3.8 arcsecond aperture. Both these spectra show clear detections of the By line and the 1 -0S(1) line, and both conclude that the molecular hydrogen line emission is thermally excited. There is some disagreement over the origin of the nuclear molecular hydrogen line emission, however. Thompson, Lebofsky and Rieke (1978) suggest the 2H lines are formed in conditions similar to those found in planetary nebulae. Hall etal. reject this interpretation on the grounds that the ratio of the By to 1-0S(1) line flux is large in planetary nebulae, and suggest that the near infrared lines arise in molecular cloud complexes. 5.2.2 The Data The data for this study consist of H (1.6-1.74pm) band spectra in 7.8 and 12.3 arcsecond apertures and K (2.05-2.4pm) band spectra in 7.8,12.3 and 19.6. All apertures were centered on the nucleus. Figure 5.2 shows these apertures relative to the CO ring and 10pm emission. The largest aperture touches on the CO ring. The 158 data reduction was performed as described in Chapter 2. The reduced spectra, with baselines subtracted, are shown in Figure 5.3. The measured line fluxes are shown in Table 5.1, along with previous measurements of the By and 1-0S(1) line fluxes. Table 5.1 Aperture Line Flux Reference (Arcseconds) (xlO-16Wm-2) 19.6 1-0S(1) 3.0±0.1 1 m By 2.77±0.05 1 12.4 1-0S(1) 2.3±0.2 1 * By 2.5±0.2 1 7.8 1-0S(1) 2.2±0.4 2 0 By 1.8±0.4 2 7.8 1-0S(1) 1.89±0.10 1 0 By 1.85±0.10 1 3.8 1-0S(1) 1.65±0.15 3 0 By 1.85±0.15 3 References 1 This thesis 2 Thompson, Lebofsky and Rieke (1978) 3 Hall etal. (1981) 159 Flux (10"‘5 Wm 96aceodKwno pcrmo G 1068. NGC of spectrum K-window arcsecond 19.6 vlnt (/ avelength W iue 5.3a Figure 160 2 m) Flux (10 24 rscn -idwsetu fNC 1068. NGC of spectrum K-window arcsecond 12.4 vlnt (yu.m) avelength W iue 5.3b Figure 161 Flux (10 . aceodK- no setu ofNC 1068. NGC f o spectrum indow -w K arcsecond 7.8 vlnt (yum) avelength W iue 5.3c Figure 162 Flux (10 24 rscn widw pcrm fNC 1068. NGC of spectrum indow -w H arcsecond 12.4 vlnt (/xm) avelength W iue 5.3d Figure 163 Flux (10 . aceodH- no setu ofNC 1068. NGC f o spectrum indow -w H arcsecond 7.8 vlnt (/zm) avelength W iue 5.3e Figure 164 5.2.3 Comparison with Previous Work The main purpose of the observations reported in this section was to determine the spatial extent of the infrared lines; they do not resolve the disagreement over the origin of the infrared lines outlined in section 5.2.1. However, the recent observations at optical and submillemeter wavelengths, published since the paper by Hall et al. in 1981, do shed some light on this issue. The suggestion by H alletal. that the nuclear H2 emission originates in molecular cloud complexes has received some support from the CO map of Myers and Scoville (1987) which shows that the molecular material is concentrated within 4 arcseconds of the nucleus. The most likely origin of the shocked lines is in cloud- cloud collisions, outflow from the active nucleus impinging on the surrounding dense molecular material, or outflows form young stars. The nuclear By flux is probably photoionised primarily by the active nucleus. This interpretation receives support from the optical line map of Baldwin, Wilson and Whittle (1987) which shows this portion of the nucleus to be dominated by high excitation lines. 5.2.4 Spatial Extent of the Infrared Lines In order to determine the spatial extent of the infrared lines the position of the continuum must be located. The signal to noise ratios on the H-window spectra are sufficiently high that the continuum level is unambiguous. The choice of the continuum in the K-window requires more explanation. The higher resolution spectrum of Hall et al. (1981) clearly shows the stellar absorption features due to C al (2.27 pm) and N al (2.21pm), so these points are likely to be below the continuum. It is possible to obtain a rough estimate of the equivalent width of these stellar features in the larger aperture spectra presented here by assuming that, in the absence of a nonstellar contribution to the continuum, the equivalentwidth of these features remain constant with aperture. The equivalent widths as measured by Hall et al. in their 4 arcsecond aperture are 0.74 and 0.57 A for Na I and Ca I respectively, and the active nucleus contributes “80% of the flux at K in the 4 arcsecond aperture (Halletal. 1981, McCarthy etal. 1982). Thecontinuum levels in the spectra presented here are » 0.4 Jy in the 8 arcsecond aperture, 0.7 in the 12 arcsecond aperture and 1 Jy in the 19 arcsecond aperture. The equivalent widths of the Na absorption features (assuming that the nucleus contributes 80% of the flux in a 165 4 arcsecond aperture) should be 1.4 A, 1.5 A an d 2 A, in the 8,12, and 19arcsecond apertures respectively, and those of Ca absorption features are 1 A, 1.15 A, an d 1.5 A. The continuum level in each aperture has been drawn so that the Na and Ca features have approximately these equivalent widths. It should be emphasised that higher resolution spectra, which resolve the Na and Ca absorption features, are necessary to locate the continuum.unambiguosly. Plots of flux versus aperture for the S(l) line, the By line and the [Fell] line are shown in Figure 5.4. It is clear that the flux in all lines is centrally concentrated. There is some suggestion that the 19 arcsecond By flux has not increased significantly above the 12 arcsecond flux, indicating that the emission does not extend further than 12 arcseconds. The 1 -0S(1) line may also show evidence for extent The extended (i.e. large minus small aperture) By flux is « lx l O'16 W nr2 and the extended S(l) line flux is 1.3x10‘16 W nr2. The exact spatial distribution of these lines cannot be determined from these measurements, and other interpretations are possible. The [Fell] line shows no evidence of spatial extent 166 Flux (10 S(l) line flux plotted against aperture. against plotted flux line S(l) Aperture (arcsec) Aperture iue 5.4a Figure 167 J t — ___ r l ___ 5 2 Flux (10 Byline flux plotted against aperture. against plotted flux Byline Aperture (arcsec) Aperture iue 5.4b Figure 168 t — r 5 2 Flux (10 [Fell] line flux plotted against aperture. against plotted line flux [Fell] prue arcsec) (a Aperture iue 5.4c Figure 169 5.2.5 Excitation Mechanisms As outlined in Chapter 1, the rotation-vibration spectrum of H2 has been observed to be excited by two processes: thermal excitation due to shocks (as in molecular outflows from young stellar objects) and UV fluorescence, observed in reflection nebulae (Dinerstein et al. 1988). The ratio of the 2-1 S(l) to the 1-0 S(l) line may be used to distinguish between these two processes, since a gas that is predominantly thermally excited will have a 2-1 S(l)line flux which is approximately 10% of the 1-0 S(l) line, whereas non-thermal excitation enhances the 2-1 S(l) line relative to the 1-0 S (l). It is clear from the 4 arcsecond spectra of Hall et al. that the line ratios in this aperture are characteristic of thermal excitation, probably behind a shock front However, in the 19 arcsecond aperture spectra the 2-1 S( 1) line is rather more prominent The flux in this line is ** 1 x 10-lfi W nr2, which is 75% that of the extended 1-0S(1) line flux of «1.3 x 10-16 Wm-2. This is more consistent with models of fluorescent excitation than with models of thermal excitation. It is possible that the low surface brightness extended component of the 1 -0S(1) line is fluorescently excited, whereas the nuclear emission is shock excited. However, as emphasised in the previous section, the resolution of the 19.6 arcsecond aperture is not high enough to determine the continuum unambigously, and so a high resolution off-nuclear spectrum is needed to confirm this discovery. Graham, Wright and Longmore (1987) have shown thatthe By to [Fell] line ratio can be used to distinguish shock excitation from photoionisation, (see Chapter 1). For a typical HII region (assuming case B recombination) this ratio is expected to be 0.06, whereas for shock excitation it is much higher. For NGC 1068 this ratio is between 1 and 10 in both apertures, indicating a substantial shocked com ponent 170 5.2.6 Origin of the Extended By Emission The extended component of the By line can be excited by two processes; a tircumnuclear starburst and/or ionising flux from the active nucleus. Here I examine these two hypotheses and discuss the implications for star formation in the disk and the energetics of the active nucleus. If the extended 10pm disk mapped by Telesco et al. (1984) and Telesco and Decher (1988) is due to a burst of star formation, there should be hydrogen recombination line emission associated with it The ratio of the ionising luminosity to bolometric luminosity (Lgy /L b o l ) of a HE region is a sensitive function of the spectral type of the ionising star. Table 5.2 shows the Lby /L b o l ratio expectedfrom HII regions ionised by stars of different spectral type. The values for the ionising luminosity were taken from Panagia (1973), and the By flux was calculated under the assumption of Case B recombination, outlined in Chapter 1. Hence the Lgy/Ltot ratio can be used to determine the spectral type of the stars that dominate the burst Table 5.2 Stellar Type Lb /L b o L xio*5 B0.5 0.56 B0 3.2 09 9.4 07 15 The distribution of 10pm emission in the central 20 arcseconds has not been measured directly since subtraction of the nuclear flux is very uncertain in this region. However, a lower limit to the flux in this region can be estimated by integrating the lowest contour of the 10pm map of Telesco and Drecher (1988) (60 mJy in a 4 arcsecond beam ) over this area. This gives a lower limit of 0.5 Jy at 10pm. Assuming that the integrated infrared flux is »10 times the 10pm flux (Wright etal. 1988), we have a lower limit to the infrared flux of 6xl0_l3 Wm-2* An upper limit to the integrated infrared flux can be set by assuming that the total flux (2x 10_11 Wm-2) in the extended disk as measured by Telesco et al. (1984) is concentrated in the inner 20 arcseconds. Assuming a value of ° 0.8x l0'16 Wm*2 (c.f. Figure 5.4) for the 171 extended component of the By flux, the lower limit to the infrared flux gives a Lfy/Ltot ratio of 26xl0-5 and the upper limit a ratio of 0.4x10-5. Therefore if the extendedBy emission is excited by a starburstthe spectral type of the stars is between BO. 5 and 05. Thus die extended By component is consistent with a recent burst of star formation in the 20 arcsecond disk. It is possible, however, that the By emission is associated with the Extended Narrow Line Region. H us region has been mapped out in [OHI] and HP by Baldwin, Wilson, and Whittle (1987), and seen in H a by Atherton, Reay and Taylor (1985) and BaHck and Heckman (1985) It is a region extending 10-12 arcseconds from the nucleus where die emission line line ratios are dominated by the high excitation tines, and also have broad (1000 km s~l) wings. This region is not situated symmetrically around the nucleus, the excitation map of Baldwin Wilson and Whittle showing it to be located preferentially along the radio lobes, suggesting that the ionising flux is preferentially beamed along the radio axis. The [OID] map from Baldwin, Wilson and Whittle (1987) is reproduced in Figure 5.5, along with the apertures for the infrared spectra shown in this study. Figure 5.5 [ODI] line map taken from Baldwin, Wilson and Whittle (1987). If a component of the extended By emission is associated with the [Om] and Hp, then it is reasonable to assume that it is excited by the same source, the nuclear radiation field. If this is the case, it is important to ask whether the ionising flux inferred from the By flux is consistent with that from other observations, such as theHP flux and the IUE data of Neugebaueretal.(1980). Assuming that the extended 172 By flux (0.8x10"16 W m*2) is associated with the Narrow lin e Region, and using standard recombination theory, an estimate of the ionising flux of 10 photons s" ^cm"2 is obtained. This is a strict lower limit to the ionising flux, since no allowance has been made for extinction or the filling factor of the Narrow Line gas. This should be compared to the value calculated from the Hp flux by Baldwin, Wilson and Whittle (1987) who also estimated a flux of 40 photons s* 1cm’2. These two estimates of the ionising flux are significantly higher than the figure of 2 ionising photons s*l cm*2 obtained by extrapolating the non-stellar continuum as presented by Neugebauer et al. (1980). Baldwin, Wilson and Whittle (1987) suggest that this discrepancy is due to the collimation of the flux along the radio lobes. If the extended By flux is indeed associated with the nuclear radiation this is additional evidence for collimation. There are three arguments in favour of the hypothesis that the extended By emission may be associated with the Extended Narrow Line Region. Firstly, the By linewidth (as measured by Hall et al. 1981) is 831200 kms* 1, which matches the [OIII], HP and H o linewidths of * 1000 kms* 1 measured by Baldwin, Wilson and Whittle (1987) and Atherton, Reay and Taylor (1985) respectively. Secondly, these lines extend out to between 10 and 12 arcseconds, as the By line may do (Baldwin Wilson and Whittle 1987 and Atherton, Reay andTaylor 1985). Finallythe [Offl]/HP ratio in the area encompassed by our 20 arcsecond aperture is characteristic of a hard’ radiation field associated with the active nucleus (Baldwin, Wilson and Whittle 1987) The star formation in the infrared disk as mapped by Telesco and Decher (1988) correlates well with the CO ring (Myers and Scoville 1987). That there is little molecular gas in the region between 4 and 20 arcseconds of the nucleus argues against the star formation hypothesis. Ionisation by the active nucleus is consistent with the distribution of the molecular gas, and the energetics and dynamics of the By line match that of the Extended Narrow Line Region. These data do not allow the exclusion of star formation as the source of the extended By emission. On balance, however, I suggest that the flux is dominated by ionising radiation from the active nucleus. This conclusion, if confirmed, undermines the common assumption that the By line traces star formation. In general, therefore, By line fluxes should only be used to determine star formation rates if it is clear there is no contribution to the flux from an active nucleus. 173 5.2.7 Distribution and Excitation of the 1-0S(1) and [Fell] Lines In section 5.2.5 it was suggested that the extended 1-0S(1) line flux may be fluorescently excited. Models of fluorescent excitation indicate that one S(l) photon is emitted for every 100 Lyman continuum photons absorbed (Kwan 1977; see Chapter 1), thus for fluorescent excitation to be a viable hypothesis there must be sufficient Lyman continuum photons to produce the observed 1 -0 S(l) line flux. The extended By flux implies a Lyman continuum flux of « 6 photons s" * cm"2 compared to 0.13 photons s"l cm‘2 in the extended component of the 1-0S(1) line. This indicates that only ** 50 Lyman continuum photons are produced for each 1-0S(1) line photon. Thus for fluorescence to be a plausible mechanism we require that a t least 90% of the Lyman continuum photons are not intercepted by the Narrow Line gas producing the By recombination line. This may occur if the filling factor of the narrow line gas is low, if the optical depth in the Lyman continuum is small, or a combination o f both. The CO map of Myers and Scoville (1987) shows that the dense molecular material is confined to within 4 arcseconds of the nucleus. In this context, it may be significant that the [Fell] line, which has a substantial shocked component, shows no evidence for extent between 8 and 12 arcseconds. This is generally consistent with the suggestion that the shocked excited lines are associated with dense molecular material, whereas the fluorescently excited emission is associated with much more diffuse m aterial. 5.3 Near Infrared Spectroscopy of NGC 4151 - Evidence for Brackett Line Variability NGC 4151 is the best studied example of a Seyfert 1 galaxy. The original motivation for this work was to determine the spatial extent of the 1-0S(1) line and the By line, to compare with the spatial extent in other galaxies. A detailed study of the spatial extent of these lines in NGC 4151 has not been possible, however, partly because the spectra are noisy and the 1 -0S( 1) line has not been clearly detected. By comparing the By line strength in the data presented in this thesis with those from other workers, it can be shown that the flux in this line has varied on timescales of a few months, whereas the flux in the 1 -0S( 1) has remained constant The implications for the origins and excitation of these lines will be discussed in this chapter. 174 5.3.1 Variability in NGC 4151 The nucleus of NGC 4151 is the canonical example of Seyfert variability. It has been well studied at all wavelengths, especially at UV and X-ray wavelengths where the contribution from star light to the integrated light is minimal (Malkan et al. 1987). The timescale for the continuum brightness at UV to double is 5-30 days, which indicates that the continuum region is « 0.0 Ipc across. The broad lines detected by IUE (e.g. CIV) also show evidence for variability on these timescales, but these these lines frequently lag behind the continuum by about 3 days, suggesting that the BLR is 1-3 light days across. In contrast the X-ray flux can double in approximately 12 hours, showing that this region is smaller (by a factor 10) than the BLR. The variability of NGC 4151 has been confirmed in the near infrared (Cutri et al. 1981),and has also been studied in the optical. An upper limit to the timescale for optical variability is 30 days (for the H a line). The timescale is not well determined and there is considerable evidence that the line and continuum fluxes vaiy on timescales shorter than this, but data on these short timescales are scanty. Another problem facing optical observers is the subtraction of the line and continuum fluxes due to the stellar component in NGC 4151. 5.3.2 The Data The calibrated spectra for this study are shown in Figure 5.6. They consist of K-band spectra in 5.4,7.8, and 19.6 arcsecond apertures. The By line has been detected in all apertures. The measured By line fluxes, 3a upper limits to the 1 -0S( 1) line flux, and continuum levels are shown in Table 5.3. Also shown in Table 5.3 are previous observations of NGC 4151 reported by Fischer et al. (1987). 175 Table 5.3 Date Aperture By flux S(D flux Continuum ref (arcseconds) (xlO -17 Wm-2) (m Jy) Ja n 1984 5.4 not observed 3.0±0.4 130 Fischer et al. 1987 Feb 1985 5.4 17±1 3.0±0.4 185 9 9 Jan 1986 5.4 5.0±0.6 < 3.0 125 This thesis 7.8 6.4±1.0 m H 19.6 8.2±0.8 9 9 176 Flux 19.6 arcsecundK-window spectrum of NGC4151. of spectrum arcsecundK-window 19.6 aeegh (yum) Wavelength iue 5.6a Figure 177 Flux (mJy) 7.8 arcsecond K-window spectrum of NGC 4151. NGC of spectrum K-window arcsecond 7.8 aeegh (jum) Wavelength iue 5.6b Figure 178 Flux (mJy) 5.4 arcsecond K-window spectrum of NGC 4151. NGC of spectrum K-window arcsecond 5.4 aeegh (/mi) Wavelength iue 5.6c Figure 179 5.3.3 Evidence for Variability Comparison of the 5 arcseccmd aperture 2 pm observations taken in January 1986 with those of Fischer et al. taken in February 1985 shows that during this period the line Byline flux has apparently decreased by a factor of 3, from *»17xl0-17 to *5xl0 -17 Win'2. The 2 pm continuum has decreased by a factor of 1.5 over the same period. There are two obvious sources of error that could produce such a discrepancy; a simple calibration error and a pointing error. In January 1984 Fischer et al. obtained a 5 arcsecond aperture spectrum of the 1-0S(1) line of molecular hydrogen which shows the 2pm continuum at *130 mJy. The By line was not observed on this occasion. This estimate of the continuum level is consistentwith our measurement in January 1986 of 125 mJy. In February 1985 Fischer et al. obtained a spectrum of the S( 1) line and the By line which shows the 2pm continuum has increased to * 185 mJy. A simple calibration error is unlikely to account for this difference between the continuum level observed by Fischer et al. in January 1984 and February 1985 since the flux in the S(l) line is the same on both occasions, rather than scaling with the continuum. Similarly, if the calibration on our 1986 spectrum has been underestimated by a factor of 1.5, it would be expected that the By line flux would be reduced by a factor 1.5 and not 3 .1 therefore conclude that a simple calibration error is unlikely to account for the variation seen in the continuum fluxes and By line fluxes. A pointing error could in principle produce discrepant line fluxes. However, the multiaperture data presented here show that in the 19.6 arcsecond aperture the By line flux is still a factor of 2 below the Feb 1985 estimate of the line flux from Fischer et al. Thus a pointing error of ^ 10 arcseconds is required to explain this discrepancy. There is other evidence that the Brackett line variability is real. McAlary and McLaren (1981) reporta By line flux of19x l0 -17 W nr2 and acontinuum level of 170 mJy in a 7 arcsecond aperture from observations dating from February 1978 and March 1979. This is similar to Fischer et al.’s 1984 data. Finally, Figure 5.7 shows the variation in the By line and 2pm continuum compared to the variation seen in the HP and Hell broad line fluxes and the optical continuum. These data were taken from taken from Peterson and Cota (1988). The optical lines and continuum show evidence for variability on timescales of 30 days, much shorter than the 12 month period between the Brackett line measurements. There is, however, a tendency for the optical lines and the continuum to decrease between January 1985 and January 1986. 180 This strongly suggests that the decrease in the By line flux is real, and is associated with the active nucleus. 181 Flux (W By, Hell and HP line fluxes during the period Januaiy 1985 to January 1986. January to 1985 Januaiy period the during fluxes HP line and Hell By, a 1 a 1 a 1 u 1 e 1 o 1 a 1 Jan 1 Nov 1 Sep 1 Jul 1 May 1 Mar 1 Jan 95 1986 1985 iue 5.7a Figure 2 8 1 Flux (mJy) Optical and infrared continuum flux densities during 1985 during to the January period densities and Optical continuum flux infrared a 1 a 1 a 1 u 1 e 1 o 1 a 1 Jan 1 Nov 1 Sep 1 Jul 1 May 1 Mar 1 Jan 95 1986 1985 iue 5.7b Figure January 1986. January 183 5.3.4 Implications for Origin and Excitation of Infrared Lines It is clear from the above discussion that the By emission from NGC 4151 contains a large, and at times dominant, contribution from the Broad Line Region of the active nucleus. This highlights the dangers of using infrared recombination line fluxes to determine star formation rates (Beck, Beckwith and Gatley 1984) in the nuclei of active galaxies without taking into account the contribution to the line fluxes from the nuclear radiation field. For example, Fischer et al. used their February 1985 measurement of the By flux to determine the star formation rate in NGC 4151. Fischer e t al also concluded that the ratio of the star formation rate determined bythe 1-0S(1) line flux to the star formation rate determined by the By flux was «0.6, which indicates that the star formation rate in NGC 4151 is sufficient to account for the 1-0S(1) line luminosity in terms of outflows from young stars. This is in contrast to the other Seyferts they studied where this ratio is between 2 and 15. Since their 1985 measurement of the By line flux overestimates the contribution from stars by at least a factor of three, this ratio for NGC 4151 becomes «1.8, which is comparable with the other Seyferts. It is significant that the 1-0S(1) line does not show any evidence for variability. This underlines the point made by Fischer et al. (based on determining the 1 -0S( 1) linewidth) that the molecular hydrogen emission is not associated with the Broad Line Region of the active nucleus. However, Fischer et al. also concluded on the basis of the 1985 By flux that there are sufficient photons in the ionising continuum to excite the molecular hydrogen emission fluorescently (Black and van Dishoeck 1987). The 1986 By flux, however, implies that the number of ionising photons at that time was too small to be able to account for the S(l) line flux. This and the lack of variation in the molecular hydrogen lines argue against fluorescent excitation. 5.4 Summary and Conclusions This chapter presents near infrared spectra of two classic Seyfert galaxies, NGC 4151 and NGC 1068. Multiaperture spectroscopy of NGC 1068 shows that the By recombination line has an extended component. The [OIH]/H (3 ratio in the area of the extended By emission is characteristic of a hard’ radiation field, and the By linewidth 184 is larger than is normally seen in star formation regions. These two observations indicate thatthe By recombination line is associated with gas photoionised by the active nucleus, although it is not possible to rule out star formation unambiguously. Assuming thatthe extended By emission is associated with the active nucleus, the ionising flux inferred from the By flux is larger than that measured by IUE. This may well indicate that the ionising photons are collimated, perhaps along the radio lobes. The By line flux in NGC 4151 is shown to vary on timescales of months; this shows clearly that the recombination line flux from this galaxy must be dominated by the active nucleus. Observations of the By recombination line from these Seyfert galaxies undermine the common assumption that the By line always traces star formation. The 1-0S(1) line in NGC 1068 also has an extended component The 2- 1S(1) line is seen clearly in the 19 arcsecond spectrum presented here, while it was not detected in the small aperture spectrum of Hall etal. (1981). This suggests that while the nuclear component of the H2 emission is certainly thermally excited, the extended emission may be fluorescently excited. In order for fluorescence to be energetically viable, however, at least 90% of the Lyman continuum photons are not intercepted by the narrow line gas producing the By recombination line. The 1-0S(1) line in NGC 4151 was not detected in the spectra presented in this thesis. However, the spectra of Fischer e t al (1987) demonstrate that this line shows no evidence for variability, while the recombination lines and the infrared continuum do. This suggests that the H2 emission in NGC 4151 is not dominated by fluorescence. 185 Chapter 6 A Statistical Study of Infrared Lines in Galaxies 6.1 Introduction A stronomers commonly use two techniques for investigating astrophysical phenomena. The first involves an in-depth study of a particular object, and the second a statistical study of a number of obj ects selected to have a key property in common. Statistical studies are particularly useful for determining broad trends in data from which physical inferences can be drawn, and also for comparing the properties of two groups of objects with differing selection criteria. The aim of the analysis in this chapter is to determine the relationship between the infrared lines studied in this thesis and the infrared continuum fluxes measured by the Infrared Astronomical Satellite, IRAS. This analysis will be carried out for two groups of objects; galaxies known to have an active nucleus, and starburstfnteracting galaxies. 6.2 Data and Statistical Analysis In this section I will outline the selection criterion for the galaxies used in this study, discusse the biases found in the data, and then describe in detail the statistical techniqu es used to determine the relationships between the infrared line fluxes and the IRAS continuum fluxes. 6.2.1 The Sample The data used for this study consist of about 30 galaxies with measured By and/or S( 1) line fluxes and detections in all four IRAS bands. A large fraction of galaxies with measured near infrared line fluxes contain active nuclei, and since the physical processes that give rise to the line and continuum radiation in active nuclei differ from those in ordinary galaxies, the sample has been subdivided into two groups, those containing active nuclei and a sample of interacting/merging galaxies. The galaxies selected were unresolved by IRAS. The data for all galaxies was taken from the Extragalactic Point Source Catalogue (Lonsdale et al. 1985). The only galaxies with cirrus flags were NGC 3256, NGC 7469 and NGC 6052. Tables 6. la and 6. lb show the By and S (l) line fluxes, the IRAS fluxes, the optical diameter, the recession velocity and the aperture size used for the infrared line observations for starburst galaxies and Seyfert galaxies respectively. The final column gives the 186 reference for the infrared line flux. All optical diameters and distances were taken from the Second Reference Catalogue of Bright Galaxies (de Vaucouleurs, de Vaucouleurs and Corwin 1976). Figure 6.1 shows the K-window spectra of two galaxies, NGC 2798 and NGC 4631 and NGC 3079 which are not presented anywhere else in this thesis,but have been included in this statistical analyisis. 187 Flux (10 19.6 arcsecond K-window spectrum of NGC4631 of spectrum K-window arcsecond 19.6 aeegh Gum) Wavelength iue . (a) 6.1 Figure 8 8 1 CM I 1 Flux (10 E E 400 0 -4 0 0 -2 0 - 200 19.6 arcsecond K-window spectrum of NGC 2798 NGC of spectrum K-window arcsecond 19.6 j _ 2.05 I _ i t — —r i— . 21 22 .5 2.3 2.25 2.2 2.15 2.1 aeegh (/xm) Wavelength t — r iue . (b) 6.1 Figure 189 t — r ] —i —i —i —i —[ i— i— i— i— |— i— i— i— i— ■]— Flux (10 96 rscn K-idwsetu ofNC 3079 NGC f o spectrum -window K arcsecond 19.6 aeegh (jum) Wavelength iue . (c) 6.1 Figure 190 Table 6.1a Starburst Galaxies G alaxy By S (l). 12 25 60 100 V C25 A Ref (NGQ lO -l^W m -2 Jy Jy Jy Jy km s’ l arcsecs 5253 14.8 0.73 2.59 12.21 30.91 29.04 209 240 13.5 3 He2-10 4.4 0.9 1.1 6.55 23.78 25.68 560 78 5.3 3 1614* 4.6 4.5 1.39 7.59 33.18 31.61 4745 79.08 19.6 1 3690* 10.0 9.0 3.73 21.57 105.4 109.7 3104 144 19.6 1 4194* 4.3 - 0.86 4.39 22.5 25.2 2528 147 5.5 2 3256* 14.8 5.7 3.23 15.5 94.21 120.3 2592 204 13.5 3 5135 2.2 - 0.68 2.49 15.94 30.05 1.0 144 5.5 2 ESO 286-IG 0.23 0.46 0.33 1.92 12.47 10.19 12783 69 4.5 3 ESO148-IG 0.64 0.42 0.3 1.67 10.87 10.02 13267 46.8 4.5 3 6240 0.6 26.8 0.57 3.52 23.21 25.88 7597 132 19.6 1 1482* 2.4 2.4 1.6 4.65 31.21 45.54 1542 93 5.3 3 2798* 3.76 3.8 0.77 3.13 22.48 28.36 1708 165.2 19.6 1 986* 2.1 0.9 0.82 3.1 22.81 48.46 1943 222 5.3 3 6052* 3.0 1.5 0.31 0.88 6.94 9.74 48.21 57 19.6 1 1672 2.2 1.9 1.48 4.06 34.47 68.25 1076 288 5.3 3 4433* 1.1 - 0.6 1.43 13.17 25.2 2907 137.4 5.5 2 520* 1.9 - 0.78 2.85 31.21 47.41 2168 287.2 5.5 2 Arp220 1.3 5.2 0.48 8.15 103.68 116.25 5420 80.4 19.6 1 2623* 0.8 - 0.36 1.75 24.06 27.53 5355 37.2 19.6 1 4631 1.0 4.2 1.82 3.01 51.15 118.6 638 909.6 19.6 1 3079 - 19.0 1.25 2.04 42.52 87.63 1212 454.8 19.6 1 7714* 4.7 - 0.49 2.81 11.07 10.92 2800 21.4 5.0 5 6810 1.2 1.8 1.06 3.51 17.84 33.65 1800 38.0 6.0 4 4038 - 2.3 1.22 3.92 38.9 74.6 1445 26.0 13.5 3 IC4687* 0.3 0.66 22.0 14.99 26.36 5000 13.99 13.5 3 1792 0.54 1.39 1.86 23.0 73.6 977 40 5.3 3 References 1 - This thesis 4 Moorwood and Oliva (1988) 2 DePoy (1987) 5 Ho,Beck and Turner (1989) 3 Kawara, Nishida and Gregory (1987) 191 Table 6.1b Seyfert Galaxies Galaxy 12 25 60 100 V D25 A Ref (NGC) 10-17Wm'2 Jy Jy Jy Jy k m s~ l arcsecs 1068 30.0 27.7 38.3 86.83 185.58 238.7 1134 414 19.6 1 3227 4.25 1.2 0.67 1.75 7.84 16.93 1050 337.2 19.6 1 7469 5.0 - 1.3 5.5 26.67 34.4 4894 106.2 5.5 2 5643 2.4 3.9 0.87 3.37 18.45 43.26 962 276 13.5 3 7552 6.5 3.0 2.99 12.01 72.52 99.46 1636 210 5.3 3 1808 6.1 2.0 4.14 15.94 96.69 134.92 769 432 5.3 3 1365 4.2 1.7 3.22 11.13 77.75 139.94 1502 588 5.3 3 7582 4.4 - 1.36 6.37 47.63 71.49 1427 92 5.3 3 613 1.4 - 0.74 2.09 19.3 48.12 1500 57.5 6.0 4 6221 3.4 - 1.51 5.31 35.96 83.23 1254 31.6 6.0 4 5506 5.7 - 1.3 3.66 8.67 9.4 1690 28.8 6.0 4 7496 3.7 - 0.33 1.53 8.42 14.81 1443 34.7 6.0 4 1275 6.3 1.03 3.63 7.09 238.7 5223.9 25 12 5 A1409-56 10.0 18.9 68.39 248.8 313.3 172.6 31.9 13.5 3 2992 1.0 0.6 1.38 6.76 13.99 1862 41.0 13.5 3 4418 1.0 0.93 9.61 43.45 32.9 2098 15.6 13.5 3 6300 1.5 0.75 2.16 13.8 39.35 988 53.95 13.5 3 6221 0.6 1.51 5.3 35.97 83.17 1253 31.9 13.5 3 4038 2.3 1.22 3.92 38.9 74.6 1445 26.0 13.5 3 References 1 This thesis 4 Moorwood and Oliva (1988) 2 DePoy (1987) 5 Fischer et al. (1987) 3 Kawara, Nishida and Gregory (1988) 6.2.2 Biases due to Distance Effects Statistical studies in extragalactic astronomy are often plagued with biases due to distance effects. In this section the main distance biases are illustrated by examining the correlation between the By fluxes and IRAS continuum fluxes for the starburst galaxies in this sample. The simplest way to look for correlations between the strength of the IRAS bands and the Byline strength is to determine the least squares correlation coefficient between the Brackett line fluxes and the IRAS fluxes. However, there is a trend for both the By fluxes and the IRAS fluxes to decrease with increasing redshift This is illustrated in Figures 6.2a- 6.2e, which show the By flux and IRAS band fluxes plotted against recession velocity. Thus a correlation between any two fluxes may simply reflect the fact that they are both independently correlated with distance. -15 -16 tt 0 O -J -17 -18 2 3 4 5 log v Figure 6.2a By flux (in Wnr2) plotted against recession velocity (in km s_1) for starburst galaxies. 193 0.6 0.4- 0. 2 - N e 0.0 - o - 0.2 - -0.4 - -06 3 4 5 logv Figure 6.2b 12 jim flux (in Jy) plotted against recession velocity (in km s*1) for starburst galaxies. 2 1 Vi N (9 O 0 2 3 4 5 logv Figure 6.2c 25 pm flux (in Jy) plotted against recession velocity (in km s_1) for starburst galaxies. 194 log v Figure 6.2d 60 pm flux (in Jy) plotted against recession velocity (in km s_1) for starburst galaxies. 2 2 - i * ■ ■■ ■ 2 .0 - ■ ^ B 1.8 - B * Q 1.6 “ _ ■ 1 m * ■ § 1.4 - m J 1.2 - 1.0 - ■ . * 0.8 - ...... T ' ' 1 " ' 1 1 2 3 4 5 log v Figure 6.2e 100 pm flux (in Jy) plotted against recession velocity (in km s*1) for starburst galaxies. This problem can be overcome by normalising the fluxes by some parameter that removes this distance dependence. In principle, this can be done by looking for correlations between By luminosities and the luminosities in each of the IRAS bands, where the luminosity is given by Lm4itlP f. Here f is the flux in the band being considered, and D is the distance to the galaxy. Problems arise however, when the 195 range in D2 is much larger than the range in f. In this case, the luminosity in each band of a given galaxy may be dominated by the square of the distance. A correlation between the luminosities in any two bands reflects the fact that both luminosities are dominated by D2, die flux introducing only a small amount of scatter For the galaxies examined here, the range in D2 « 104 whereas the range in fluxes (both for Brackett line fluxes and IRAS bands) is * 10. Thus this effect must be investigated in detail. This spurious correlation can be illustrated by examining the correlations obtained between the By line luminosity and the 100pm IRAS band using random By and IRAS fluxes. Uniformly distributed random Brackett line and 100pm IRAS fluxes were assigned to each galaxy, the dispersion in the random fluxes being the same as in the observed fluxes. All random fluxes were obtained using the random number generator on a Casio fx-3800P pocket calculator. Table 6.2 shows the values of the means and standard deviations of the real fluxes and the random fluxes. The observed redshift of the galaxy was then used to calculate a (random) luminosity. Figure 6.3 shows these random fluxes plotted against each other, to demonstrate that there is no correlation between them. Figure 6.4a shows the By luminosity plotted against the 100 pm luminosity using the observed fluxes, while Figure 6.4b is the same plot except that the luminosities have been calculated using random fluxes. Inspection of these two figures shows that the ’random flux’ plot (least squares correlation coefficient 0.71) is as well correlated as the ’real flux’ plot, (coefficient 0.69). This suggests thatthe correlation between the By line luminosity and the 100pm luminosity is due to this distance bias, and therefore some other method must be used to investigate any correlation between the intrinsic strength of the Brackett lines and the IRAS bands. Table 6.2 Real Random 1 0 0 |jm By 1 0 0 Jim By m ean 50.0 3.6 53.7 6.01 s.d 36 4.1 37 3.1 196 3 2 - O 1 O ? m l % o H— 1 i — -17.0 - 16.5 -16.0 -15.5 Log BY Figure 6.3 Random By flux (in Wm*2) plotted against random 100 pm flux -10 a -i i - 8 9 100 imn Luminosity Figure 6.4a By luminosity plotted against 100 pm luminosity, using observed fluxes 197 -7 -12 H------1------1------1 ...... 6 7 8 9 10 11 100 pm Luminosity (random) Figure 6.4b By luminosity plotted against100 pm luminosity, using random fluxes The practice of correlating luminosities is widespread, and many cases can be found in the literature. For example, Sekiguchi (1987) proposed a model to explain the infrared emission from starburst and normal galaxies in which the infrared luminosity is the sum of two components; a warm component (T“80-90K) associated with dust in HII regions and a cool component (T«30K) arising from galactic disks. He claimed that the warm component and the cool component luminosities (calculated on the basis of his model) correlated with the observed integrated infrared luminosity. The basis for this claim is two plots, reproduced in Figure 6.5, showing the warm and cool component luminosities plotted against the observed IRAS infrared luminosity. The variables are apparently well correlated, however, the range in fluxes is * 10 and the range in distance-squared is *70 and so the possibility of spurious correlations must be investigated. 198 Warm and cold component luminosities plotted against infrared luminosity, taken from Sekiguchi (1987) In order to test the reliability of these correlations, random fluxes were assigned to each galaxy. The flux of a galaxy was set by randomly assigning to it the observed flux of another galaxy. The random luminosity was then calculated from this spurious flux. The results of this analysis is shown in Figure 6.6, where the spurious warm and cold component luminosities are plotted against the spurious infrared luminosity. The correlations obtained using random fluxes are as good as those obtained using the real fluxes. The correlations obtained by Sekiguchi may be spurious, and further analysis is required to demonstrate their validity. 199 Figure 6.6a Random warm component luminosity plotted against random infrared luminosity e e E e u e e u Figure 6.6b Random cool component luminosity plotted against random infrared luminosity The final problem to be examined in this section is whether the range of aperture sizes used for the By fluxes (5 to 19 arcseconds) will have effect on the analysis. This effect is likely to he minimal, because the area encompassed within an aperture on any given galaxy is dominated by the distance of the galaxy. For example, a 5 arcsecond aperture encompasses an area of 0.48 Kpc on a galaxy with a recession velocity of 1000 km s’1 (approximately the distance of the closest galaxy), 2 0 0 and 4.8 Kpc on a galaxy with a recession velocity of 10,000 km s- 1 (the recession velocity of the furthest galaxy). A 19 arcsecond aperture on a galaxy with a recession velocity of 1000 km s_1 will encompass an area of 1.8 Kpc, less than the area within a 5 arcsecond aperture for a galaxy at 10,000 km s_1. The difference in size between the apertures used for the Brackett line fluxes and the (much larger) IRAS apertures will be discussed in the results section. The main aim of the discussion in this section is to highlight die major problems associated with extragalactic correlation analysis. Investigating correlations between fluxes is valid so long as there is no systematic change in flux with distance. Correlations between luminosities are often spurious, and can result in plotting distance-squared on both axes. Luminosity-luminosity plots are frequently shown in the literature, and can lead to misleading results, as in the paper by Sekiguchi (1987) described above. In the following section two techniques are described which circumvent these problems. 6.2.3 The analysis The problems outlined in the previous section can be overcome by two methods. The first involves selecting a subset of the galaxies that have a sufficiently small spread in distance that there is no systematic decrease in flux (in any band) with distance. Then the fluxes in any IRAS band can then be compared directly with the By flux. This method is relatively straightforward, butwasteful in that a large fraction of the sample is thrown away. The second involves the use of non-parametric partial correlation analysis (Macklin 1982). This is a technique used to determine the statistical correlation between any two variables, in the case where the two variables are dependent on one or more other variables. This method will be described in detail in the next fewparagraphs, starting with a description of the non-parametric correlation analysis, and then widening the discussion to include non-parametric partial correlationanalysis (i) Non-parametric analysis The simple linear or Pearson’s correlation coefficient is essentially an estimate of the residuals to be expected if the data are fitted to a straight line using a least squares fit It is given (Barford 1985) by: £i (xj - 2 0 1 where x* and yj represent two independent samples with mean values is distributed like Students t-distribution, which is approximately normally distributed about 0 with unit variance when n is large (e.g.Cramer 1946). Thus alarge (|t |> 1) value of t suggests that the null hypothesis can be rejected. A useful test of significance is the probability (Pr) that 111 has its observed value or larger just by chance. A small value of Pr (e.g. Pr = 10-3) indicates that the null hypothesis can be rejected with confidence. The Pearson’s correlation coefficient is limited, however, since it requires that the variables being considered are normally distributed. In particular, it is very sensitive to ’outliers’ - points with xj or yi values several standard deviations away from the mean. This limitation can be overcome if non-parametric correlation analysis is used. (e.g. Press et al. 1986) The basis of this technique is that the xj andyt are ordered from highest to lowest, and then each assigned a number Rj and Sj that reflects its rank within the sample. Statistical tests are then performed on the two sets of consecutive integers, Ri and Si. This means that the distribution function from which the two variables are drawn is well known ( i.e. integers uniformly distributed between 1 and n, where n is the number of data points) and so overcomes the problems associated with the Pearson’s correlation coefficient. One of the most commonly used non-parametric estimates of the strength of a correlation is the Spearman Rank Order Correlation Coefficient, rs. This is obtained simply by substituting the Ri for the xi and the Sj for the y\ in equation (1). Ii (Ri - The significance of rs can also be derived by analogy with Pearson’s correlation coefficient; with r s substituted for r, the t-statistic is distributed like Students distribution. 2 0 2 (ii) Partial Correlation Analysis Partial correlation analysis has been developed to determine the statistical correlation between any two variables, in the case where the the two variables are dependent on one or more other variables. LetX and Y be random variables, X is dependent on variable A, Y on variable B, and both are dependent on R. Then LogX = LogXR + LogA (6.4) LogY = LogpR + LogB (6.5) where X and p are constants. We need to know whether the correlation between X and Y is due to their common dependence on R, or to a correlation between the variables A and B - ie we require the correlation coefficient between the variables A and B, tab . This correlation coefficient is given by equation (6.3) With some algebra it is possible to show that rAB = r XY,R w hererxY,R is the Partial Correlation Coefficient of X and Y. This is given by rxY,R° (6.6) v(l-rYR 2)(l-rxR2) where rxY is the Spearman correlation coefficient between the variables X and Y, ryR the Spearman correlation coefficient between Y and R etc. It can be shown that if correlation of X and Y is due entirely to their separate dependences on R (i.e. the population partial correlation coefficient denoted by pxY,R»is zero) that the statistic t’ = rxY,R- \ f ^ (6-7) is distributed according to Students t-distribution. Once again, the non-parametric partial correlation coefficient and its significance can be derived by substituting the rank of a variable in the sample for its value. 203 6.2.4 Comparison of Partial Correlation Analysis with Restricted Redshift Range Analysis Both of the methods (restricting the redshift range and using partial correlation analysis), outlined in section 6.2.2 have been used to determine the relation of the By fluxes to the IRAS bands for the starburst galaxies. The results of this analysis are given in this section, to demonstrate that partial correlation analysis on the whole sample gives the same results as ordinary analysis on the restricted redshift range sam ple. In Table 6. la the galaxies selected to be included in the restricted redshift sample have been marked with an asterisk (*). Figures 6.7a-e show plots of flux vs redshiftfor this sample, and Figures 6.8a-d show the By fluxes plotted against the flux in the IRAS bands. There is a clear correlation between the By fluxes and the IRAS bands, which cannot be due to a common dependence on distance because Figures 6.7a-e demonstrate that there is no systematic decrease of flux with redshift Table 6.3 shows the Spearman Rank correlation coefficient (rs) and the significance of this result (Pr) derived from Students distribution for the plots shown in Figure 6.8. Table 6.4 shows the results of the partial correlation analysis on the whole sample of starburst galaxies. Table 6.3 Spearman correlation analysis for restricted redshift range sample 1 2 25 6 0 10 0 r s 0.82 0.93 0.75 0.72 Pr <0.005 <0.005 0.008 0.01 204 Table 6.4 Partial correlation analysis, whole sample starburst galaxies 1 2 25 60 10 0 u 0.47 0.55 0.024 0.05 Pr <0.025 0.007 0.26 0.82 -15.8 -16.0 -16.2 s CD -16.4 e -16.6 -16.8 -17.0 31 3.2 3.3 3.4 3.5 3.6 3 7 3.8 log v Figure 6.7a B y (Wm-2) flux plotted against recession velocity (km s* l)for the restricted redshift range sample 205 0.6 0.4- 0 .2 - N 0.0- O O - 0.2 - -0.4 - - 0.6 - - 3.1 3.2 3.3 3.4 3.5 3.6 3.7 3.8 log v Figure 6.7b 12 jim flux (in Jy) plotted against recession velocity (km s_ l)for the restricted redshift range sample (S o o ml 3.1 3.2 3.3 3.4 3.5 3.6 3.7 3.8 log v Figure 6.7c 25 pm flux (in Jy) plotted against recession velocity (km s-^for the restricted redshift range sample 206 2.2- ■ 2.0- ■ 1.8- g 1-6- ■ ■ ■ § 1.4- mI ■ ■ ■ ■ 1.2- ■ ■ ■ 1.0 - ■ 0.8-...... 1 .... 1 " 1 ...... 1" ”F ” 1 3.2 3.3 3.4 3.5 3.6 3.7 3.8 log v Figure 6.7d 60 pm flux (in Jy) plotted against recession velocity (km s-^for the restricted redshift range sample 2.2 2.0 I 8 § 16 SO 1.4 wJ 1.2 1.0 0.8 3.1 3.2 3.3 3.4 3.5 3.6 3.7 3.8 log v Figure 6.7e 100 pm flux (in Jy) plotted against recession velocity (km s- J)for the restricted redshift range sample 207 -17.0 -16.8 -16.6 -16.4 -16.2 -16.0 -15.8 LOG By Figure 6.8a By flux (in Win-2) plotted against 12pm flux (in Jy) for the restricted redshift range sample in N o o -17.0 -16.8 - 6.6 -16.4 -16.2 -16.0 -15.8 LOG By Figure 6.8b By flux (in Wm-2) plotted against: 25 pm flux for the restricted redshift range sample 208 2 . 2 -j ■ 2.0 - ■ 1.8 - § I * ' ■ ■ ■ § 1 .4 - ■ ■ ■ ■ 1.2 - ■ s 1 .0 - 0 . 8 ------i— i------1—* ------1 ' i 1 "i------17.0 -16.8 -16.6 -16.4 -16.2 -16.0 -15 8 LOG By Figure 6.8c By flux (in Win'2) plotted against " \ 60 pm flux for the restricted redshift range sample 2.2 2.0 1.8 o o 1.6 o 1.4 1.2 1.0 0.8 -17.0 -16.8 -16.6 -16.4 -16.2 -16.0 -15.8 LOG Br Figure 6.8d By flux (in W nr2) plotted against: 100 pm flux for the restricted redshift range sample The results obtained with the restricted redshift range sample show that the By fluxes are correlated with all four IRAS bands, but that the significance of the correlation is greater, by an order of magnitude, at 12 and 25pm (Pr <0.005) than at 209 60 or 100jim (Pr «0.01) The partial correlation analysis of the whole sample shows this discrepancy as well; there are highly significant correlations at 12 and25jim with the By flux (Pr « 0.03 and 0.007 respectively), whereas the 60 and 100pm flux is barely correlated (Pr = 0.26-0.28). The results from the partial correlation analysis for the whole sample agree with those from the Spearman correlation analysis for the restricted redshift range sample. This demonstrates that partial correlation analysis removes the distance dependence in the flux-flux correlations, and so partial correlation analysis will be used in the remainder of this Chapter. 6.3 Starburst Galaxies The results of the partial correlation analysis for the sample of starburst galaxies will be presented in section 6.3.1, and the implications discussed in section 6.3.2. 6.3.1 Results The results of the partial correlation analysis for the starburst galaxies are shown in Tables 6.5 and 6.6. Table 6.5 shows r, t and Pr for the By flux correlated against the IRAS bands, and Table 6.6 shows r, t and Pr for the S(l) line flux correlated against the IRAS bands. Table 6.5 Brackett line correlations for starburst galaxies 22 objects Band r t Pr 12pm 0.47 2.36 0.028 25 pm 0.55 3.02 0.007 60pm 0.24 1.14 0.26 100pm 0.05 0.22 0.82 210 Table 6.6 S(l) line correlations forstarburstgalaxies 20 objects Band r t Pr 12pm 0.07 0.3 0.76 25 pm 0.06 0.25 0.8 60pm 0.34 1.62 0.12 100pm 0.46 2.29 0.03 The results presented here show that the By line flux is well correlated with the 12 and 25 pm fluxes but is not significantly correlated at 60 or 100 pm. This is in contrast to the S( 1) line flux which is only significantly correlated at 100 pm. 6.3.2 Brackett Line Correlations It is surprising to And that the By line luminosities are correlated with the luminosity in any of the IRAS bands, since the aperture sizes used for the Brackett line observations (up to 19 arcseconds) are small compared to the IRAS aperture of approximately 300 arcseconds. The correlations at 12 and 25 pm are strikingly better than at 60 and 100 pm. It is possible that this effect may be due to the fact that the IRAS aperture at 60 and 100 pm (*» 5 x 3 arcminutes) is larger than at 12 and 25 pm (« 1x4 arcminutes), so that the size of the IRAS aperture at short wavelengths is more closely matched to the aperture for the By line measurements. However, good correlationsare n ot seen between the 12 and 25 pm fluxes and the S(l) line flux, suggesting that the wavelength dependence of the IRAS apertures has a minimal effect on the By correlations. It seems likely therefore that the correlations observed are real. The fact that the 12 and 25pm luminosities are correlated so well with Brackett line luminosities suggests that both the 12 and 25pm emission must be concentrated in the central regions of the galaxy encompassed by the smaller aperture used for the Brackett line observations. Conversely, the scatter in the correlation at 60 and 100pm suggests that the infrared emission may be more extended at these wavelengths. The following discussion concentrates three points. These are the relative sizes of the emitting regions for the By line and the infrared emission, whether the observed correlations can be understood in terms of star formation, and the origin of the 1 Opm emission. (i) Size of the emitting region The conclusion that the 12 an d 25 pm emission is centrally concentrated whereas the emission in the far infrared bands can be verified by looking at the correlation between the luminosity in the IRAS bands and the size of the galaxy, if the entire galaxy is within the IRAS aperture. The size of the galaxy can be estimated from the optical diameter D25. If the emission at 12 an d 25 pm is centrally concentrated with a scale size that is independent of D2 5, there should be no correlation between the luminosity in these bands and D25. However, if there is an extended component to the longer wavelength emission that depends on the size of the galaxy, the luminosity in these bands should correlate with D25. Thus correlation analysis can be used to determine the relation between luminosity and physical size. The results of this analysis are shown in Table 6.7. The galaxies with D25 > 300 arcseconds have been excluded from this analysis, since they are larger than the IRAS aperture. The12pm luminosity may be dependent on physical size (Pr «0.05), and the 2 5 pm flux essentially uncorrelated (Pr «0.5). O nly th e 100pm luminosity shows a significant correlation (Pr * 0.005) with physical size. This is consistent with there being an extended component at 100pm, while the mid-infrared emission is centrally concentrated. There is, however, direct evidence that the 10pm emission is centrally concentrated in many of the galaxies in this sample, which will be outlined below. Table 6.7 Correlation of D25 with IRAS fluxes starburstgalaxies 23 objects Band r t Pr 12pm 0.41 2.04 0.05 25pm 0.13 0.6 0.55 60 pm 0.34 1.65 0.11 100pm 0.56 3.1 0.005 212 Simple 10 pm mapping was obtained by Wright etal. (1988) for NGC 1614, NGC 2798 and NGC 6240. They found that NGC 1614 has 10pm emission that extends 4-5 arcseconds to the south and east of the nucleus, which corresponds to an em itting region o f **2Kpc. NGC 2798 w as found to b e extended on a scale o f *800pc at both 10 and 20pm. The 10pm emission has been mapped in NGC 3690 and IC 694 (collectively Arp 299) by Gehrz, Sramek and Weedman (1983), and the 20pm emission by Joy et al. (1989). This has been described in detail in Chapter 4; both sources are extended at 10 and 20pm, out to a distance of at least lKpc. As an alternative to maps at 10pm, an estimate of the spatial extent can be derived by comparing the 10pm flux in a small aperture from a ground based telescope to that obtained in the large 12pm IRAS aperture with suitable colour corrections. This technique is particularly useful for estimating the extent of the 10pm emission where more detailed measurements are not available, and is also used in statistical studies (Devereux 1987, Hill, Becklin and Wynn-Williams 1988; DePoy Becklin and Wynn-Williams 1988; Carico 1988). The ’compactness’ of the 10pm em ission is defined by RlO = 7gF (6.8) w here fjo is the small aperture 10pm flux, fj2 the IRAS flux and F the colour correction. Table 6.8 shows the compactness values for several of the sample galaxies, taken from DePoy (1987) Also shown is the approximate scale size for the emission ro, assuming that the spatial distribution is exponential, i.e. e " ^ . 213 Table 6.8 Galaxy Compactness Scale Size xlOO (Kpc) NGC 1614 85 0.2-0.5 NGC 4194 65 «0.2 NGC 4433 80 “ 0.2 NGC 5135 30 0.5-1 NGC 6240 95 «0.2 NGC 520 70 «0.2 NGC 2623 70 0.2-0.5 A rp 220 100 <0.2 These estimates of the scale sizes of the 1 Opm emitting regions in the sample galaxies therefore indicate that the mid- infrared emission is not uniformly distributed over the disks but has a central component «500 - 1000 pc in extent. The statistical analysis presented here shows that the By line flux is well correlated with the 12 and 25 pm emission, suggesting that the emission at these wavelengths has a central component which dominates the total flux. This suggestion is supported by simple mapping of the 10pm emission and estimates of the compactness, both of which show that the emission is typically 500 -1 OOOpc in extent. (ii) Origin of the mid-infrared emission It is widely assumed that the By recombination line flux arises primarily in HII regions. It is logical to assume that the correlation of the By line flux with the 12 and 25 pm fluxes arises because the mid-infrared emission is produced by thermal radiation from dust in and around HII regions. The aim of this section is to test this hypothesis. This will be achieved by comparing the integrated mid-infrared emission with the starburst luminosity predicted on the basis of the By flux. The integrated infrared luminosity was calculated by assuming that the 12 and 25 pm fluxes are characterised by a blackbody of emissivity u2. The blackbody 214 temperature was calculated for each galaxy from the ratio of the 12 and and 2 5 pm fluxes given by: F25 = /V25\ 5 e*12 - 1 /g g\ F12 'vi2' e** - 1 * ’ ' where F12 andF25 are the fluxes at 12 and 25pm, U25 and u 12 are the frequencies corresponding to 25 and 12pm, xi2“ hu^/kT and X25“ hu25/kT. T is the temperature and h and k have their usual meanings. Equation (1) can be rearranged to calculate the temperature, T -1 /F 2 5 v 12 T = t(V25-V 12 ) In (6.10) k iFi2vy. It has been assumed that ex >1. The mid-infrared flux for any galaxy can then be calculated by integrating the flux in a modified blackbody with a temperature inferred from equation (2). The ratio of the integrated flux in a blackbody to the flux at any given frequency u is given by, Fint (6.11) vFv where Fint is the integrated flux, F0 the flux at a frequency v, x = hu/kTand 00 te = 7 -V - 1 (6.12) n - l n 5 Fint can be expressed in terms of the flux at any frequency, F0, from equation (3), since x is known from T. The integrated infrared flux of a galaxy has been estimated from the observed 25pm flux and equation (3). The bolometric flux of the starburst has been estimated from the By flux and the starburst models described in Chapter 3. The integrated mid-infrared flux is plotted against the By flux for all of the starburst galaxies in the sample in Figure 6.9. The integrated flux calculated from the 215 By flux is shown for two starburst models; a ZAMS burst, and an evolved burst with a lower mass cut off of * 5 M o. Figure 6.9 shows that for most galaxies the observed mid-infrared flux is less than the bolometric flux predicted from the By flux. This is to be expected if some of the luminosity of the starburst is radiated at longer wavelengths. However the observed mid-infrared flux is typically at least 30-50% that of the bolometric flux, suggesting that a substantial fraction of the starburst luminosity is radiated at these wavelengths. The strong correlation between the By flux and the 12 and 25pm fluxes can therefore be understood in terms of a starburst model, where a large fraction of the luminosity of a centrally concentrated starburst is emitted by dust at mid-infrared wavelengths. 216 nertdmi-nrrde sin(nWnr)potdaantB lx(nWnr2) n W (in flux By against plotted r2) n W (in ission em id-infrared m Integrated Integrated Mid Infrared Flux iue 6.9 Figure 217 (iii) Origin of the lOjim emission The relation between the By flux (or radio free -free flux) and the 10 pm flux in galactic HU regions is given by (see Chapter 1) (6.13) where Fio is the 10pm flux density and Ffiy is the By flux. It is therefore possible to predictthe 10pm flux associated with the By emission assuming that the ratio of the By flux to 10pm flux is similar to that observed in galactic HE regions. The predicted 10pm flux can then be compared with the observed 10pm Auxin order to investigate whether the observed 10pm flux is consistent with star formation. Table 6.9 shows 10pmfluxes available from the literature for the starburst galaxies in this sample, and these are plotted against the extinction-corrected By flux in Figure 6.10. The nearby interacting system Arp 299 has been split up into four regions for this study; the 10pm sources A, B and C, and the extended 10pm emission around source A. These sources have been described in Chapter 4. The extinction correction applied to the By line flux for the merging galaxies (NGC 520, NGC 3256, NGC 6240, NGC 1614, NGC 4194, Arp 220 and NGC 6052) has been described in Chapter 3, and the extinction correction applied to Brackett lines for the components of Arp 299 has been described in Chapter 4. The extinction for He 2-10 was taken from Johansson (1987), and that for NGC 2798 from Keel et al. (1985). Figure 6.10 also shows the region of this plot occupied by galactic HII regions, described in Chapter 1. 218 Table 6.9 Galaxy 10|m Flux Aperture Reference (mJy) (arcseconds) NGC 520 310±15 5 DePoy(1987) NGC 2623 80±12 4 Wright etal.(1988) NGC 3256 17001150 15 Graham etal. NGC 6240 124115 4 Wright etal (1988) NGC 1614 480135 4 Wright etal (1988) NGC 4194 320145 Rieke and Low (1972) A rp 220 480 Soiferetal. (1984) NGC 2798 5201104 8 Wright et al (1988) H e 2-10 1700 11 Cohen & Barlow (1974) NGC 4631 >14.65 6 Cutri & McAlary (1985) A rp 299 Source A (nuc) 5201130 5 Gehrzetal. (1983) Source A (extended) See chapter 4 Source B 9001140 5 Gehrzetal. (1983) Source C 3251130 !5 Gehrzetal. (1983) 219 Figure 6.10 Ground based 10 pm flux plotted against By flux for starburst galaxies. It is striking that almost all of the galaxies shown in Figure 6.10 have 1 Opm/By ratios larger than that seen in galactic HD regions, a conclusion also reached by Ho, Beck and Turner (1989), Turner, Ho and Beck (1987) and Beck, Turner and Ho (1986) (see Chapter 1). This excess 10pm emission is usually parameterised by F io th e ra tio R = where Fio is the flux at 10pm (in Jy), and FsoHzis the flux at 220 5GHz (also in Jy). This ratio has a value ** 10 in HE regions ( see Chapter 1), whereas the galaxies in this sample have values of R m 20-100. Other radio and 10 pm observations of extragalactic star forming regions have confirmed that high values o f R are not uncommon. Rieke (1976) first noticed that R can be «20 for luminous extragalactic star forming regions. Other examples include in the starburst region of NGC 1068 where R *30 (Becklin, Wynn-Williams and Scoville 1985) and the nucleus of NGC 2903 where R a 70 (Wynn-WIlliams and Becklin 1985). In the following discussion reasons for the high R values will be considered. In galactic HH regions a high value of R can be obtained in compact HE regions where the emission measure is sufficiently high (> 109pc cm-6) that the plasma is optically thick at radio wavelengths, thus suppressing the radio flux (Habing and Israel 1979, Thompson 1987). However,this phase is short lived in HE regions (typically 105 years), so this is unlikely to be the explanation for the high R values found in extragalactic systems. Another possibility is that the excess 10pm emission is due to a population of small («10 A) dust grains (e.g. Wynn-Williams and Becklin 1985). These small grains have a very small heat capacity, and hence can be raised to a very high dust temperature (T*» 1000K) by the absorption of a single UV or optical photon (Sellgren 1984). There is spectroscopic evidence that small grains are present in the nuclei of starburst galaxies. The composition of these small grains is not known in detail, but they are associated with emission features at 3.3,8.7 and 11.3pm which may suggest that the primary constituent is some form of polycyclic aromatic hydrocarbon (PAH). These mid-infrared features are common in starburst galaxies (Roche and Aitken 1985, Aitken and Roche 1982), extragalactic HE regions (Phillips, Aitken and Roche 1984), but are rarely seen in active galaxies (Roche et al. 1984, Aitken and Roche 1985, Desert and Dennefield 1988). These features have been observed in several of the starburst galaxies in this sample; Phillips et al. (1984) found 8 and 11 pm emission features in the starburst nuclei He 2-10, and Joseph (private communication) has detections of the 3 pm feature in IC 694, NGC 1614 and NGC 2798. It seems plausible, then, that the excess 10pm emission in the sample being considered here is associated with a population of small grains in the starburst region. There is strong circumstantial evidence that there is an enhanced population of small grains in the starburst region, because large grains are easily broken up into small grains by the action of fast (50-500km-1) shocks from supemovae (McKee et al. 1987). There is direct evidence for fast shocks in the nuclei of the sample galaxies 221 from measurement of the [Fell] line, discussed in Chapter 3. Rieke and Low (1975) and Ho, Beck and Turner (1989) have both noted that the 10pm emission in starburst galaxies is spatially coincident with the synchrotron radio emission, which suggests that the small grain emission is associated with supemovae. There are two possible heating mechanisms in this scenario. One is that the interstellar radiation field is higher than average in the starburst region. Alternatively, the small dust grains may not be heated by interstellar photons but shock heated in supernova remnants (Ho et al. 1989). These two theories will be discussed below. There is some evidence that the first effect may contribute to the 12pm emission. Figure 6.11 shows the 12/25 vs 60/100 IRAS colour-colour diagram for the galaxies in the sample. The crosses are points taken from Boulanger and Perault (1988) and show the effect of increasing the interstellar radiation field incident on a dust cloud composed of both large and small grains, parameterised in terms of the equilibrium temperature of the big grains. The large grains in the cloud radiate in thermal equilibrium with the interstellar radiation field, while the small grains emit at much higher temperatures when they absorb a photon. The effect of increasing the radiation field is thus to increase the non-thermal emission from the small grains and at the same time increase the the temperature of the large grains. Itis clear that the galaxies fall approximately along the curve defined by increasing the radiation field. 222 log 60/100 Figure 6.11 ERAS colour-colour diagram showing the positions of the starburst galaxies (filled squares). The crosses were taken from Boulanger and Perault (1988) and show the effect of increasing the intensity of the radiation field. The solid line is a quadratic fit to these points. The suggestion that shock heating of the small grains was made by Ho, Beck and Turner (1989) and was motivated by the observation that the 10pm emission is spatially coincident with the non -thermal radio emission. If this effect is important it is difficult to see why the 12pm IRAS luminosity is well correlated with the B y luminosity, because a large contribution to the 10pm flux from shock heated grains would introduce scatter into this relation. The ERAS colours are consistent with increasing the radiation field. My conclusion is therefore that the small grains are heated by photons from the young stellar population, and that the 10pm emission is enhanced in the central region by a combination of an excess population of small grains and increased radiation field. 223 6.3.3 S(l) Line Correlations In contrast to the Brackett line correlations, the S(l) line flux is notwell correlated with the 12 or 25 pm flux, and is only convincingly correlated at 100 pm. This lack of correlation is somewhat surprising, since early observations of H2 emission in starburst galaxies suggested that the S( 1) line was excited by outflows from young stellar objects (e.g Fischer et al. 1983, Joseph Wright and Wade 1984). These results argue against the S(l) line emission being associated directly with star formation, either shock excited by outflows or fluorescently excited by UV photons from the young stars. They also do not support the suggestion by Ho, Beck and Turner (1989) that the lOpmfluxis dominated by shock heated small grains, since this theory would predict a relation between the line emission from shocked gas and the 12pm flux. In section 6.5.1 it was suggested that the 100pm flux had an extended component in starburst galaxies. The correlation between the S(l) line and the 100 pm flux therefore suggests that the shocked gas is also extended in starburst galaxies. The lack of correlation between the S(l) line flux and the other IRAS bands is also explained if the S( 1) line has an extended component It is possible that this result reflects the number of interacting and merging galaxies in the sample, and is consistent with the result obtained in Chapter 3 that the shocked gas is more extended than the infrared emission. 6.4 Active Galaxies In section 6.4.11 present the results of the partial correlation analysis for the active galaxies, which are discussed in section 6.4.2. 6.4.1 Results The results of the partial correlation analysis for the active galaxies are shown in Tables 6.10 and 6.11. Tables 6.10 and 6.11 show r, t, and Pr for By and S(l) respectively, correlated against IRAS bands. 224 Table 6.10 Brackett line correlations for active galaxies 12 objects Band r t Pr 12pm 0.78 3.85 0.004 25 pm 0.84 4.8 0.0009 60pm 0.74 3.36 0.008 100pm 0.53 1.92 0.086 Table 6.11 S( 1) line correlations for active galaxies 13 objects Band r t Pr 12pm 0.57 2.32 0.04 25 pm 0.28 0.98 0.34 60pm 0.08 0.26 0.79 100pm 0.5 1.96 0.075 The correlations of the infrared lines and IRAS bands for active galaxies are significantly different from the correlations obtained with starburst galaxies. The Brackett lines are well correlated with all of the IRAS bands for active galaxies, in contrast to the starburst galaxies where only the 12 and 25pm emission is significantly correlated. Correlations are also seen between the 12 and 100pm fluxes and the S(l) line in Seyferts. 225 6.4.2 Brackett Line Correlations The By flux observed in this sample of active galaxies may arise in HE regions or may be associated with gas photoionised by the active nucleus. This question was discussed at some length with reference to the well known Seyfert Galaxies NGC 1068 and NGC 4151 in Chapter 5, and it was concluded that in these two galaxies the By flux is dominated by gas in the narrow line region in NGC 1068, and in the broad line region of NGC 4151. It will therefore be assumed that the By line in the active galaxies discussed here is associated with the active nucleus. The good correlations seen between the By flux and the IRAS bands suggest that the IRAS emission is centrally concentrated in the region encompassed by the aperture used for the Brackett line observations, and dominated by the active nucleus. This indicates that the infrared emission is dominated by dust heated in the narrow/broad line regions of these galaxies. There is, however, significantly more scatter in the correlation at 100pmthan at the other wavelengths, suggesting that the 100pm emission has an extended component. This suggestion can be tested once again by investigating the correlation between the IRAS bands and the optical diameter D25. The result of this analysis is shown in Table 6.12. Table 6.12 Correlation of D25 with IRAS fluxes Activegalaxies 14 objects B a n d r t Pr 12pm 0.4 1.33 0.21 25 pm 0.39 1.26 0.23 60pm 0.5 1.77 0.1 100pm 0.62 2.38 0.04 The only convincing correlation is seen between the 100 pm flux and D25, which supports the suggestion that the scatter in the correlation with the By flux is due to an extended component 226 The extended component at 100 pm may be due to a drcumnuclear starburst, a phenomenon that is frequently seen in Seyfert galaxies (e.g. Keel 1984, Wilson 1988). The archetypal example of a circumnuclear starburst is NGC 1068, which was discussed in detail in Chapter 5. Here the mid-infrared emission (As 30pm) is dominated by dust heated by the active nucleus in the narrow line region, while the far infrared emission contains an extended component associated with the starburst Other galaxies in this sample are known to have circumnuclear star formation, for example NGC 7469 (Cutri et al. 1984, Wilson et al. 1986), NGC 7582 (Morris et al. 1985) and NGC 1365 (Phillips et al. 1983). A number of objects have been classified as 'composite ' by Veron and Veron-Cetty (1986), including NGC 1808, NGC 6221 and NGC 613. The correlation analysis described here can be compared other studies of Seyfert galaxies by Rodriguez Espinosa, Rudy and Jones (1987), and Dultzein- Hacyan, Moles and Masegosa (1988). Dultzein-Hacyan, Moles and Masegosa (1988) found good correlations between the HP line luminosity and 12 and 25 pm luminosities for a sample of Seyfert 2 galaxies, but did not report significant correlations between the Hp luminosity and the 60 and 100 pm luminosities. They did not however, perform any quantitative statistical tests, so a detailed comparison with the results presented here is not possible. Rodriguez Espinosa, Rudy and Jones (1987) carried out a statistical analysis of the relation between the optical-ultraviolet 3pm and 10pm fluxes and the IRAS 25, 60 and 100pm fluxes. They found essentially no correlation between the IRAS fluxes and the optical-ultraviolet fluxes. Significant correlations were found between the 25 pm fluxes and the 3 and 10 pm fluxes but the significance of the correlations between the 3 and 10 pm fluxes and the 60 and 100 pm fluxes were reduced. Their interpretation of these results was that the far infrared (A^30 pm) was dominated by circumnuclear star formation. The main conclusion of Rodriguez Espinosa, Rudy and Jones (1987) - that there is evidence for circumnuclear star formation in Seyfert galaxies - does not differ significantly from the conclusions outlined above. However, the result that the optical-ultraviolet fluxes are uncorrelated with the infrared fluxes disagrees with the Brackett line correlations presented here. There are two possible reasons for this discrepancy. The first is that the Seyferts considered by Rodriguez Espinosa, Rudy and Jones (1987) are optically selected. In contrast, the Seyferts in this study are infrared objects, and there is only one object, NGC 7469, in both samples. There may be systematic differences between optically-selected and infrared-selected Seyferts. The second is that Rodriguez Espinosa, Rudy and Jones (1987) have made 227 no effort to correct for the optical-ultraviolet fluxes for extinction, and that this has introduced scatter into the relationships observed. Carletonetal. (1987) concluded that there is evidence for significant amounts of reddening (Av » 0.2 -1.5) in a sample of Seyfert 1 galaxies, which suggests that the second effect may be im portant Further analysis is needed to discriminate between these two possibilities. 6.4.3 S(l) Line Correlations The S( 1) line shows weak correlations with the 12 and 100pm flux, and is uncorrelated with the 60 and ■» ;pm fluxes. This is similar to the results obtained for the starburst galaxies, where the S(l) line flux was only weakly correlated with the 100pm flux. This suggests that the H2 emission is not directly related to the nuclear activity, and argues against the idea that fluorescent emission dominates the flux in Seyfert galaxies. 6.5 Summary The correlations of the infrared lines and the IRAS bands for a sample of starburst galaxies show that the By line flux is well correlated with the 12 and 25pm fluxes, that the emission at these wavelengths is centrally concentrated (with an exponential scale size **200 -1000pc) and that the correlation can be explained if the mid-infrared emission is due to thermal reradiation from hot dust in HH regions. The 10pm flux is larger than that predicted from simple star formation models and the excess emission is probably due to non thermal emission from small grains in the starburst region. In contrast, the S(l) line is only weakly correlated with the 100pm flux, suggesting that the S( 1) line is not directly associated with the starburst activity. The correlation of the S(l) line with the 100pm flux may suggest that the S( 1) line has an extended component, which probably reflects the number of merging galaxies in the sample. This supports the results in chapter 3 that the shocked gas is more extended than the infrared emission. The correlations of the infrared lines and the IRAS bands for the active galaxies show dramatic differences from the correlations found for the starburst galaxies. The By flux is well correlated with all of the IRAS bands, suggesting that in active galaxies the bulk of the infrared luminosity is associated with the active nucleus. The correlation of the By flux with the 100pm flux has rather more scatter than the other bands. This may be due to an extended component associated with a 228 circumnuclear starburst The S( 1) line is not well correlated with any of the IRAS bands suggesting that it is not powered directly by the active nucleus. 229 Chapter 7 Summary, Conclusions and Suggestions for Further Work 7.1 Summary There are two major themes in this thesis. The first is the use of Brackett line spectroscopy to characterise extragalactic star formation, and the second is the use of multiaperture spectroscopy and statistical techniques to investigate the spatial extent of infrared lines. A simple starburst model was used in Chapter 3 to investigate the star formation in merging galaxies. It was found that that values of Q (the ratio of the star formation rate inferred from the By flux to that inferred from the infrared luminosity) for most of the mergers is consistent with Zero Age Main Sequence star formation or an evolved starburst with a high (5 Mo) lower mass cut off. In contrast to these findings, multiaperture spectroscopy of the Seyfert galaxy NGC 1068 suggests that the By line is associated with narrow line gas photoionized by the active nucleus. The By line flux was found to vary on timescales of months in NGC 4151, showing that the recombination line flux in this Seyfert is dominated by gas in the broad line region. The conclusion that the recombination line flux in Seyferts is dominated by the active nucleus provides a physical basis for the suggestion in Chapter 3 that Seyfert galaxies have systematically larger infrared luminosities for a given By luminosity than do merging galaxies. There is evidence from the statistical analysis in Chapter 6 that the spatial distribution of the infrared emission is systematically different in starburst and active galaxies. The 12 and 25 pm fluxes are significantly correlated with the By fluxes in starburst galaxies, butthe the 60 and 100 pm fluxes are not significantly correlated the Brackett line fluxes. This suggests that the 12 and 25 pm emission are concentrated within the aperture used for the By observations. In contrast, the fluxes in all 4 ERAS bands are well correlated with the By flux in a sample of Seyfert galaxies. This suggests that the infrared luminosity in Seyfert galaxies is dominated by the active nucleus. The spatial extent of the shocked gas was investigated by multiaperture spectroscopy of the [Fell] and S(l) lines in three merging galaxies in Chapter 3. It 230 was found that the the S(l) and [Fell] had exponential scale sizes of approximately 1.5 kpc in the all the galaxies. This is in contrast to the exponential scale size of the infrared emission, which is about 0.5 kpc. This strongly suggests that the extended shocked gas is not directly associated with the source of the infrared luminosity, and is most likely powered by the merging process. These results are consistent with those obtained in Chapter 6, where it was found that the S(l) line in a sample of starburst/interacting galaxies is more extended than the infrared emission. There are two further major conclusions from multiaperture S(l) line spectroscopy. Spectroscopy of the S( 1) line in NGC 1068 shows that this line is centrally concentrated, with some evidence for an extended component Comparison of the large aperture spectra in Chapter 5 with the 4 arcsecond spectrum of Hall et al. (1981) shows that the 2-1 S( 1) line is rather more prominent in the large aperture spectra. The obvious interpretation of this is that the extended S( 1) line is fluorescently excited, while the nuclear S( 1) line is thermally excited. This conclusion, however, depends critically on the location of the continuum in the larger apertures, which is difficult to determine in low resolution spectra. Simple S(l) line mapping in the interacting system Arp 299 and multiaperture spectroscopy in Arp 220 show that there is a discrepancy between the spatial distribution of the molecular hydrogen traced out by S(l) line spectroscopy and aperture synthesis maps of CO. These discrepancies underlie the fact that aperture synthesis maping in CO is insensitive to low surface brightness extended emission. 7.2 Matters Arising - UV Fluorescence in Spiral Galaxies? In Section 4.2.3 it was stressed that the position of the continuum in the K window spectra of NGC 1068 was determined partly because the points between the By line and the 2-1 S( 1) line were almost certainly below the continuum, because of the presence of the Nal stellar absorption line at 2.21 pm. This may have some bearing on the recent claim of Puxley et al. (1989) to have detected H2 emission that is dominated by fluorescent excitation. Figure 7.1 shows their spectrum of NGC 1808, which incidentally contains a Seyfert nucleus. The position of the Nal absorption feature is indicated. It is clear that Puxley et al have located the continuum at the lowest possible position, and if there is absorption at 2.21 pm then a higher continuum should be drawn.. For example, if the absorption equivalent width of the Nal feature is a third of the emission equivalent width of the By line (as it is in NGC 1068) then the continuum is 231 rather higher in this region than indicated by Puxley etal.. In this case, the flux in the 2- 1S(1) line is reduced to less than 25% of the 1-0S(1) line flux, more consistent with shock excitation than fluorescent excitation. It is perfectly possible that Puxley et al. have detected fluorescently excited H2 emission, but their claim that the H2 emission is dominated by fluorescently excited gas is dependant on there being essentially no Na absorption. It is more probable that there is a shocked component to the H2 emission. In the case of NGC 1808 the morphology may be similar to NGC 1068; the fluorescently excited gas surrounding a shock excited nuclear component 232 b. NGC 1808 CMI E OJ CD t- cn OJ J £ 0 . V=1-0 S(1) 2.2 2.3 2.4 Wavelength (pm) Figure 7.1 Spectrum of NGC 1808, taken from Puxley et al. 1989 233 In their analysis, Puxley et al calculate the number of photons available for fluorescent excitation on the basis of the B y flux. They calculate Ruv (the ratio of the number of photons required to produce the observed 1 -0S(1) line flux to the number available -see Chapter 1) for each galaxy in the case where the continuum is produced by a burst of star formation. The value of Ruv is dependent on the upper mass cut off of the star burst, and values of Ruv > 1 are only found when the upper mass cut off ^ 30M . The maximum value of Ruv is 3.14, and for NGC 4536 Ruv = 1.18. Thus although fluorescent excitation is energetically possible, these values of Ruv are uncomfortably close to the minimum value of 1. This constraint is weaker, however, if not all of the 1-0 S( 1) line flux is fluorescently excited. Puxley et al. also speculate on why the H2 emission in spiral galaxies is different from that observed in other extragalactic systems where the emission is shock excited. Assuming that a large fraction of the 1-0S(1) line in these spirals is shock excited, the difference does not seem so m arked. 234 7.3 Suggestions for Further Work In this section I describe some observations that would help clarify some of the major unresolved issues in this thesis. I will concentrate on the potential of line imaging and further spectroscopic work. 7.3.1 Line Imaging Until very recently, infrared astronomers essentially relied on single detector instruments (with the exception of one dimensional arrays used in spectroscopy), lim ited spatial information was gained by multiaperture photometry, and crude imaging was achieved by raster scanning of the single element In 1986 the Santa Barbara Research Corporation developed a two dimensional InSb photovoltaic focal plane array specifically designed for infrared astronomy; in particular a high quantum efficiency (^80%), a high filling factor (^90%), low noise from detectors and readout device, and a high dynamic range. The readout device is a switched FETmultiplexer which can simultaneously read out the 3596 elements of the array. The resolution and field of view of a camera depends on the f-ratio of the telescope and any optics, but the resolution is typically 1 arcsecond per pixel and the field of view an arcminute. The UKIRT infrared camera, for example, has three scales available; 0.6,1.2 and 2.4 arcseconds per pixel This provides the infrared astronomer with an imaging device (in the 1-5pm region) thatis comparable to optical CCD arrays. There are a number of questions raised in this thesis that could be resolved by line imaging. For example, Figure 4.3 tentatively suggests that the spatial distribution of the By line in NGC 1068 is coincident with the high excitation [OIH] line, which could be easily verified by direct imaging. Imaging would also show whether the By line shows any evidence for asymmetry, particularly along the jet axis. The feasibility of such a study can be demonstrated by calculating the approximate on source- integration time required to obtain a 10a detection of the extended By line in NGC 1068. The extended flux is *8xlO-17W nr2, and assuming this is uniformly distributed between radii of 2 and 10 arcseconds the surface brightness is ®2.7xl0-19 Wm'^arcsecond2. Using the specifications of the Kitt Peak National Observatory infrared array (KPNO Users Manual) and assuming that the pixel size of the array is «1 arcsecond, that the exposure time is 40 seconds and that the array is in use on a 4m telescope, the total on-source integration time fora 10a detection per pixel is about 70 minutes. In order for a calibrated line map to be produced, a continuum image will have to be subtracted. Allowing for observations of standards, it is feasible to 235 produce a By line map in half anight This is comparable to the time required to produce the spectra presented in Chapter 4. Infrared line and continuum imaging will be used to investigate the spatial distribution of gas and stars in galaxies in more detail than ever before. High resolution Brackett line imaging can be combined with existing H a images to produce extinction maps. Brackett line and 2pm continuum images can be used to compare the spatial distribution of stars and ionised gas. Detailed knowledge of the spatial extent of the 1 -0S(1) line can be compared with aperture synthesis maps in CO, and the hypotheses that the shocked emission in starburst galaxies is primarily due to the interaction of supemovae with the interstellar medium can be tested by comparing the distribution of the 1 -0S( 1) line and [Fell] lines with the non-thermal radio emission. The suggestion that the shocked gas in mergers is more extended than in Seyferts is readily tested by line imaging. In Chapter 5 it was suggested that the excess 12pm flux was due to non-equilibrium emission from small grains that were photoionized; if this is the case the By line might be expected to follow the 10pm emission. This project will have to wait until amid infrared array has been developed. 7.3.2 Future Spectroscopic Observations In Chapter 2 it was pointed out that the UKT 9 CVF spectrometer was rather more sensitive than the CGS 2 grating spectrometer. The sensitivity of CGS 2 has recently been improved, so that is now possible to obtain high resolution extragalactic spectra. This is crucialfor accurately determining the continuum, to enable quantitative modelling of the molecular hydrogen vibration-rotation lines. The most immediate question that falls in this category is to discover whether the extended H2 emission in NGC 1068 is fluorescently excited; this requires one or more off nuclear, small (5 arcsecond) aperture, high resolution (X/5X = 300) spectra to fully resolve the Nal absorption line and the 2-1 S( 1) line. Higher resolution spectra of the galaxies observed by Puxley et al. will help determine the relative contribution of shocked and fluorescently excited gas. In the longer term, the two dimensional arrays described in the previous section can be used in long slit spectroscopy, where one dimension is spatial and the other dispersive. This would be ideal for determining the spatial distribution of the shock excited and fluorescently excited 1 -0S( 1) line in different galaxy types. 236 Low resolution K- band spectra are particularly useful for the type of statistical analysis described in Chapter 5, even if the results presented in Chapter 4 suggest that they should used with caution for detailed modeling of the H2 lines. The main shortcoming of the analysis described in Chapter 5 is the wide range of galaxy types in the sample. An obvious extension of this work is to obtain By fluxes (and preferably B a fluxes as well to allow an estimate of the extinction) for a well defined sample of nuclear starburst galaxies, such as M arkarian galaxies with HA region nuclei. Partial correlation analysis can be used to test whether the relationship between the ionizing continuum and the 12pm IRAS flux holds for a better defined sample. The large scale ’energy budget’ of galaxies can be better understood by looking for correlations between the ionizing luminosity and the luminosity at other wavelengths such as optical, X-ray and radio. The relationship between star formation and the interstellar medium can be better understood by looking for correlations between ionizing luminosity and CO mass, and HI mass. In Chapter 5 it was suggested that there is an excess of small grains in the starburst region of galaxies because larger grains are broken up by the action of fast shocks from supemovae. This hypothesis can be tested by long slit spectroscopy of the 10 pm feature in galaxies to see whether this feature is enhanced in the central starburst area. The suggestion that the small grains are produced in shocks can be tested by combining the results of long slit 10 pm spectroscopy with line imaging of shocked H2 and [Fell]. 7.4 Concluding remarks. The major objectives of the work in this thesis is explore ways in which infrared spectroscopy can be used to understand physical processes in galaxies. The major conclusions are that the infrared luminosity in merging galaxies can be understood in terms of star formation, the shocked lines in merging galaxies are powered by the interaction, and the discovery of extended fluorescent molecular hydrogen in NGC 1068. These results demonstrate the utility of infrared spectroscopy as a tool for extragalactic research. The full potential of infrared spectroscopy will be realised in the next few years with the advent of sensitive infrared spectrometers and imaging devices. 237 Appendix 1 Starburst Parameters The simple analytical starburst model described in Chapter 3 approximated the stellar quantities as power laws over various mass ranges. In this appendix I give full details of the derived quantities. Throughout this appendix the stellar luminosities and masses are measured in solar luminosities and masses. Al.l The Initial Mass Function The initial mass function for the starburst models was taken from Scalo (1986). Figure Al.l shows t|i(M) where ¥(M )dM is the number of stars kp-2 formed in the mass range M-M+dM plotted as a function of mass. The corresponding fits are also shown. The parameterisation adopted is given by: 2.5 X 1.5- X x £ 0.5 - $ - 0.5 - -15 - - 2.5 - 1.5 -0 5 0.5 1.5 LogM Figure Al.l Initial mass function from Scalo (1986), with power law fits 238 A 1.2 Mass-Luminosity Law The mass-luminosity law adopted was taken from Tinsley (1980) for stellar masses less th an 50 Mo. and from Maeder (1980) for masses > 50 Mo* T he m odels of Maeder (1980) include the effects of convective overshoot for large mass stars. Figure A 1.2 shows the data from Tinsley (1980), plotted as crosses and the data from Maeder (1980) plotted as triangles. The lines indicate the fits derived. They are tfe)" 048 lm ) 16 M thyu & 6 5 < m < s o m g tfe )- 1258® 1'5 m > 5 o m ° The parameters derived for large M > 50 M o were obtained from McDowell (private communication). 239 Figure A1.2 Stellar mass plotted as a function of luminosity A 1.3 Mass Main-Sequence Lifetime The relationship between mass and main sequence lifetimes were taken from Scalo (1986). Figure A1.3 shows the mass plotted against main sequence lifetime taken from Scalo (1986), along with the derived fits. M -3.0 II / \ ■ 104 | M < 5 M o f I m o J B 309Ci r— v 1-9 5 < M < 1 0 M o l E I m o J ( I m ) / M \ -1.17 - 4 2 7 10 < M < 2 0 M q (yearsj I m o J (TiojL'S f M s-0.83 - 1 5 5 2 0 < M < 5 0 M o (yearsj i M o J ( Tms’s r M \-0.55 ■55 r- M > 5 0 M o (yearsj M o ) 240 11 Log m Figure A1.3 Main sequence lifetime as a function of mass A 1.4 Mass - Ionising Photon Flux The number of ionising photons as a function of stellar luminosity has been tabulated by Panagia (1973). These values were used to obtain a relationship between the luminosity and emission of ionising photons in photons s*1. These were used in conjunction with the mass-luminosity relationships derived in section A 1.2 to obtain mass-ionising photon flux equations. The number of ionising photons vs.luminosity taken form Panagia (1973) is shown in Figure A1.4, along with the derived fits ( S ^ 7 i) - 3'4s,,0"(-fe)U1 L>s4l'1°*1® fcaS n) - ‘•M« “ * (“h f ‘1 - M * ">‘I® The mass-ionising photon flux relationships are (S ^ S T i)'3'47lloI’(M5),'“ 5 * M < 3 4 M o M < M < S O M o 241 iDhotons s-ij s-ij iDhotons /_J5L LogN Ionizing photon flux versus luminosity versus flux photon Ionizing x _ £ n -- 1.95 iue A1.4 Figure xlO46 f - J r i181 M > 50 M Mo i181 r J - xlO46 f MoJ l References Aitken, D.A., and Roche, P.F., 1985, M.N.R.A.S., 213, 777 Aitken, D.K., Roche, P.F., Allen, M.C., and Phillips, M.M., 1982, M.N.R.A.S . , 199, 31p Ailken, D.K., Roche, P.F., and Phillips, M.M., 1981, M.N.R.A.S, 196, lOlp Allen, P.A., James, T.J., Hyland, A.R., 1985, AstrophysJ., 291, 280 Arp, H., 1966, A das o f Peculiar Galaxies, Pub Californian Institute Technolo gy Arums, L.,Heckman, T., and Miley, G., 1987, A stro n o J . 94, 831 Atherton, P.D., Reay, N.K., Taylor, K., 1985, M.N.R.A.S.,2 1 6 ,15p Baldwin, J.A., Wilson, A.S. Whittle, M.,1987, Astrophys. J., 319, 84, Balick, B., and Heckman, T., 1985, Astron. and Astrophys., 90,197 Barford, N.C., 1985, Experimental Measurements: Precision, Error and Truth, John Wily & sons Barr, P., and Mushotzky, R.F., 1986, N a tu re 320, 421 Barrow, G.M.,1962, in Introduction toM olecu;arSpectroscopy, McGraw Hill Beck, S.C., and Beckwith, S., 1984, MNRAS, 207, 671 Beck, S.C., Beckwith, S., and Gatley, I., 1984, AstrophysJ, 2 7 9 , 563 Beck, S.C., Turner, J.L. and Ho, P.T.P., AstrophysJ.. 309, 76 Becklin, E.E., Wynn-Williams, C.G., and Scoville, N.Z., 1985, AstmphysJ., 2 9 7 ,6 0 7 Beckwith, S., Persson, S.E., Gatley, I., 1978, AstrophysJ.LetL , 219, L.33. Beckwith, S., Persson, S.E., Neugebauer, G., and Becklin, E.E., AstrophysJ., 223, 464 Black, J.H. and van Dishoeck, E.F., 1987, AstrophysJ\, 322, 412 Black, J.H., and Dalgamo, A., 1976, A s t r o p h y s J203, 132. Boulanger, F., andPerault, M., 1988, AstrophysJ., 330, 964 Bowers, R.L., and Deeming, T., 1984, in Astrophysics II: Interstellar Matterand G alaxies, Jones and Bartlett Brocklehurst, M., 1971, M.N.R.A.S., 153, 471 Burton, M.G., Geballe, T.R., and Brand, P.W.J.L., 1988, MNRAS, 231, 617 Campkins, H., Rieke, G.H., and Lebofsky, M.J., A stro n J., 90, 896 Canto, J.1985, NearbyMolecularClouds, ed. Serra, G., Springer Verlag. 243 Carico, D.P., Sanders, D.B., Soifer, B.T., Elias, J.H., Matthews, K., and Neugabauer, G., 1988, A stro n J., 95, 356 Carleton, N.P., Elvis, M., Fabbiano, G., Willner, S.P., Lawrence, A., and Ward, M., 1987, AstrophysJ., 318, 595 Casoli, F., Combes, F., Dupraz, C., Germ, M., Encrenaz, P. and Salez, M., 1988, Astron. & Astrophys., 192, L.17. Cohen, M., and Barlow, M.J., 1974, A s tr o p h y s J., 193, 49 Cramer, H.A., 1946, in M athematical Methods o f Statistics, Princeton University Press Cutri, R.M., Aitken, D.K., Jones, B., Merrill, K.M., Peutter, R.C., Roche, P.F., Rudy, R.J., Russell, R.W., Soifer, B.T., and Willner, S.P., 1981, AstrophysJ., 245, 818 Cutri, R.M., and McAlary, C.W., 1985, AstrophysJ., 296, 90 Cutri, R.M., Rudy, R.J., Rieke, G.H., Tokunaga, A.T., and Willner, S.P., 1984, 280, 521 DeGrijp, M.H.K. Miley, G.K., Lub, J., and de Jong, T., 1985, Nature, 314, 240 Delbouille, L., Roland,G., Brault, J., Atlas of the Solar Spectrum, University of Arizona Press DePoy, D.L., 1987, Ph.D thesis, University of Hawaii. DePoy, D.L., Becklin, E.E., andGeballe, T.R.,1987, AstrophysJ., 316, L63 DePoy, D.L., Becklin, E.E., and Wynn-Williams, C.G., 1986, AstrophysJ., 307,116 Devereux, N., 1987, AstrophysJ., 323, 91 Dinerstein, H.L., Lester, D.F., Carr, J.S., Harvey, P.M, 1988, Astrophys. J ., 327, L27 Doyon, R., and Nadeau, D., 1988, AstrophysJ., 334, 883 Drain, D.T., Roberge, W.G., and Dalgamo, A., 1983, AstrophysJ., 264, 485 Emerson, J.P., Jennings, R.E., and Moorwood, A.F.M., 1973, AstrophysJ., 184, 401 Field, G.B., Somerville, W.B., and Dressier, K., 1966, Ann Revs Astron and A stro p h ys, 4, 207 Fischer, J., Geballe, T.R., Smith, H.A., Simon, M., and Storey, J.W.V., 1987, AstrophysJ., 320, 667 Fischer, J., Sanders, D.B., Simon, M., Solomon, P.M., 1985, AstrophysJ. 293, 508. Fischer, J., Simon, M., Benson, J., and Solomon, P.M., 1983, AstrophysJ, 273, L27 244 Hollenbach, D.J., and MCKee, C.F., 1980, A stro p h ysJ., 241, L47 Hutchings, J.B., and Neff, S.G., 1987, A s tr o n J ., 93, 14 Hutchings, J.B., and Neff, S.G., 1988, A s tr o n J .96, 1575 Israel, F.P., 1988, in Millimetre and Submillimetre Astronomy, ed. Wolstencroft, R.D., and Burton, W.B., Kluwer Academic Publishers. Issacman, R., 1984, Astron.andAstrophys, 130, 151. Johansson, L., 1987, Astron andAstrophys, 182, 179 Johnson, H.L. and Mendez, M.E., 1970, A stro n . J., 75, 785. Joseph, R.D. and Wright, G.S., 1985, M.N.R.A.S., 214, 87. Joseph, R.D., Miekle, W.P.S., Robertson, N.A., and Wright, G.S., 1984 M.N.R.A.S 209, 111 Joseph, R.D., Wright, G.S. and Wade, R., 1984, N a tu rq 311, 132. Joy, M., and Lester, D.F., 1988, AstmphysJ., 331, 145 Joy, M., Lester, D.F., Haney, P.M. and Frueh, M., 1988, A s tr o p h y s J., 307, 110. Joy, M., Lester, D.F., Harvey, P.M., Telesco, C.M., Decher, R., Rickard, L.J., andBushouse, H., 1989, Astrophys. J., 339, 100 Joy, M., Lester, D.F., Harvey, P.M., Telesco, C.M., Decher, R., Rickard, L.J., and Bushouse, H., 1989, A s tr o p h y s J., 339, 100 Kawara, K., Nishida, M and Gregory, B., 1987, AstrophysJ.Lett, 321, L.35. Kawara, K., Nishida, M., and Taniguchi, Y., 1988, AstrophysJ., 328, L44 Keel, W.C., 1984, A s tr o p h y s J., 282, 75. Keel, W.C., Kennicutt, R.C., Hummel, E. and Van Der Hulst, J.M., 1985, A s t r o n J90, 708. Kimiaki, K., Nishida, M., and Gregory, B., 1987, A s t r o p h y s J321, L.35. Kormendy, J., 1985, AstrophysJ.Lett., 292, L.9. Krolik, J.H., McKee, C.F., and Tarter, C.B., 1981, AstrophysJ., 249, 422 Kuiper, G.P., 1963, Comm Lunar and Planetary Lab, 1,1 Kuiper, G.P., AstrophysJ., 1947, 106, 251 Kuiper, G.P., Wilson, W. and Cashman, R.J.,1947, Astrophys.J., 106, 243 Kwan, J., 1977, Astrophys. J. 216,713 Kwan, J., and Krolik, J.H., 1981, AstrophysJ., 250, 478 246 Lang, K.R., 1978, in AstrophysicalFormulae, Springer-Verlag Larson, H.P, 1980, Ann Revs Astron and Astrophys, 18, 43 Lawrence, A., Rowan-Robinson, M.M. Leech, K., Jones, D.H.P., and Wall, J.V., 1989, M.N.RA.S., 240, 329 Lawrence, A., Ward, M., Elvis, M., Fabbiano, G., Whinner, S.P., Carleton, N.P. andLongmore, A., 1985, Astivphys.J219, 117. Lester, D.F., Harvey, P.M., and Carr, J., 1988, AstrophysJ., 329, 641 Lester, D.F., Joy, M., Harvey, P.M., Ellis, H.B., and Parmer, P.S., 1987, Astrophys. J.,321,755 Lester, D.F., Joy, M., Harvey, P.M., Ellis, H.B., and Parmer, P.S., 1987, Astrophys. J., 321,755 London, R., McCray, R., and Chu, S-I, Astrophys.J., 217, 442 Longmore, A.J., Robson, E.I. and Jameson, R.F., 1986, M.N.R.A.S., 221, 589. Lonsdale, C.J., Helou, G., Good, J.C., and Rice, W., 1985, CataloguedGalaxies and Quasars Observed in die IRAS Survey Jet Propolusion Lab Lonsdale, C.J., Persson, S.E. and Matthews, K., 1984, AstrophysJ., 287, 95. Low, F.J., and Rieke, G.H., 1974, in M ethods o f Experimental Physics, Volume 12, P a r t A -Astrophysics, ed Carleton, N., Academic Press Lowe, R.P., Moorehead, J.M., and Wehlau, W.H., 1979, A s t r o p h y s J228, 191. Lynden-Bell, D., 1969, N ature; 223, 690 Macklin, J.T., 1982, M.N.R.A.S., 199, 1119 Maeder, A., 1980, Astron & Astrophys, 92, 101 Malkan, M.A., Alloin, D., and Shore, S., 1987, Exploring the World with thelUE Satellite; Kondo, Y., Wamsteker, W., Boggess, A., Grewing, M., De Jager, C., Lane, A.L., Lrnsky, J.L., and Wilson, R., D. Reidel, Dordrecht Matthews, W.G., and Capriotti, E.R., 1985, Astrophysics o f Active Galaxies and Quasi-StellarObjects, Ed J.S. Miller, Oxford University Press McAlaiy, C.W., and McLaren, R.A., 1981, AstrophysJ., 250, 98 McAlay, C.W., Rieke, G.H., Lebofsky, M.j. and Stoke, J.T., 1986, A stro p h y s. ., 301, 105. McCarthy, D.W., Low, F.J., Kleinmann, S.G., and Gillett, F.G., 1982, Astrophys.J., 257, L7 McCarthy, P.J., Heckman, T. and Van Breugel, W., 1987, A stron J., 92, 2. 247 McKee, C.F., Chemoff, D.F. and Hollenbach, D J., 1984, Oalacticand ExtragalacticInfraredSpectroscopy., ed Kessler, M.F., and Phillips, J.P., Dordrecht, Reidel. McKee, C.F., Hollenbach, D.J., Seab, C.G., and Helens, A.G.G.M., 1987, Astrophys J., 318, 674 Merrill, K.M., and Ridgway, S.T., 1979, Ann Revs Astrxm and Astrophys, 17, 9 Meurs, E.J.A., and Wilson, A.S., 1984, AstronandAstrophys, 136, 206 Moorwood, A.F.M., and Oliva, E., 1988, Astron & Astrophys., 2 0 3 , 278 Morris, S.L., Ward, M.J., Whittle, M Wilson, A.S., and Taylor, K., 1985, M.N.R.A.S., 216, 193 Myers, S.T. and Scoville, N.Z., 1987, AstrophysJ.Lett, 312, L.39. Nakagawa, T., Nagata, T., Geballe, T.R., Okuda, H., Shibai, H.and Matsuhara, H,, 1989, Astmphys. J., 340, 729 Neff, S.G., and Ulvestad, J.S., Astron J.,1988, 96, 841 Neugabauer,G. Morton, G., Oke, J.B., Becklin, E.E., Daltabuit, E., Matthas, K., Persson, S.E., Smith, A.M., Soifer, B.T., Torresa-Peimbert, S. and Wynn- Williams, C.G., 1980 , Astrophys. J., 238, 502 Norman, C., and Scoville, N., 1988, AstrophysJ., 332, 124 Nonis, R.P., 1985, M.N.R.A.S., 216, 701 Oliva, E., and Moorwood, A.F.M., Astron & Astrophys 203, 278 Omodaka, T., Hayashi, M., Hasegawa, T., 1984, AstrophysJ., 282, L77 Ondrechen, M.P., 1985, A stro n .J .t 90, 8. Osterbrock, D.E., 1977, AstrophysJ., 215, 733 Osterbrock, D.E., 1985, Astrophysics o f Active Galaxies and Quasi-SteUar O b je c ts ed. Miller, J.S., Oxford University Press Osterbrock, D.E., 1989 Astrophysics of Gaseous Nebulae and Active Galactic N u d e i, Freeman, San Fransisco. Panagia, N., 1973, Astron. J., 78, 9 Panagia, N., 1978, in InfraredandSubmillimetreAstronomy, ed. Fazio, G., D. Reidel Pedlar, A., Booler, R.V., Spencer, R.E. and Stewart, O.J., 1983, M.N.R.A.S ., 202, 647. Pedlar, A., Dyson, J.E, Unger, S.W, 1985, M.N.R.A.S., 2 1 4 , 463 Persson, S.E., McGregor, P.J., Duncan, D.K., Lanning, H. and Geballe, T.R., ****, XVI. E.S.L.A.B. Symposium, ed. M. Keisler and J.P. Phillips. 248 Peterson, B.M., 1988, P. A .S .P ., 100, 18 Peterson, B.M., and Cota, S.A., 1988, AstmphysJ., 330, 111 Phillips, M.M., 1977, A s tr o p h y s JI, 215, 746. Phillips, M.M., 1978, AstivphysJ., 226, 736. Phillips, M.M., Aitken, D.K. and Roche, P.F., 1984, M.N.R.A.S., 207, 25. Phillips, M.M., Turtle, A.J., Edmunds, M.G., and Pagel, B.E.J., 1983, M.N.R.A.S., 203, 759 Pounds, K.A., and Turner, T.J., 1987, in VariabilityofGalacticandExtragalactic X -ray Sources, ed. Treves, A., Assodazione per l’Avanzamento delTAstronomia Press, W.H., Flannery, B.P., Teukolsky, S.A., and Vetterling, W.T., 1986, Numerical Recipes: The A rt ofScientific Computing, Cambridge University Press. Puxley, P.J., Hawarden, T.G., and Mountain, C.M., 1988, MNRAS, 234, 29p Reay, N.K., Atherton, P.D., Walton, N.A., 1986, M.N.R.A.S., 218, 13p Rees, M., 1984, Ann Revs Astron and Astrophys, 22, 471 Rees, M., 1986, in Proceedings o f the 119 die Symposium International Astronomical Union, Ed Swarp, G., and Kapahi, V.K., D. Reidel Rieke, G.H., 1976, AstrophysJ., 206, L I 5 Rieke, G.H., and Low, F.J., 1972, A s tiv p h y s J., 176, L95 Rieke, G.H., Cutri, R M., Black, J.H., Kailey, W.F., McAlary, C.W., Lebofsky, M.J. and Elston, R., 1985, AstrophysJ, 290, 116. Rieke, G.H., Lebofsky, M.J. and Low, F.I., 1985, A s tr o n J ., 90, 900. Rieke, G.H., Lebofsky, M.J., and Walker, C.E., 1988, AstrophysJ., 325, 697 Rieke, G.H., Lebofsky, M.J., Thompson, R.I., Low, F.J. and.Tokunaga, A.T., 1980, AstrophysJ., 238, 24. Rieke, G.H.and Lebofsky, M.J., 1981, AstrophysJ., 250, 87. Rieke, G.H.and Lebofsky, M.J., 1985, AstrophysJ\, 304, 326 Roche, P.F., Aitken, D.K., Phillips, M.M. and Whitmore, B., 1984, M.NR.A.S., 207, 35. Roche, P.F., Aitken, D.K., Phillips, M.M., and Whitmore, B., 1984, M.N.R.A.S., 207, 35 Roche, P.F., and Aitken, D.K., , 1985, M.N.R.A.S., 213, 789 Rodriguez Espinoza, J.M., Rudy, R.J., and Jones, B., 1986, AstrophysJ, 309, 76 249 Rowan-Robinson, M., and Crawford, J., 1985, RutherfordAppletonLab. Report, RAL85-086 Sanders, D.B. and Mirabel, I.F., 1985, AstrophysJ.Lett, 298, L.31. Sanders, D.B., Scoville, N.Z., Sargent, A.I., and Soifer, B.T., AstnophysJ., 324, L55 Sanders, D.B., Scoville, N.Z., Young, J.S., Soifer, B.T., Schloerb, F,P., Rice, W.L. and Danielson, G.E., 1986, AstrophysJ.LetL, 305, L.45. Sanders, D.B., Soifer, B.T., Elias, H.J., Madore, B.F., Matthews, K., Neugebauer, G. and Scovile, N.Z., 1988, A s tr o p h y s J., 325, 74. Sanders, D.B., Young, J.S., Scoville, N.Z., Soifer, B.T., and Danielson, G.E., 1987, AstrpphysJ., 312, L5 Sargent, A.I., Sanders, D.B., Scoville, N.Z. and Soifer, B.T., 1987, AstivphysJ., 312, L.35. Sargent, A.I., Sanders, D.B., Scoville, N.Z., and Soifer, B.T., 1987 in S ta r Formation in Galaxies, ed Lonsdale Persson, C.J., NASA conf publ NASA CP- 2 4 6 6 Scalo, J., 1986, Fund Cosmic Physics, 11,1 Schleorb, F.P., and Loren, R.B., 1982, Symposium on Orion Nebula to Honor HenryDraper, ed Glassgold, A.E., Huggins, P.J., and Schucking, E.L., New York Academy of Sciences Schwartz, R.D., Cohen, M. and Williams, P.M., 1987, AstrophysJ., 322, 403. Schweizer, F., 1982, AstrophysJ.,2 5 2 , 455. Schweizer, F.1983, in Internal Kinematics and Dynamics o f Galaxies, ed. Athanassoula, E., Reidel Dordrecht Scoville, N.Z., Sanders, D.B. and Clemens, D.P., 1986, AstrophysJ.Lett, 310, L.81 Scoville, N.Z., Sanders, D.B., Sargent, A.I., Soifer, B.T., Scott, S.L. and Lo, K.Y., 1986, AstrophysJ.LetL, 311, L.47. Scoville, N.Z., Soifer, B.T., Neugebauer, G., Young, J.S., Matthews, K. and Yerka, J., 1985, AsrrophysJ., 289, 129. Second Reference Catalgue of Bright Galaxies, de Vaucaleurs, G., de Vaucaleurs A., and Corwin, H.G., 1976Austin: University of Texas Press Sellgren, C. 1984, AstrophysJ., 2 7 7 , 623 Sekiguchi, K., 1987, AstrophysJ, 316, 145 Sewart, F.D., Hamder, F.R., Murdin, P. and Clance, D.H., 1983, AstrophysJ., 267, 698. Seyfert, C.K., 1943, Astrophys. J., 97, 28 250 Shull, J,M., and Beckwith, S., 1982,Ann Revs Astron andAstrophys, 20, 163 Simon, M., FeUi, M., Cassar, L., Fisher, T. and Massi, M., 1983, A s tr o p h y sJ 266, 623. Simon, M., Righini - Cohen, G., Fischer, J., Cassar, L., 1981, AstrophysJ., 251, 552 Slipher, V.M.,1917, Proc Amer Phil Soc, 56, 403 Smith, A., Larson, H.P., Fink, U., 1981, AstrophysJ., 244, 835. Smith, E.P., Heckman, T.M., Bothun, G.D., Romanishin, W. and BaHck, B., 1985, AstrophysJ., 306, 64. Smith, H.A., Larson, H.P. and Fink, U., 1981, Astrophys J., 244, 835. Smith, M.G., Weedman, D.W., and Spinrad, H., 1972, AstrophyLetters, 11, 21 Soifer, B.T., and Pipher, J.L., 1978, Ann Revs Astron & Astrophys, 16, 335 Soifer, B.T., helou, G., Lonsdale, C.J., Neugebauer, G., Hacking, P., Houck, J.R., Low, FJ., Rice, W., and Rowan-Robinson, M., 1984, AstwphysJ., 278, L71 Soifer, B.T., Sanders, D.B., Neugebauer, G., Danielson, G.E., Lonsdale, C.J., Madore, B.F., and Persson, S.E., 1984, AstrophysJ., 283, LI Soifer, B.T., Sanders, D.B., Neugebauer, G., Danielson, G.E., Lonsdale, C.J., Madore, B.F., and Persson, S.E., 1986, AstrophysJ., 303, L41 Spinrad, H., and Wing, R.F., 1969, Ann Revs Astron & Astrophys, 7, 249 Sramek, R.A. and Weedman, D.W., 1986, AstrqphysJ., 302, 640. Sternberg, A., andDalgamo, A., 1989, AstrophysJ'., 338, 197 Stone, J.L., Wilson, A.S., and Ward, M.J., 1988, AstrophysJ., 330, 105 Telesco, C.M., and Decher, R. 1988, preprint Telesco, C.M., and Gatley, I., 1984, AstrophysJ., 284, 557 Telesco, C.M., Beddin, E.E., Wynn-wiliams, C.G., Haiper,D.A, 1984, Astrophys. J., 282, 427 Telesco, C.M., Becklin, E.E., Wynn-wiliams, C.G., Harper,D.A, 1984, Astrophys. J., 282, 427 Telesco, C.M., Decher, R. and Gatley, I., 1985, AstrophysJ., 299, 896. Thompson, R.I., 1982, AstrophysJ., 257, 171. Thompson, R.I., 1984, AstrophysJ, 283, 165. Thompson, R.I., 1987, A s t r o p h y s J312, 784. Thompson, R.I., Lebofsky, M.J., Rieke, G.H.,1978,Astrophys. J., 222, L49 251 Thronson, H.M., Campbell, M.F., and Harvey, P.M., 1978, A stron J.. 83, 1581 Tinsley, B.M., 1980, Fund Cosmic Physics, 5, 287, Toomre, A. andToomre, J., 1973, S d .A m ., 299, 38. Toomre, A., andToomre, J., 1972, AstrophysJ., 178, 623 Treffers, R., 1976, AstrophysJ., 209, 793. Treffers, R.R., 1979, AstrophysJ., 233, L17 Turner, J.L. and Ho, P.T.P., 1985, AstrophysJ.Lett, 299, L.77. Turner, J.L., and Ho, P.T.P., 1983, AstrophysJ., 235, 392 Turner, J.L., Ho, P.T.P. and Beck, S.C., 1987, AstrophysJ., 313, 644. Ulrich, M.H., 1978, Astrophys. J., 219, 424. Ulvestad, J.S. Neff, S.G. and Wilson, A.S.,1987, Aston. J., 92, 22 Ulvestad, J.S., 1982, Astrophys. J., 259, 96. Ulvestad, J.S., Neff, S.G. and Wilson, A S., 1987, A s tr o n .J ., 92, 22. Veron-Cetty, M-P., Veron, P., 1986, Astron & Astrophys, Suppl., 66, 335 Wachter, K.W., Strauss, M.A., and Filippenko, A.V., 1988, AstrophysJ., 330, 91 Walker, C.E., Lebofsky, M.J., and Rieke, G.H., 1988, A s tr o p h y s J., 325 687 Ward, M.J., Geballe, T., Smith, M., Wade, R., and Williams, P., 1987, AstrophysJ., 316, 138 Weedman, D.W., 1977, Ann.Revs.Astron.andAstrophys., 15, 69. Willner, S.P., Soifer, B.T., Russel, R.W., Joyce, R.R. and Gillett, C.C., 1977, AstrophysJ., 217, L.121. Wilson, A.S. and Ulvestad, J.S., 1983, AstrophysJ., 275, 8. Wilson, A.S. and Ulvestad, J.S., 1987, Astrophys. J., 319, 105 Wilson, A.S., and Heckman, T.M., 1985, Astrophysics o f Active Galaxies and Quasi-StellarObjects, ed. Miller, J.S., Oxford University Press Wilson, A.S., Baldwin, J.A., Sun Sze-dung, Wright, A.E., 1986, AstrophysJ., 310, 121 Wright, G.S., Joseph, R.D. and Meikle, W.P.S.,1984, N ature; 309, 430. Wright, G.S., Joseph, R.D., Robertson, N.A., James, P.A., and Meikle, W.P.S., 1988, M.N.R.A.S., 233, 1 Wynn-Williams, C.G., and Becklin, E.E., 1974, PASP, 86, 5 252 Wynn-Williams, C.G., and Becklin, E.E., 1985, A stroph ysJM 290, 108 Wynn-WiUiams, C.G., Becklin, E.E., and Scoville, N.Z., 1985 A stro p h y s. 7.,297, 607 Young, J.S., and Scoville, N.Z., l9S2tAstrophys. 7, 258, 467 253) (6.3)