Infrared Spectroscopy of Super Novae
Infrared Spectroscopy of Super novae
by
Jason Spyromilio
Astrophysics Group Blackett Laboratory Imperial College of Science, Technology and Medicine London SW7
A thesis submitted for the degree of Doctor of Philosophy of the University of London and for the Diploma of Membership of the Imperial College
August 1989
1 Abstract
Infrared spectroscopic observations of snpernovae 1986G and 1987A are presented and discussed. Late-time observations of SN 1986G at CTIO show no evidence for any spectral features. Contemporary observations at the AAT show a low S/N feature at 1.64 fim. The observations of SN 1987A are unique both in spectral and temporal coverage. The spectra exhibit forbidden lines of iron group elements believed to arise from the radioactive decay of 56Ni and of silicon. In addition allowed lines of hydrogen, oxygen, silicon, sodium, magnesium and calcium are identified. Models for the interpretation of these spectra are presented. Modelling of hydrogen lines in the early-time spectra of SN 1987A suggests that large deviations from Boltzmann populations are present in the envelope. A late-time model for Type la supernovae is also presented. The energy deposition of the 7-rays resulting from the radioactive decay of 56Ni is calculated through the use of Monte Carlo techniques and a self-consistent temperature and ionisation structure is calculated. A model spectrum is compared with the spectra of SN 1986G. Finally we discuss the molecular emission present in the spectra of SN 1987A and present models of the vibrational-rotational spectra of carbon monoxide and singly ionised carbon monoxide and compare them with the data.
2 Acknowledgements
I would like to thank the many people who helped me during the course of the work described in this thesis. My thanks go to Peter Meikle under whose careful supervision this work was under taken. He provided a constant source of enthusiasm, essential help and advice. I would also like to thank him and John Quenby for dealing with the bureaucratic problems arising from my non-SERC student status.
Special thanks are due to David Allen for educating me in the techniques of infrared spectroscopy and for his constant enthusiasm during observing runs. Without his help and advice the quality of data presented here would certainly be much poorer. Also special thanks are due to Dick Learner for his patience in our regular tutorials on molecular and atomic physics and for convincing me that modelling the molecular emission in SN 1987A was ‘straightforward’. I would also like to thank my friends and colleagues for stimulating and helpful discus sions. Thanks go to Phil Andrews, Chris Bell, Rene Doyon, James Graham, Phil James, Bob Joseph, Steve Matcher, Robert Renton, Sunil Sidher, Tim Sumner, Gian Varani and Martyn Wells. Finally I would like to thank my parents for their support and encouragement during my student years.
3 Contents
1 Introduction 11
1.1 Observed characteristics of supernovae ...... 12
1.1.1 Spectroscopic classification of supernovae...... 12
1.1.2 Electromagnetic emission properties of supernovae...... 15
1.2 Explosion mechanisms...... 19
1.3 The importance of infrared observations...... 22
1.3.1 Early time observations...... 22
1.3.2 Late time observations t > 150 d a y s...... 23
1.4 Conclusions...... 24
2 Techniques of Infrared spectroscopy of Supernovae 25
2.1 Introduction...... 25
2.2 Devices ...... 26
2.2.1 Circular Variable Filter (C V F )...... 26
2.2.2 Cooled Grating spectrometers (CGS)...... 27 4 2.3 Observing techniques...... 27
2.4 Factors influencing the quality of the d a ta ...... 28 2.5 Flux calibration and removal of Telluric features...... 29
2.6 Wavelength calibration...... 30
3 Infrared spectroscopy of Supernovae 1986G and 1987A 31
3.1 Supernova 1986G...... 31 3.1.1 The early time observations...... 32 3.1.2 The late time observations ...... 32 3.2 Supernova 1987A...... 43 3.2.1 The observations ...... 43
3.2.2 Results ...... 46 3.3 Conclusions...... 65
4 Interpretation of the Infrared spectra of supernovae 79
4.1 Early time spectra...... 79 4.1.1 Theory of continuum formation...... 80 4.1.2 Theory of P-Cygni profile formation...... 81 4.1.3 Modelling the early time spectra of SN 1987A...... 86 4.2 Late time spectra...... 91 4.2.1 The 7 rays ...... 91 4.2.2 Non-thermal electron energy deposition...... 92
5 4.2.3 The thermal and ionisation balance 95
4.2.4 The final spectrum...... 100 4.2.5 Comparison with observations ...... 100 4.3 Conculsions...... 102
5 Molecules in Supernovae 103 5.1 The infrared spectra of diatomic molecules...... 103 5.1.1 Energy level structure...... 103 5.1.2 Selection rules and band structure...... 107 5.1.3 Radiative rates ...... 108 5.2 Carbon monoxide emission in SN 1987A...... 109 5.2.1 The discovery of emission bands ...... 109
5.2.2 The location of the emitting C O ...... 110 5.2.3 Modelling the emission from C O ...... 110 5.2.4 Interpretation of the CO emission ...... 114
5.3 Other molecules in SN 1987A ...... 116
5.4 Molecular formation in SN 1987A ...... 117
6 Conclusions and further work 119
References 122
A Publications 129
6 List of Figures
3.1 Spectrum of SN 1986G AAT March 1987 ...... 38
3.2 Spectrum of SN 1986G CTIO 4m March 14th 1987 ...... 38
3.3 Spectrum of SN 1986G CTIO 4m March 16th 1987 ...... 39
3.4 Spectrum of SN 1986G AAT April 1987 ...... 39
3.5 Spectrum of SN 1986G AAT May 1987 ...... 40
3.6 Spectrum of SN 1986G CTIO 4m June 21st 1987 ...... 40
3.7 Spectrum of SN 1986G AAT March 1988 ...... 41
3.8 Binned spectrum of SN 1986G AAT April 1987 ...... 41
3.9 Binned spectrum of SN 1986G AAT May 1987 ...... 42
3.10 March 1987 J window spectrum of SN 1987A...... 66
3.11 March 1987 H window spectrum of SN 1987A...... 66
3.12 March 1987 K window spectrum of SN 1987A ...... 67
3.13 March 1987 L window spectrum of SN 1987A...... 67
3.14 June 1987 J window spectrum of SN 1987A...... 68
3.15 June 1987 H window spectrum of SN 1987A...... 68
7 3.16 June 1987 K window spectrum of SN 1987A...... 69
3.17 June 1987 L window spectrum of SN 1987A...... 69
3.18 Septemberl987 J window spectrum of SN 1987A ...... 70
3.19 Septemberl987 E window spectrum of SN 1987A ...... 70
3.20 Septemberl987 K window spectrum of SN 1987A ...... 71
3.21 Septemberl987 L window spectrum of SN 1987A ...... 71
3.22 Septemberl987 M window spectrum of SN 1987A...... 72
3.23 Novemberl987 J window spectrum of SN 1987A ...... 72
3.24 Novemberl987 H window spectrum of SN 1987A ...... 73 3.25 Novemberl987 K window spectrum of SN 1987A ...... 73
3.26 Novemberl987 L window spectrum of SN 1987A ...... 74
3.27 Decemberl987 J window spectrum of SN 1987A ...... 74
3.28 Decemberl987 H window spectrum of SN 1987A ...... 75
3.29 Decemberl987 K window spectrum of SN 1987A ...... 75
3.30 Decemberl987 L window spectrum of SN 1987A ...... 76
3.31 February 198® J window spectrum of SN 1987A...... 76
3.32 Februaryl980 IT window spectrum of SN 1987A...... 77
3.33 Februaryl980iT window spectrum of SN 1987A...... 77
3.34 Februaryl98g L window spectrum of SN 1987A...... 78
4.1 Supernova co-ordinate system ...... 83
8 4.2 Optical depth effects on P Cygni line profile...... 84 4.3 Density power law effects on P Cygni line profile...... 85
4.4 Blending of two P Cygni lines...... 86
4.5 Pcygni model spectrum for J w indow ...... 88
4.6 Pcygni model spectrum for H window...... 89 4.7 Pcygni model spectrum for K window...... 89 4.8 Pcygni model spectrum for L w indow ...... 90
4.9 Hydrogen departure coefficients...... 90 4.10 7 ray deposition curve...... 92 4.11 Division of primary electron energy ...... 94 4.12 Cooling curve for F elll...... 97 4.13 Effects of scattering on line profile ...... 101 4.14 Model Fe spectrum...... 102
5.1 Molecular energy level potential ...... 104 5.2 Hund’s coupling cases...... 106 5.3 CO model day 1 9 2 ...... 113 5.4 CO model day 255 ...... 114
5.5 CO model day 255 (excluding CO+) ...... 115
5.6 CO model day 284 ...... 116 5.7 CO model day 349 ...... 117
9 List of Tables
3.1 Log of observations of SN 1986G...... 33 3.2 Apertures chops and calculated backgrounds for spectra of SN 1986G . . . 36 3.3 Log of observations of SN 1987A...... 44 3.4 Evolution of the continuum in spectra of SN 1987A ...... 47 3.5 Features observed in the day 18 spectra of SN 1987A ...... 49 3.6 Line identifications for the day 110-349 spectra of SN 1987A ...... 50 3.7 Hydrogen line ratios in the spectra of SN 1987A ...... 57
3.8 FWHM of lines in the spectra of SN 1987A...... 63
3.9 Redshifts of emission line maxima in the spectra of SN 1987A ...... 64
4.1 Results of self consistent m odel ...... 99
5.1 CO model param eters...... 112
10 Chapter 1
Introduction
Observations of supernovae in the form of bright ‘guest’ stars have been carried out for approximately 2 millenia. The first such objects to be studied in a systematic fashion were the Renaissance supernovae of 1572 A.D. and 1604 A.D. which also happened to be the last supernova in our galaxy. Naked eye supernovae do not appear more than half a dozen or so times per millenium (Clark & Stephenson, 1982). However with the advent of the telescope and deliberate searches the number of observed events has risen to an average of approximately 20 per year in the last 50 years. Nearly all of these supernovae have been faint extragalactic events. Few supernovae have been brighter than 12th magnitude when discovered and even fewer have been observed to brighten after discovery. So we are indeed very fortunate that on February 23rd 1987 in the Large Magellanic Cloud (LMC) the nearest and brightest supernova for the last 300 years was discovered by Ian Shelton (Madore & Kundel 1987). With the present set of instrumentation and photon collecting facilities available to us it has provided an opportunity to observe and analyse the events accompanying the death of a star, not only at a uniquely high signal to noise ratio but also for an unprecedented length of time after the initial explosion. In this chapter we shall discuss some of the characteristics of supernovae and some of the current theories which are being applied to these phenomena.
11 12 Chapter 1. Introduction 1.1 Observed characteristics of supernovae
It is important to note that even though the electromagnetic emission of a supernova is spectacular, it represents a minor manifestation of the total event. Excluding the Sun, SN 1987A is the only astronomical object from which non-electromagnetic emission, namely neutrinos, has been detected opening up the new field of extragalactic neutrino astronomy. Approximately 100 times more energy escaped in the form of neutrinos than was deposited in the ejecta of SN 1987A. Most of the energy deposited in the ejecta is in the form of kinetic energy (Woosley & Weaver 1986) with only one hundredth of this energy emerging in the form of electromagnetic emission. Although at various stages of its evolution a supernova emits radiation from the 7-ray end of the spectrum to the radio almost all supernovae have been discovered in the optical. For most supernovae very little spectroscopic data are available and almost all of that is in the optical.
1 .1.1 Spectroscopic classification of supernovae
Supernovae initially were thought to be a single class of event primarily because SN 1937C in IC 4182 and SN 1937D in NGC 1003, which formed the basis for the early discussions of these events, both happened to be of the same spectroscopic class. The discovery of SN 1940C in NGC 4275 led Minkowski (1941) to separate supernovae into two distinct classes which he named named Type I’s and Type ITs. To the present the division of supernovae into these two distinct classes has remained valid with Type I’s now themselves being split into Type la’s and Type lb’s. Only a few supernovae do not appear to fit into either category. The classification is primarily based on the observed spectroscopic properties of the supernovae. The light curves exhibited can also be used to classify supernovae but are less reliable and clearly require more observations. On the basis of the light curve classification Type II supernovae are divided into Type lip’s which exhibit a plateau phase during their first 100 days and Type II/’s which exhibit a linear decline of their light curves after maximum light. A result of the spectroscopic classification is that pre-spectrographic supernovae cannot be confidently assigned to either class although the study of the remnants may hold clues as to what type of supernova was responsible for the observed remnant. The simplest spectroscopic distinction between Type I’s and Type II’s is that the former do not exhibit hydrogen lines in their spectra (Oke & Searle 1974). Chapter 1. Introduction 13
Kirshner et dl. (1973) argued that the spectra of the two types of supernovae axe formed in the same fundamental conditions. The overall energy distribution in these early (t <100 days) spectra implies the presence of an opaque, cooling, and expanding atmosphere (Arp 1961; Mustel 1971) with an effective temperature which ranges from 10,000 K for the peak brightness epoch to <^5000 K after a few weeks. The character of the spectral lines is similar in both classes. The lines exhibit P-Cygni profiles consisting of a broad emission line centered about the rest wavelength of the transition (in the rest frame of the parent galaxy) and a broad blueshifted absorption trough. Such lines can be formed by resonant scattering in an expanding atmosphere (see chapter 4). The spectrum evolves in a very similar fashion for both types of supernova with the continuum fading but most strong emission lines persisting throughout the evolution of the supernova from a stellar to a nebular-like object.
Type I supernovae
Type la’s are probably the most homogeneous class of supernovae with two very good standards, SN 1937C and SN 1972E, both of which had extensive coverage during the first year of their evolution (Minkowski 1939; Kirshner et dl. 1973). The excellent agreement between the optical spectra of these two objects at various stages of their evolution has defined spectroscopic templates for Type la events. The early spectra exhibit P-Cygni lines superimposed on a strong continuum, whereas the late-time spectra exhibit strong emission features on top of a weaker continuum. As noted above the main difference between this type of supernova and Type II’s is the lack of strong hydrogen lines in their spectra. In order to identify the transitions responsible for the features observed modelling of the spectra is required, as the large number of overlapping Doppler broadened lines make it difficult to identify any particular feature as the result of a single transition. There has been some success in modelling Type la spectra with most prominent features being reproduced using transitions of He, O, Mg, Si, S and Ca (but no H) with near solar abundances (Branch et dl. 1983). Later spectra of Type la supernovae require Fe to be included in the models, with spectra a year after maximum being fitted almost entirely by transitions of Fe II and Fe III (Kirshner & Oke 1975). SN 1986G, the nearest Type la since SN 1972E, was discovered (Evans 1986) in NGC 5128 and has provided us with unique early-time infrared spectra of such an event (Frogel et dl. 1987; Graham private 14 Chapter 1. Introduction
communication). In these spectra absorption features due to Na I and Mg I are identified but a few features remain unidentified. These early-time spectra together with late-time infrared spectra of SN 1986G will be discussed in chapter 3. Type lb supernovae are characterised by the absence of the 6150 A absorption feature which Pskovskii (1968) identified as the blue-shifted 6347 A and 6371 A Si II lines. Initially classed as peculiar Type I’s they are now believed to arise from a totally different kind of progenitor star and could account for more than half the Type I supernovae observed. The prototype Type lb supernova is SN 1983N in M83, for which multifrequency observations are described in Panagia et al. (in preparation). Late-time optical spectra of SN 1985F (Filippenko & Sargent 1986) exhibit strong emission lines due to transitions of [0 I], Ca II, [Ca II] and [C I]. Late-time infrared spectra of SN 1983N were discussed by Graham et al. (1986), who reported the detection of a strong emission feature at 1.64 pm. This may be, as Grahamet al. (1986) suggested, due to [Fe II] or, as Oliva (1987) claimed, due to [Si I]. Type lb supernovae have also exhibited radio emission which is absent from Type la’s.
Type II supernovae
A supernova is classed as Type II if it exhibits hydrogen lines in its spectrum (Minkowski 1941; Oke Sz Searle 1974). A few days after the explosion the optical spectrum exhibits a relatively smooth continuum. All lines are weak except for Ha. Weak P-Cygni profiles can be seen in the wavelengths of H/? and H 7 (Kirshner et al. 1973). A few weeks later the continuum is redder and additional lines of Mg, Ca and Na become prominent exhibiting P-Cygni profiles. At later times [O I] and [Ca II] lines become prominent in emission while Ha remains strong (Uomoto & Kirshner 1986). SN 1987A resulted in a large amount of optical spectra. The spectra are generally consistent with the evolution described above with lines of Ba, and Sc also being identified (Williams 1987). The infrared observations of SN 1987A, the only Type II supernova with extensive coverage in this spectral region, are discussed in detail in chapter 3. Briefly the infrared data may be described as consistent with the optical data in their continuum and line properties. Chapter 1. Introduction 15
1.1.2 Electromagnetic emission properties of supernovae
Radio
The Type II supernovae SN 1979C and SN 1980K exhibited quite similar radio light curves (Weiler et dl. 1981) but with different turn-on times (SN 1980K became visible at 6 cm almost immediately after the outburst whereas SN 1979C only became detectable after 1 year). Given the supernova size the large radio luminosity implies an effective temperature T > 104 K i.e. the origin of the radiation is non-thermal. If the emission originates in the supernova then a central pulsar is required to generate the required magnetic field and relativistic particles (Pacini & Salvati 1981). Chevalier (1982) suggests that the circumstellar interaction region is a more likely location of the radio emission. Rayleigh-Taylor instabilities can generate sufficient random motions to build up magnetic fields and may also generate relativistic electrons. Alternatively shock acceleration may produce the necessary relativistic electrons. Radio emission from SN 1987A was detected, from day 2 to day 11 (after the explosion), at a variety of frequencies ranging from 0.843 GHz to 8.4 GHz and exhibited an initial rise but then declined below detection limits (Turtle et dl. 1987). Chevalier & Fransson (1987) find that the observed radio luminosity is compatible with the circumstellar interaction model but some absorbing medium is required in order to explain the slow rise exhibited. They suggest that the early turn-on of the radio emission implies a relatively low density circumstellar medium. After fading the supernova has not been detected in the radio again. However Chini et al. (1988) have reported millimetre observations and the detection of emission at 1.3 mm on day 562. Subsequent observations at the same wavelength have only resulted in upper limits on the flux coming from the supernova.
Although radio emission is absent from Type la events, Type lb’s have been detected in this spectral region. In particular SN 1983N and SN 1984L were detected at 6 and 20 cm (Weiler et al. 1986). Here again the turn-on is abrupt, but it occurs closer to the optical outburst than in Type II’s. Indeed in the case of SN 1983N the turn-on was 11 days before maximum optical light. The mechanisms for the production of the radio emission in Type lb’s are believed to be similar to those taking place in Type II’s. 16 Chapter 1. Introduction
IR 1—100 pm
Infrared emission from supernovae can generally be divided into two categories: emission from dust and the rest. Emission from dust is usually considered in the context of the infraxed excess observed in some supernovae ( e.g. SN 1979C, SN 1980K.) It is generally believed to arise from circumstellar dust which is heated by the UV and optical radiation around the period of maximum light and is observed as an echo (Dwek 1983; Graham et oil. 1983). Alternatively dust may form inside the ejecta as they cool (Dwek et dl. 1983; Dwek 1988a) and although no conclusive evidence exists for such a process, recent infrared and optical observations of SN 1987A indicate that dust may have started condensing in the ejecta around day 600 after the explosion (Danziger et dl. 1989; Spyromilio et al. 1989). SN 1987A has also exhibited emission from circumstellar dust (Roche et al. 1989). Any such dust will have to lie outside the evaporation radius which therefore determines the minimum delay between the maximum of the optical light curve and the emergence of an infrared excess due to the emission from dust. It is evident from the lack of substantial reddening that only a small fraction of the optical radiation is captured by the dust and therefore the subsequent re-emission in the infrared becomes detectable only at late times. At early times the infraxed spectrum is very similar to the optical. There is a near thermal continuum, which dominates the emitted flux for the first 100 days, with P-Cygni lines superimposed on it. In the case of SN 1987A the spectrum evolved with the continuum fading and forbidden lines becoming prominent in emission. No other supernova has had extensive spectral coverage at late times. It has been suggested (Axelrod 1980) that after approximately 2-3 years an ‘infrared catastrophe’ will occur resulting in most of the emitted radiation being in the fax infraxed band (see chapter 4). The infraxed catastrophe has yet to be observed in a supernova possibly because of insufficient extended temporal coverage. SN 1987A will surely prove to be the testbed for this effect.
Optical 3000-10000 A
Most photometric and spectroscopic observations of supernovae have been made in this wavelength band. The optical display staxts soon after the UV flash emerges, when the optically thick region - the photosphere - cools. At very early times the optical spectrum Chapter 1. Introduction 17 can be approximated by a black-body (T~104 K) with the temperature later dropping to about 5000 K. Most of the features in the early spectra exhibit P-Cygni profiles, presum ably due to resonant scattering in the expanding envelope. However, the hydrogen lines and in particular Ho: have, until recently, proved difficult to model using simple resonant scattering models. However Hoflich (1987), by combining non-LTE effects with scattering, has successfully fitted the early spectra of SN 1987A including Ha. The early evolution of the optical emission can be explained in terms of the photosphere receding into the ejecta and cooling. At later times the optical spectra exhibit a weak continuum with strong emission lines which arise predominantly from forbidden transitions.
UV and soft X-ray
The first electromagnetic emission from a Type II supernova is predominantly UV radia tion. Falk (1978) and Klein & Chevalier (1978) have shown that as the shock reaches the photosphere of the progenitor star a ~103 sec burst of ultraviolet and very soft X-rays is produced with a peak luminosity of order 1045 ergs/s corresponding to an effective tem perature of order 105 K, but with a non-Planckian spectrum. As the shock propagates through the less dense regions above the photosphere of the progenitor a further burst of hard X-rays (hu > O.lMeV), lasting a few tens of minutes, may arise with a slight delay from the softer burst. Soft X-rays from SN 1980K were detected by the Einstein Observatory 35 days after light maximum but quickly faded below detection limits (Canizares et dl. 1982). These X- rays are clearly not associated with the UV flash. A plausible explanation for this emission is the inverse Compton scattering of the optical photons emitted at the photosphere off fast electrons in the circumstellar material (Canizares et al. 1982). Alternatively the X-rays may be produced in the shocked supernova gas when it interacts with the circumstellar material (Chevalier 1982).
Most of subsequent UV radiation appears in a continuum depressed by heavy line blanketing due to transitions of ionised species. The lines observed in the UV spectra may arise from at least three distinct regions. These axe the ejecta of the supernova, the circumstellar material and the interstellar medium. SN 1979C, a Type II, exhibited narrow interstellar absorption features (Panagia et al. 1980) which confused the initial 18 Chapter 1. Introduction analysis of the data. The velocity widths inferred for the emission lines in these spectra were also low, implying that they were formed in circumstellar material (probably the wind of the progenitor). However, reanalysis of these data taking into account the effect of the interstellar lines on the line widths showed that the lines were broad and therefore formed just above the photosphere (Fransson et al 1984).
Spectra of SN 1981B, a Type la, (Panagia 1982) are concentrated about the time of maximum light. After interstellar reddening has been taken into account the spectra indicate a continuum level which is a factor of approximately 4 lower than that of the contemporary optical spectra. The spectra exhibit distinct bands above this weak contin uum. These spectra can be best explained by considering the effect of scattering by Fe II and Co II of the photospheric emission (Branch &: Venkatakrishna 1986). IUE spectra of SN 1983N, a Type lb, are similar to those of SN 1981B (Panagia et al. in preparation). Spectra of SN 1987A taken shortly after discovery suggest a rapid evolution of the spectrum during the eaxly stages of the supernova (Cassatella et al. 1987). The first spectra, taken 14 hours after discovery, match neither type of event. However, by day 3 after the explosion, the long wavelength portion of the UV spectrum resembles a Type la spectrum. The absence of a substantial radio flux at early times (Turtle et al. 1987) argues for a low mass of circumstellar material. Therefore it is not unreasonable that SN 1987A exhibited a Type la like spectrum (Chevalier & Fransson 1987). Computed UV spectra (Lucy 1987) suggest that line blocking, principally by lines of Fe II and Fe III, accounts for for the main features in the UV at early times. However, after ~ 100 days, narrow UV emission lines, presumably from circumstellar gas became visible (Cassatella 1987). The lines have been used by Lundqvist & Fransson (1987) to derive the properties of the circumstellar material and the UV flash. The UV spectrum of SN 1987A at these later times is contaminated by the emission from two other blue stars ~ 1-2 arcsec away. Until the supernova has faded enough to allow us to obtain spectra of the ‘background’ stars it will be difficult to unambiguously analyse these spectra.
Hard X-ray and 7-ray
As mentioned above some hard X-rays may be produced during the initial outburst. How ever, the emergence of 7-rays and hard X-rays as the supernova envelope becomes optically Chapter 1. Introduction 19 thin is fax more crucial to our understanding of the evolution of the supernova. The pro duction of radioactive 56Ni during the initial explosion is central to the current theories relating to the late time evolution of supernovae. It is believed that it is the radioac tive decay of this 56Ni to 56Co and subsequently to 56Fe that provides the energy for the late time light curve of Type II events and probably the entire electromagnetic event in Type I’s. The radioactive decay of the longer lived 56Co produces 7-rays that are Compton degraded into X-rays. Only SN 1987A has been near enough to detect these 7-rays and X-rays. These were detected by the X-ray satellites Ginga and Kvant (Dotani et al. 1987; Makino 1988; Sunyaev et al. 1987), the Solar Maximum mission satellite (Matz et al. 1988a,b), and balloon-borne detectors (Sandie et al. 1988; Cook et al. 1988). The time at which7 -rays and hard X-rays emerge depends on the level of mixing of radioactive 56Co in the outer layers (Kumagai et al. 1988; Pinto & Woosley 1988; Ebisuzaki & Shibazaki 1988). The fact that these7 -rays and X-rays were detected soon (~150 days) after the explosion suggests that the 56Co has been mixed into the hydrogen envelope of SN 1987A (Kumagai et al. 1988).
1.2 Explosion mechanisms
The theoretical scenarios leading to the explosions of Type II and Type I supernovae are totally different, with regard to both the progenitor evolution and the mechanism that ends the star’s life. The common characteristic is the explosion energy of order 10 51-1053 ergs. In both type’s almost all of the explosion energy goes into the kinetic energy of the ejecta with only a small fraction contributing to the electromagnetic event. Central to our understanding of supernovae is the nature of the progenitors. In general, Type II supernovae are found in the spiral arms of Sab-Sd galaxies (Tammann 1982) and are not observed in elliptical galaxies. This seems to point to a population of massive stars as the progenitors of Type II supernovae. Type la’s have been observed in all types of galaxies. However, their occurrence in elliptical galaxies suggests they have low mass progenitors (Tammann 1978). Type lb’s on the other hand may be associated with H II regions in spiral galaxies (Panagia & Laidler 1988) suggesting a massive progenitor. In this section we shall briefly describe the ‘standard’ explosion models for the different types of supernovae. 20 Chapter 1. Introduction
Type la
Type la supernovae are believed to be the result of thermonuclear explosions in white dwarfs. The lack of hydrogen lines in the spectra and the large masses of iron seen the late time spectra (Kirshner & Oke 1975, Meyerott 1980, Axelrod 1980) supports this scenario. The fusion energy released may propagate through the star as a detonation or a deflagration. In the former case a supersonic front carries the thermonuclear burning through the stax whereas in the latter case the burning propagates subsonically. The ‘standard’ model for Type la events involves an accreting carbon-oxygen white dwarf in which degenerate thermonuclear runaway occurs. This is caused by the temperature dependence of the carbon burning reactions and the isolation of a small central region from the convective layers above it thus preventing cooling of the burning region. The runaway therefore does not occur at the centre of the stax but at one or more locations on the shell above the central region. The thermonuclear fusion results in ~10 51 ergs being deposited in the star. The burning produces ~1 M© of iron-group elements with more than half being in the form of radioactive 56Ni which then provides the energy for the electromagnetic emission.
Type lb
Type lb supernovae could result from an explosion similar to that occurring in Type la’s with the ignition of helium rather than carbon occurring further off-centre than for Type la’s with a detonation being the more favoured outcome (Branch & Nomoto 1986). Panagia & Laidler (1988) argue for a binary configuration for the progenitor stax in which the primary stax evolves into a relatively massive star (~5 M©) which explodes when the secondary reaches the red supergiant stage. This model would explain both the composition of the progenitor, inferred from the spectra, and the radio emission observed. Alternatively a Type lb supernova could result from the collapse of the highly evolved core of a Wolf- Rayet star - t.e. a massive stax that has lost its hydrogen envelope (Wheeler & Levreault 1985; Ensman & Woosley 1988). As in Type la’s the electromagnetic event is driven by the radioactive decay of 56Ni produced by the detonation or the shock produced by the core collapse. Chapter 1. Introduction 21
Type II
Type II events are generally believed to be the consequence of core collapse in red super- giant stars. As mentioned previously Type II supernovae are characterised by the presence of hydrogen lines in their spectra. Therefore the progenitor star has to maintain its hy drogen envelope up to the moment of the explosion. The limits on the mass range for a star that will result in a Type II supernova are fixed at the lower end by the mass of a star that will become a white dwarf rather than undergo core collapse and at the higher end by the mass of a star that will maintain its hydrogen envelope until the explosion. This mass range is thought to be 8-40 M® (Woosley & Weaver 1986 and references therein). The core collapses when electron degeneracy pressure can no longer support it (i.e. when the mass exceeds the Chandrasekhar limit). The process of core collapse is aided by electron capture by iron-group elements and photodisintegration of the iron core. Once these pro cesses start occurring the core contracts rapidly. When nuclear densities are exceeded the core stops contracting and ‘bounces’. What follows is not clear. A small fraction (~ 1%) of the gravitational energy released heats and ejects the outer mantle/envelope but most of the energy appears in the form of neutrinos. The explosion may be a prompt or a delayed one. In the former case pressure waves propagating outward accumulate where the Mach number equals one, forming a shock which then propagates outward through the progenitor depositing its energy in the mantle and envelope of the star. However this mechanism will only work if the shock does not lose too much of its energy when escaping from the core. If the shock stalls we are left with a core still containing approximately 1053 ergs of energy mostly in the form of neutrinos. The delayed explosion mechanism may still produce a supernova. As the trapped neutrinos diffuse out and escape they may deposit their energy in the dense shock region to restart the outward motion. Whether prompt or delayed , the explosion will result in an energetic shock travelling through the relatively dense regions of the mantle and envelope of a massive star. The first evidence of this shock will emerge when it reaches the surface of the star and the UV and soft X-ray pulse is produced. The early electromagnetic emission arises from the shock-heated and ionised envelope of the progenitor star which remains optically thick for some time. The late-time electromagnetic emission is powered by the radioactive decay of 56Ni produced in the shock-heated ejecta at the time of the explosion or by radiation from the compact remnant formed in core collapse. 22 Chapter 1. Introduction
We axe now in the fortunate position of knowing the nature of one star that did explode and produce a Type II supernova. The progenitor of SN 1987A was Sanduleak -69 202 (Gilmozzi et al. 1988) a B3 la star (Sanduleak 1969; Rousseau et d. 1978) which does not fit the description of the progenitor given above. This does not mean that the picture described above is wrong. In many respects SN 1987A was an unusual event. Although classified as Type II in view of the presence of hydrogen lines in its spectrum, it was subluminous, exhibited large expansion velocities (Hanuschik & Dachs 1987), and produced an unusual light curve. All these observed characteristics can be explained in terms of the progenitor star’s compact nature. To date the electromagnetic emission from SN 1987A can be understood in terms of the effects of the shock travelling through the ejecta and the radioactive decay of 56Ni. No conclusive evidence exists for any substantial contribution by a compact stellar remnant. The fact that neutrinos were detected from the supernova at KAMIOKANDE-II (Hirata et al. 1987) and at IMB (Bionta et al. 1987) is the most important vindication of the theory that core collapse can lead to a supernova explosion.
1.3 The importance of infrared observations
1.3.1 Early time observations
We have described above how the infrared spectra, at early times, are very similar to the optical spectra. So why study them? For the use of/supernovaet>jfeTE as distance indicators knowledge of how the continuum behaves at longer wavelengths is crucial. It is important to note that the use of/supernovae^-31 as distance indicators is independent of the distance ladder. They therefore can be used as calibrators for the distance ladder. As in the infrared the lines observed are widely spaced in energy they suffer less from blanketing than their IJV and optical counter parts. This has dual advantages when using the Baade-Wesselink method for the determination of the distance to supernovae. The absence of strong line blanketing makes the continuum easier to determine. In addition the P-Cygni lines are ‘cleaner’ allowing for an easier determination of the velocity of the absorption trough. The presence in the infrared of the higher series of hydrogen provides an opportunity to determine the populating mechanism for the high principal quantum number (3 < n < 15) Chapter 1. Introduction 23 energy levels and therefore the conditions in the ejecta at these epochs (see chapter 4). In the early infrared spectra of SN 1987A transitions of elements not observed in the optical can also be found ( e.g. Strontium).
1.3.2 Late time observations t > 150 days
The late time IR observations of supemovae have a dual purpose. One is to observe the actual ejecta and the other is to look for emission from dust. To account for the late-time luminosity there must be a source of stored energy. Possibilities are energetic particles either from radioactive decays or a pulsar formed by the collapsing core in the case of a Type II or even possibly a Type lb supernova. In the case of Type la’s it is almost certain that no compact remnant remains, so radioactivity must be the sole source of energy. In Type II’s and possibly Type lb’s radioactivity is almost certainly the principal source for the first couple of years after the explosion. The energetic photons from the radioactive decay will either escape or Comptonize until the energy is deposited in the gas in the form of relatively high ionisation and a quasi-thermal electron gas. The electron gas then collisionally excites the low lying levels of the atoms in the ejecta. Many of these transitions are observed in the infrared. The products of the explosive nucleosynthesis that occurs when a supernova explodes are of considerable interest. The infrared may be the only spectral region where we can test this theory. Since in the infrared the effects of line and continuum opacity are lower than in other spectral regions these lines can provide unique information about the core of the supernova. Also, when the electron gas has cooled sufficiently it will only be able to excite the forbidden ground-state fine-structure lines and the ‘infrared catastrophe’ will occur with almost all the (optical/infrared) flux coming out at far-infrared wavelengths (Axelrod 1980). The onset of the ‘infrared catastrophe’ will depend on the conditions in the ejecta and on the possible energy source. In the late time spectra dust emission may become prominent. Observations of emission from the dust improve our understanding not only of the physics of circumstellar material but also, in the case where the dust may form inside the ejecta, of the chemistry of dust condensation. 24 Chapter 1. Introduction 1.4 Conclusions
It is clear that many aspects of supernovae and their theory remain unanswered. For example the nature of the progenitors of Type lb’s, the mechanism for the production of the radio emission, the precise mechanism of the explosion for all types of supernovae, the timing of the infrared catastrophe are all uncertain. Some of these questions may be answered through future observations of supernovae and others through more detailed theoretical work. Infrared observations are already playing a crucial role in the case of SN 1987A and have produced some of the most exciting results for past supernovae ( e.g. Graham et dl. 1986). Chapter 2
Techniques of Infrared spectroscopy of Supernovae
In this chapter we shall describe the basic techniques used in infrared 1 spectroscopy as applied to supernovae. Some of the methods described here are specific to our observations of SN 1987A. To take full advantage of the proximity of SN 1987A both new observing procedures and some new software were developed (within the framework of the STAR- LINK Figaro data reduction package) in order to produce the best possible data from this unique object. Therefore some of the techniques described below are non-standard and reflect the advantages and disadvantages of working with the infrared spectrometers available at the Anglo-Australian Telescope. The details of the observing programmes which form part of this work are presented in chapter 3.
2.1 Introduction
Working in the infrared has both advantages and disadvantages when compared with the optical region where most supernova spectroscopic work has been carried out until recently. The disadvantages are manifold. The terrestrial atmosphere contains molecules, especially water, which are most efficient absorbers in the infrared. They reduce the transmission of the atmosphere depending on the the wavelength region and the amount of water 1 Infrared here refers to the 1-5/im band.
25 26 Chapter 2. Techniques of IR spectroscopy of SNe vapour between the telescope and the object. Observations are usually carried out in parts of the spectrum, known as windows, with relatively good atmospheric transmission but which nonetheless contain some narrow absorption lines. The atmosphere is also quite a strong emitter in the infrared due to molecular emission, in particular from OH. Another problem for the infrared astronomer is the thermal emission from the telescope and the associated instruments especially at the longer wavelengths. The advantages however are also manifold. The interstellar extinction in the infrared is lower than in the optical allowing observations of objects obscured at shorter wavelengths. The most abundant molecules have strong transitions which are observed in the infrared, and the low-lying forbidden lines of the iron-group elements can also be found in the infrared. Finally a variety of strong allowed lines of hydrogen, helium, calcium, silicon, magnesium and sodium are present in infrared spectra. For the study of supernovae however the infrared has the key advantage of allowing the observer to look ‘through’ the ejecta at useful resolutions (i.e. A A/A 2.2 Devices Very high resolutions can be obtained by using spectrometers such as the Michelson In terferometer but these devices require very bright objects. However, there is little call for such high resolution (~ 10*'Observations of supernovae since they generally exhibit intrin sically broad lines and axe usually too faint anyway for this technique. The common-user infrared spectrometer is usually either a circular variable filter device or a cooled grating spectrometer. 2.2.1 Circular Variable Filter (CVF) A Circular Variable Filter is a disk shaped interference filter having a wedge cross-section so that the transmission wavelength varies almost linearly with azimuth around the disk. By rotating the CVF the wavelength range can be scanned. Until quite recently this was the standard infrared spectrometer available to the community and some of the observations described later used such a device. One of the advantages of the CVF is that it can easily be accommodated within a broad band photometer dewax. However, the attainable Chapter 2. Techniques of IR spectroscopy of SNe 27 resolution is limited to about 100. Moreover, if light landing on the CVF extends over a large region the resolution can bp degraded.^'Hie wavelength coverage of these devices is usually restricted to one octave^due to technical and observational restrictions. 2.2.2 Cooled Grating spectrometers (CGS) Although single element CGS instruments are in use and two-dimensional detector arrays are now being incorporated into these devices, most CGS spectrometers currently available are based on one-dimensional detector arrays. In general in a CGS the image from the telescope is projected on to the grating which is then tilted at different angles in order to send light of a chosen wavelength to a particular detector in the array. The array may have its detectors spaced or closely adjacent. If the array has spaced detectors the grating has to be scanned in order to cover the regions of the spectrum that fall into the gaps between the detectors, and the number of steps taken in the scan depends on the desired sampling. If the detectors are adjacent then a small array will cover only a short portion of the spectrum and many settings will be needed. Order overlap is a problem if the grating is used in high order. The CGS with an array of detectors obviously has a multiplex advantage over the CVF. However, the higher resolutions attainable tend to eliminate any multiplex advantage. Another advantage is that the CGS can provide coverage throughout the near-infrared wavelength region (1-5 pm). 2.3 Observing techniques In infrared spectroscopy normally the study object is fainter than the sky in its vicinity due to emission from the telescope, instrumentation and the earth’s atmosphere. To remove this excess emission a . part of the sky ( one with no stellar object of its own) a few arcseconds away from the object but still in the parent galaxy is selected and its signal is subtracted from that of the field containing the object. This is done by swapping the telescope beam from the object to the blank field and back again at a frequency which depends on the integration time but is usually of order a few Hz. This is called chopping and is usually done by using a wobbling secondary mirror. However this technique will not 1 remove gradients in the background emission -* 28 Chapter 2. Techniques of IR spectroscopy of SNe . 3 - - , Any spatial structure in the emission from the sky, for example from the parent galaxy background can also produce gradients which are not removed. An improvement is effected if the entire telescope is moved so that the background field and the object field alternate. The subtraction then will also remove the first derivative of any background gradient. This is known as nodding or beamswitching. If the dominant source of noise is not too strong (e.g. in the J h E windows) the background subtraction can be done by nodding alone. When using a CGS with spaced detectors it is possible to perform the chopping cycle at each setting of the grating or only at the end of each scan. At long wavelengths where the sky is intense the former is preferable because small changes in the sky dominate. This observing mode is more time consuming than when the scan is performed at each position of the chopper, as the instrumental overheads (e.g. chopper settle time) axe higher. For calibration purposes spectrophotometric standard stars are also observed. A typ ical observing scheme might involve loops, in which the object is first observed, then the standard, and finally a comparison lamp for wavelength calibration. The wavelength range might then be changed and the comparison lamp observed, followed by the standard and then the object. The loop is designed to minimize a) the effects of the telescope on the instrument, b) changes in the sky emission, and c) observing time lost due to the slewing of the telescope from object to standard and back. 2.4 Factors influencing the quality of the data In conventional observations beyond 1 pm almost all measurements are limited by noise from the sky and background emission rather than the instrument (Allen & Barton 1981). However, at high resolutions photon shot noise may become the dominant source of noise. Following the convention of Allen & Barton all background noise (including the telescope) will be referred to as sky noise. Assuming the detectors do not exhibit a 1 // component in their readout noise long integrations will improve on the signal to noise ratio. However since we are chopping the sky is being sampled at a frequency which decreases as we move to longer integrations. Therefore the gain from increasing the signal may be offset in long integrations by the presence of a l / f component in the sky noise which Allen & Barton showed was present throughout the infrared windows. Chapter 2. Techniques of IR spectroscopy of SNe 29 The spectrophotometric quality of the data is affected by variations in the transmission of the atmosphere and by turbulence in the atmosphere which degrades the seeing. Vari ations in the transparency of the atmosphere should not affect the spectroscopic quality of the data as long as the frequency at which the spectrum is sampled and chopped is higher than that with which the transmission of the atmosphere varies. In a spectrum produced by a scanning grating spectrometer, such as the one used for the observations of SN 1987A, the change in the seeing manifests itself as a seeing ripple. As the grating scans and the seeing fluctuates, the same pattern of variation will be observed on each detector at differing wavelengths. This can be removed in the data reduction process since it has a pattern that repeats from one detector to another. Standard data reduction routines can then be used to model the ripple from detector to detector. Thus by dividing the original spectrum by the model seeing ripple spectrum the ripple may be removed from the data. In long integrations the seeing ripple tends to become less noticeable but as mentioned above the sky I ff component then may become important. 2.5 Flux calibration and removal of Telluric features In order to remove the telluric features and calibrate the spectra the standard star is observed at very similar air mass to the object so that telluric features in the object are closely matched in the standard. By dividing the object spectrum by the standard spectrum and then multiplying by a model spectrum of the standard the true spectrum (above the atmosphere) of the object may be retrieved from the data. It is particularly convenient if a suitable standard is available such that it is close to the object in both right ascension and declination. This allows the observer to keep the same standard throughout the observing run as both object and standard will rise and set at more-or-less the same time. There are additional advantages from such a choice of standard. The time taken to move from one field to another is minimized which saves observing time. In the case of SN 1987A a suitable standard was not available. A secondary standard was selected which was used to remove the telluric features but needed in turn to be flux calibrated with respect to a known spectrophotometric standard. Spectrophotometric standards are usually chosen to have an almost featureless spectrum which can easily be modelled. However the secondary standard stars may be less accommodating and care should be 30 Chapter 2. Techniques of IR spectroscopy of SNe taken that spurious features are not introduced in the spectra due to structure in the standard star spectrum that is not modelled. 2.6 Wavelength calibration In addition to determining the wavelength of the features observed in the spectra it is also very important to accurately match the standard and the object spectra in wavelength otherwise spurious features may result when the division is done. The wavelength can be computed from the known characteristics of the CVF being used or from the grating equation in the case of a CGS. However, it is highly desirable to confirm the results of such a calculation by observation of a suitable calibration source. A spectrum of a comparison lamp can be taken just before or after an observation, before any movement of the telescope which might induce flexure in the instrument. In the absence of a lamp suitable for the range under investigation, telluric features of known wavelength in the spectra can also be used for calibration if they occur in strong, featureless regions of the object’s spectrum. If the standard being used is a secondary standard (see above) then absorption lines in its spectrum, e.g. hydrogen lines, of known wavelength may be used to calibrate its spectra. In some instruments it may also be possible to extract the spectrum of the night sky (t.e. the un-chopped spectrum) containing the OH emission bands which may then be used to calibrate the data. Chapter 3 Infrared spectroscopy of Supernovae 1986G and 198TA In this chapter we shall describe the observations of two supernovae. SN 1986G and SN 1987A were of Type la and Type II respectively. They were both observed in the infrared at early and late times thus providing us with a unique set of data. The data are presented and identifications for some of the prominent features are provided. Where possible the conditions in the ejecta that would give rise to the emission observed are also discussed. 3.1 Supernova 1986G Supernova 1986G was discovered before maximum light by R.O. Evans in NGC 5128 (Centaurus A) on May 3.5 (UT) 1986 (Evans 1986). Strong interstellar Na I D and Ca II H & K absorption lines were observed at the redshift of NGC 5128 (Tonry & Strauss 1986; Heathcote, Cowley & Hutchings 1986) suggesting that the supernova was reddened substantially by the dust lane that cuts across this peculiar galaxy. Optical spectra showed absence of hydrogen lines and the presence of the blue shifted absorption at 6200 A (Phillips et al. 1987) thus classifying the supernova as a Type la. The distance to the parent galaxy is uncertain to a factor of two with estimates ranging from 3 Mpc (Hesser et al. 1984) to 6.9 Mpc (Sandage & Tammann 1981). 31 32 Chapter 3. IR Spectroscopy of SNe 1986G and 1987A 3.1.1 The early time observations Although the early time observations of this supernova were not made by the author, for completeness a brief description of the data will be presented. Frogel et dl. (1987) at CTIO and Graham (private communication) at XJKIRT have obtained 1-2.5 pm spectra during the first month after the explosion of SN 1986G (see Frogel et dl. Fig. 2). The spectra axe characterized by a strong thermal continuum (T ~ 104 K) upon which several ‘narrow’ (Vfwhm ~ 10,000 km/s) absorption features are superimposed. The most striking feature however is the very deep, broad (V fwhm ~ 50,000 km/s) absorption feature at 1.3 pm. The ‘narrow’ features in the K window are almost certainly P-Cygni profiles. The trough to peak velocities of these features are of order 7000 km/s which is comparable to the velocities observed in the P-Cygni fines in the contemporary optical spectra. Frogel et dl. identify the Na I 2.207 and 2.337 pm fines as well as the 2.058 pm He I fine to be responsible for the narrow K window features. Graham independently identified the Na I fines but suggests that Mg II at 2.14 pm or Ca II at 2.141 pm may be the cause of the absorption at 2.064 pm. Frogel et dl. also suggest that Mg I at 1.7109 and 1.736 pm matches the features seen in theH window, and Graham suggests that the emission peak at 1.086 pm may be due to a blend of Mg II 1.093 and Fe II 1.086 pm. There does not seem to be a promising candidate for the broad feature at 1.3 pm. If this feature is due to a single atomic transition exhibiting a P-Cygni type profile it’s optical depth must be ~2000 (Graham, private communication). However, no single transition has the correct wavelength coincidence and cosmic abundance to be a plausible candidate. It seems more likely that the feature is due to a blend of several transitions lying very close in wavelength. In the spectra taken at UKIRT coverage was also obtained in between the atmospheric windows revealing an asymmetric deep and broad absorption feature at 1.9 pm. Graham suggests that this feature may be due to the Ca I multiplets with transitions near 2 pm. 3.1.2 The late time observations SN 1986G presented a unique opportunity to acquire late time spectra of a Type la supernova. The aim of the programme was to detect the 1.64 pm feature which had been seen in the spectra of SN 1983N (Graham et al. 1986). This feature was identified by Graham et al (1986) to be emission from [Fe II] whereas Oliva (1987) suggested that Chapter 3. IR Spectroscopy of SNe 1986G and 1987A 33 Table 3.1: Log of observations of SN 19866G Date Telescope Time allocated Time on SN 1986G Observers Jan 22 ’87 AAT 1/2 night 0 Allen Feb 15 ’87 AAT 1/2 night 0 Allen Mar 13 ’87 AAT 1 night < 1 night Allen Mar 14-16 ’87 CTIO 3 nights < 3 nights Graham, Meikle Apr 14 ’87 AAT 1/2 night 6 hours Allen, Meikle May 1-4 ’87 IRTF 4 x 4 hours 2 x 4 hours Graham, Spyromilio May 14 ’87 AAT 1/2 night 1.5 hours Allen June 19-21 CTIO 3 nights 3 nights Andrews Mar 7-8 ’88 AAT 2 nights 1 night Allen, Spyromilio it may be [Si I] that produced the emission. The iron, if present, would result from the radioactive decay of 56Ni formed in the explosion. It’s presence in the spectra of a Type la supernova would be strong evidence for the radioactive decay powering mechanism for this class of supernova. To investigate this we were allocated time on the AAT, CTIO (4m) and the IRTF telescopes. All spectra were taken in the H atmospheric window and the observations are listed in Table 3.1. All instruments used are common user and are described in detail in the respective telescope manuals. We shall describe briefly each instrument and the mode in which it was used as well as the conditions during each observing run so as to make comparisons between different runs possible. 1. The FIGS spectrometer contains an array of 16 InSb detectors sensitive over the range 1-5 pm. A choice of gratings offers resolutions (A/A A) of about 500 or 1500. In both cases the resolution is somewhat degraded if large entrance apertures are used. The array is spaced and so the grating is scanned over 13 increments in order to fully sample the spectrum or over fewer increments if full sampling is not required. The grating was scanned before the chopper was moved to observe to sky. The low resolution grating was used in the oversampling mode (13 increments) for the observations on the 13th of March 1987. The entrance aperture was 5.9 arcsecs. Problems with the computers and the chopping 34 Chapter 3. IR Spectroscopy of SNe 1986G and 1987A secondary lost some of the night. The observations on the 7th of March 1988 were carried out in the under sampling mode (7 increments) with the low resolution grating. The entrance aperture used was 3.5 arcseconds. These observations were being made almost two years after the discovery of SN 1986G and as such we were unsure as to whether we would be able to detect any emission from the supernova at all. We therefore chose not to oversample in order to improve S/N. This was acceptable since the emission line we were attempting to detect was expected to cover at least a few resolution elements the fact that we had not oversampled would not put its detection in doubt. The smaller entrance aperture was selected in order to reduce the galaxy background. 2. IRS is a grating spectrometer which offers two resolution modes. It has an 8 element InSb array sensitive from 1-5 pm. An optional monitor channel which is fed through a beam splitter is also available. This monitor channel can be used to reduce the effects of seeing, guiding errors or poor atmospheric transmission by dividing the other detectors’ output by the monitor’s. This instrument was used on the 14th and 16th of March 1987 in the low resolution fully sampling mode. In this mode two interleaved spectra are taken. A number of chopping and beamswitching cycles were performed for each segment of the spectrum before the grating was moved. The effect of observing in this mode can be seen in the size of the error bars which alternates from point to point with a two point cycle. Problems with computers at the telescope lost some observing time during this run. Further runs on the 19th, 20th and 21st of June 1987 were carried out using this instrument without fully sampling. In this mode the instrument takes two spectra side by side, the second spectrum starting where the first ended. Again a full integration cycle was performed for each spectrum before the grating was moved. 3. IRPS contains a standard CVF with a resolution of 100. Data were taken with this instrument on the 14th of April 1987 and the 14th of May 1987 both times using the 3.5 arcsec circular aperture. During both runs cloud prevented us from observing the supernova for the whole duration of the run. In particular in April we observed the supernova for 6 hours whereas in May only 1.5 hours of coverage was obtained. 4. CGAS is a grating spectrometer with two available gratings and uses a 32 element InSb array. It also has a fixed entrance aperture of 2.7 arcseconds diameter. This instrument was used for two half nights the 1st and 2nd of May 1987 in its low resolution mode. The weather was poor during the runs and as the final sensitivity of the instrument was a Chapter 3. IR Spectroscopy of SNe 1986G and 1987A 35 factor of five lower than expected no useful data were collected. For all the observations standard data reduction techniques were used. The AAT data were reduced by David Allen at the AAT and the July 1987 CTIO data by Phil Andrews at CTIO. In all observations the spectrophotometric standard was BS 4903 a G3V dwarf with an H magnitude of 4.58 (Allen & Cragg 1983). When a black-body fit to the theoretical continuum from BS 4903 was used its temperature was assumed to 5700 K. (The STARLINK package Figaro does not use a black-body to fit the continuabut a model spectrum based on observations of spectrophotometric standard stars). The March 1987 data from CTIO were reduced by the author at ICSTM with the use of routines developed for this purpose. The data from runs which had compatible apertures and chop-throws were coadded. The late time data The spectra can be seen in Figs 3.1 through 3.7. In order to compare data taken with such a wide range of instruments we have binned the spectra from IRPS to the lowest effective resolution (bin width/A) in the data set, namely that of IRS. A weighted (according to the individual error bars) mean and error was calculated in the binning process. The binned data can be seen in Figs 3.8 and 3.9. From maps taken at the AAT and photometry taken at CTIO the background con tinuum due to the galaxy was calculated for different chop directions, chop throws and apertures. The calculated backgrounds for each run can be seen in Table 3.2. The unbinned IRPS spectra show little evidence for any significant feature. From inspection of the binned spectra however it is possible that a line may be present at 1.64 pm in both spectra taken with IRPS at the AAT. The significance of the highest point in the binned IRPS spectra taken in April 1987, relative to the calculated continuum that should be present in the spectra, is 2.8 o. In contrast the data taken with IRS in March 1987 do not show any such feature but confirm the choice of galaxy continuum for small apertures. The contemporary to the IRS data, FIGS data show no evidence for a feature at 1.65 pm. It is clear that in March 1987 both FIGS and IRS while detecting the background from the galaxy did not detect the supernova. In April and May 1987 IRPS detected the background but also possibly the supernova. The July 1987 data from 36 Chapter 3. IR Spectroscopy of SNe 1986G and 1987A Table 3.2: Apertures chops and calculated backgrounds for spectra of SN 1986G Date Telescope Instrument Aperture f Chop Background f Mar 13 ’87 AAT FIGS 5” .9 □ 12” NS 8 (2) Mar 14 ’87 CTIO IRS 6” □ 13” NS 9(2) Mar 16 ’87 CTIO IRS 3” □ 7” NS 1.5 (.2) Apr 14 ’87 AAT IRPS 3”.5 0 12” NS 2 (.5) May 14 ’87 AAT IRPS 3” .5 0 12” NS 2 (.5) June ’87 CTIO IRS 6” □ 9” NS 9(2) March 7 ’88 AAT FIGS 3”.5 □ 12” NS 2.5 (.6 ) Notes for table 3.2 f Aperture shape: D:square, 0 :circular $ Backgrounds in 10-13 erg/s/cm 2//mi, (rterror) IRS are consistent (in part) with the detection of the background alone whereas the FIGS data from March 1988 do not show significant evidence for the continuum from the galaxy. This may be due to the small aperture used for those observations which limited the flux entering the instrument. A longer integration in that case may have shown the galaxy continuum. However since the object of the observations was not the galaxy but the supernova which was clearly not detected the observations were halted. If the ‘feature’ observed with the IRPS at the AAT were real what could have caused it? The question arises whether the spectra taken were of the same part of the sky and if so whether they axe the result of instrumental effects? Since for all the observations discussed here the supernova was too faint to guide on, we offset from a nearby star and guided on that star. We have checked that the offset used in all observations was of the right size and in the right direction. Problems with the guiding though unlikely would be of order 1-2 arcseconds and as such would not have prevented the detection of the supernova in the 6 arcsecond aperture which was used both at the AAT and at CTIO in March 1987. Therefore we believe that at least some of the spectra have to be of the same part of the sky. A leak in the CVF filters in IRPS around the 1.65 pm region could Chapter 3. IR Spectroscopy of SNe 1986G and 1987A 37 also have produced the observed spectra. However, tests for this effect did not show any evidence for such a leak. A temporal change in the supernova in a timescale of tens of days is unlikely. The 1.65 pm line if present is due to collisionally excited lines of singly ionised iron or neutral silicon. In order to make such lines ‘switch on’ a rapid recombination to those ionisation states would be required. Theoretical models of supernovae (see chapter 4) indicate that the timescale for the change in the ionisation state in the ejecta is of order a few months. Therefore it seems unlikely that this could be the cause of the observed spectra although it cannot be totally discounted. Alternatively the emergence of the lines could be the result of an optical depth effect such that the emitting region remains hidden until the photosphere recedes deeper into the ejecta. However since no significant continuum (above the galaxy background) from the supernova was detected in March 1987 it also seems unlikely that such an optically thick region existed. We believe that no significant detection of an emission line was made. However, the spectra may be used to place limits on the mass of Fe II (or Si I) present in the ejecta of SN 1986G. In chapter 4 we shall compare the data with predictions from theoretical models. 38 erg/cm2 /s//im) F\ (10 —'132 erg/cm /s/yum) Figure 3.2: SN 1986G CTIO 4m March 14th 1987 14th March 4m CTIO 1986G SN 3.2: Figure Figure 3.1: SN 1986G AAT March 1987 March 1986GAAT SN 3.1: Figure aeegh (yt/m) wavelength Chapter 3. IR Spectroscopy of SNe 1986G and 1987A and 1986G SNe of Spectroscopy IR 3. Chapter Chapter 3. IR Spectroscopy of SNe 1986G and 1987A and 1986G SNe of Spectroscopy IR 3. Chapter cm2/s/jLim) Fx (10 13 erg/cm2/s//im ) Figure 3.3: SN 1986G CTIO 4m March 16th 1987 16th March 4m CTIO 1986G SN 3.3: Figure Figure 3.4: SN 1986G AAT April 1987 April AAT 1986G SN 3.4: Figure 40 hi O Figure 3.9: Binned spectrum of SN 1986G AAT May 1987 1.8 wavelength (ptm) Chapter 3 . IR Spectroscopy of SNe 1986G and 1987A 43 3.2 Supernova 1987A The discovery of supernova 1987A by Ian Shelton (Madore & Kundel 1987) in the Large Magellanic Cloud prompted an unprecedented number of observations of a SN event. We have used the Anglo-Australian Telescope to observe SN 1987A in the J, JJ, AT, L and M near infrared atmospheric windows at regular intervals after the explosion. Here we present only data from the first year after the explosion of SN 1987A. 3.2.1 The observations For all the data presented here we used the infrared spectrometer FIGS on the AAT. The spectra were always oversampled (13 increments). At low resolution most of the J, H or K atmospheric windows, or more than half of the L window, can be covered at a single setting; at high resolution five or six observations must be spliced together to cover each one of the atmospheric windows. A log of the observations can be found in Table 3.3. Owing to the fact that FIGS had been designed and optimised for faint objects, the techniques required to attain data of high signal-to-noise had not been developed. SN 1987A provided an incentive in this regard, and examination of the data presented here reveals a progressive improvement in their quality with time despite the fact that the supernova was fading. The technique developed, and adopted in the later spectra, is as follows. 1. We used an f/36 chopping secondary for sky subtraction. The 13 grating steps were scanned before the chopper changed state to allow the sky spectrum to be measured. This option was chosen because modulation by seeing was always a larger source of noise than fluctuations of the sky, even during daylight and twilight observations. Sampling times per grating position were typically 40-100 ms, depending on the seeing. 2. To remove telluric absorptions we observed a standard star positioned as closely as pos sible in time and zenith distance to the supernova. Early attempts proved unsatisfactory due either to variations in the telluric bands or to the difficulty of removing the absorp tion lines in the standard. Eventually we established the A7 dwarf BS 2015 as our local standard. It has the great advantage of proximity to SN 1987A, but has strong, broad 44 Chapter 3. IR Spectroscopy of SNe 1986G and 1987A Table 3.3: Log of observations of SN 1987A Date Epoch Resolution Local Observing conditions (1987/88) (days) Standard Mar 13 18 Low BS 1294 photometric Jun 13 110 Low BS 2256 photometric, poor seeing (>3 arcsec) Jun 15 112 High (J) BS 2256 photometric Sep 3 192 High BS 1443 photometric, mediocre seeing (2-3 arcsec) Nov 5 255 Low BS 2015 photometric Dec 4 284 High BS 2015 non-photometric Feb 7 349 High BS 2015 photometric but very poor seeing 1. ‘High’ and ‘low’ resolution refer to A/AA of 1500 and 500 respectively. All the L window spectra are low resolution. 2. Other spectra obtained but not shown are: (a) a low-resolution J spectrum on day 110 (b) an incomplete, high-resolution J spectrum on day 255 (c) a high-resolution M window spectrum on day 255, obtained in poor conditions (d) low-resolution JHKL spectra on day 283. Chapter 3. IR Spectroscopy of SNe 1986G and 1987A 45 hydrogen absorption lines across which the continuum (and telluric absorption features) must be interpolated before division. BS 2015 was in turn flux-calibrated against BS 1294, a G dwarf spectrophotometric standard (Allen & Cragg 1983). A further disadvantage of BS 2015 is its weakness at the longer wavelengths, so we instead used a Carinae as the local standard for observations in the 4.5 to 4.8 pm region, although we were forced to revise our definition of ‘local’ in this case. 3. Whenever time permitted we observed first at low resolution using the largest available aperture (5.9 arcsec square) to derive spectrophotometric data. We then covered the spectral windows at high resolution and with a smaller aperture (usually 3.5 arcsec), allowing a small overlap between portions. The low-resolution spectrophotometry was then used to verify the splicing of the high-resolution data and if necessary to scale the result. The high-resolution spectra were always of higher quality, both because oversampling the lines shows better the effects of blends and because more data points could be measured between strong telluric absorptions. When non-photometric conditions or very poor seeing introduced additional uncertainty we scaled the spectra to match appropriate broad-band photometry from the SAAO (Menzies et al. 1987; Catchpole et al. 1987, 1988; Whitelock et al. 1988). 4. Data reduction was undertaken at the ICSTM and the AAO using some of the routines of the Figaro package supported by STARLINK. The techniques were standard, except that for the earlier data no comparison spectra were available to provide wavelength cali bration. Instead, recognisable telluric features in absorption (in the programme objects) or emission (in the sky), and absorption lines in the standard star provided confirmation of the wavelengths computed from the grating equation. (See discussion in chapter 2). The limitations of the data should be noted. The resolution of FIGS depends on the entrance aperture or, more typically, on the seeing and telescope tracking during the observation. At worst the nominal values of 1500 and 500 maybe reduced to about 1000 and 350. The uncertainty in the absolute flux calibration is 5-10% and possibly higher in regions of high atmospheric absorption. Wavelengths are generally reliable to ±1 pixel; the pixel size ranges from 0.00025 pm at J in high resolution to 0.003 pm at L in low resolution. Telluric absorption has been removed extremely effectively except for a few very deep features. In particular, water vapour bands cause irregularities in the red wing of the 1.13 pm line, and together with methane, throughout the entire spectral region 2.9- 46 Chapter 3. IR Spectroscopy of SNe 1986G and 1987A 3.3 pm. In the M window, incomplete removal of telluric CO R(0) and water probably account for most of the split in the Pfund/? feature, with its steep sides and atypically narrow width also having been caused by telluric absorption, mainly CO R(l) and P(l). However, we do not rule out the possibility that a contribution to the split in the Pf/3 line came from absorption by cold CO located in the wind produced by the progenitor while in its RSG phase. For the J window observations we used the gratings in fourth order, and there is possible slight contamination by fifth order light at the long wavelength end and by third order at the short. 3.2.2 Results For a full description of the first year’s data see Meikle et al. (1989). Here we present only an overview of the infrared spectra of SN 1987A up to day 349 after the explosion accompanied by an analysis of what we feel are some of the more exciting aspects of the spectra. In particular the analysis of the molecular emission which was detected in SN 1987A will form a separate chapter and has been omitted from this section. The spectra can be seen in Figs 3.10 through 3.34. i /w o f acquired *K\s 9 ooo f/rzuU>L^ c\ G=> m p py^S^MfeJK'b«\ ttrt»5 uoo(K* v f ‘Ss Overview of the infrared spectra 3-10 —3.3^ ^ Table 3.4: Evolution of the continuum in the spectra of SN 1987A Epoch Continuum (10-10 erg/s/cm 2//im) (Days) 1.25/zm 1.65/nn 2.2pm 3.6pm 18 3000 1200 450 120 110/2 7100 3300 1400 470 192 1200 550 240 65 255 360 250 100 25 284 200 160 70 20 349 110 65 27 9 4800 K. This is again in agreement with the value derived from photometric studies by Catchpole et al. (1987). The P-Cygni troughs were weaker than those observed on day 18. On day 110 the lines accounted for ~ 10% of the total flux. By day 192 the P-Cygni troughs had essentially disappeared and almost 40% of the total flux was originating in the broad emission lines. One of the most difficult problems faced in the analysis of the spectra was the placement of the continuum. This was done by fitting a polynomial function through what seemed to be parts of the spectrum which did not have any specific spectral feature. On average 3-4 points per spectrum were selected. The values for the continuum levels at the centre wavelengths of each atmospheric window can be found in Table 3.4. The evolution of the continua was qualitatively similar to that observed for the bolo- metric light curve (Catchpole et al. 1987,1988). Between days 18 and 110 the continuum flux increased about threefold and then fell to approximately 20% of its value on day 192. The subsequent decline can be reasonably represented by an exponential with an e-folding time ranging from 65 to 75 days depending on the atmospheric window in question. The evolution of individual line intensities was more complicated than that of the continuum. Most lines declined slower than the continuum although this was almost certainly due to the fact that the lines after day 110 were no longer driven by continuum photons. 48 Chapter 3. IR Spectroscopy of SNe 1986G and 1987A Line identifications The line identifications which we present in Table 3.5 were established through the use of Branch’s (1987b) guide to line identification in the early supernova. For the later spectra no detailed predictions have yet been published. Fransson & Chevalier (1987) have produced model spectra for the late time supernova though they did not include in these models a large number of atomic species. In fact in the 1-5 pm region of the spectrum therefore we used atomic data tables (Wiese et dl. 1966, 1969; Kurucz & Peytremann 1975) to predict the intensities of features that may be present in the infrared spectra of a Type II supernova. The candidate transitions for the interpretation of the data tended to be low excitation transitions from low ionisation ions and lines that may be strong in recombination spectra. Additionally the identifications where possible were cross-checked with features present in the contemporary optical spectra taken at SAAO (Menzies et dl. 1987; Catchpole et al. 1987, 1988; Whitelock et al 1988) or taken from the AAO archive. The late spectra axe quantitatively described in Table 3.6. The line identifications can also be seen in some of the figures. We shall now briefly discuss some of the more interesting features of the spectra. Chapter 3. IR Spectroscopy of SNe 1986G and 1987A 49 Table 3.5: Features observed in the day 18 spectra of SN 1987A ^peak Intensity Vfc/ue Vtrough Vp eak Vred Identification ^rest ( p m ) lO" 1 0 (km/s) (km/s) (km/s) (km/s) ( p m ) erg/s/cm2 f 1.0683 ? | 1.0685 1.068 lies in the C I 3s3 P°-3p3D (1) j 1.0691 Py trough | 1.0707 j 1.0729 [ 1.0754 f 1.0914 Mg II 3d2 D-4p2 P° | 1.0915 1.095 6 . 8 + |_ 1.0952 Sr II 4d2 D-5p2 P° 1.0915 + -6700 -4300 +300 +4200 Paschen y 1.0938 f 1.1747 1.16- ~ 4 C I 3p3 D-3d3 F° | 1.1754 1 . 2 2 + 11.1802 Ca II 5s 2 S-5p2 P° (5) f 1.1836 [ 1.1947 1.283 1 2 -5100 -3700 +300 +4500 Paschen /? 1.2818 1.64 ~ 1 -4000 -2500 - 0 blend Brackett 12 1.6407 1.684 - 0.5 blend -3300 +600 blend Brackett 11 1.6806 1.737 ~0.5 blend -2800 + 1 0 0 +2400 Brackett 10 1.7362 2.165 1 . 0 -5500 -3500 -70 +2550 Brackett y 2.1655 3.739 0.35 -4100 -2600 -50 +2900 Pfund y 3.7395 4.053} > 2 . 0 -7000 -4200 +150 > +3600 Brackett a 4.0512 Notes for table 3.5: v&/uc 5 ^trough, Vpeak> vre(f give the respective velocities w.r.t. the terrestrial rest frame of the blue edge, absorption trough, emission peak and red edge of the P-Cygni profile. $ Partial spectral coverage of feature. The error in line intensities is typically ± 1 0 %. Presence of before the intensity indicates that the value is much less certain due to problems with blending, continuum placement or telluric absorption. 50 Chapter 3. IR Spectroscopy of SNe 1986G and 1987A Table 3.6: Line identifications for the day 110-349 spectra of SN 1987A Epoch ^peak Intensity Identification rest 0 ^ o (days) (/im) 1 ( p m ) erg/s/cm2 112 1.082 ~35 He I 2s3S-2p3P° 1 .0 8 3 0 192 1.087 50 (1) 255 1.087 35 + 284 1.085 20 [S I] 3p43P-3p41D 1.0820 349 1.087 15 (IF) 112 1.095 200 192 1.095 80 H I Paschen y 1.0938 255 1.095 30 + 284 1.097 20 [Si I] Sp^-Sp^ 1.0991 349 1.098 10 (2F) 112 1.131 ~35 192 1.132 36 0 I 3p3P-3d3D° 1.1287 255 1.130 28 + 284 1.132 25 [S I] 3p43P-3p41D 1.1306 349 1.131 12 (IF) 112 1.14 ~ 30 192 1.14 ~ 15 Na I 3p2P°-4s2S f 1.1382 255 1.14 ~ 8 (3) [ 1.1404 284 1.14 ~ 4 349 1.14 ~ 3 112 1.166 ~ 15 192 1.166 ~ 7 255 1.166 - 4 ? 284 1.166 ~ 3 349 1.167 - 1.5 112 1.178 ~ 4 192 1.177 ~ 3.5 f 1.1690 255 1.178 ~ 2.5 K I 42P°-3d2D | 1.1770 284 1.176 - 1.5 (6) [ 1.1773 349 1.177 ~ 1 Chapter 3. IR Spectroscopy of SNe 1986G and 1987A 51 Table 3:.6 (continued). Epoch ^peak Intensity Identification ^rest i o (days) (/im) o (pm) erg/s/cm 2 112 1.18 t 192 1.188 ~ 12 255 1.188 ~ 6 Mg I 3p 1P°-4s1S 1.1828 284 1.188 ~ 2.5 349 1.19 ~ 1.5 112 1.20 t 192 1.202 20 f 1.1984 255 1.201 12 | 1.1992 284 1.202 10 j 1.2032 349 1.203 5 Si I 4s3P°-4p3D | 1.2104 112 1.228 5 (4) 1 192 1.227 1.5 | 1.2271 255 1.227 0.8 [ 1.2396 284 1.228 0.5 349 1.230 0.1 192 1.258 3.6 [Fe II] a6D9/2-a4D7/2 1.2570 255 1.259 4.5 + 284 1.259 3.3 K I 4p2P°-5s2S f 1.2432 349 1.260 2.9 (5) [ 1.2452 112 1.285 100 192 1.285 90 H I Paschen (3 1.2818 255 1.285 44 + 284 1.285 25 A11 4s2S-4p2P° \ 1.3123 349 1.286 13 (4) [ 1.3151 192 1.488 ~ 0.4 Mg I 3d 3D-4f3F° 1.4878 349 1.490 0.2 110 1.505 - 15 192 1.505 3.2 255 t X Mg I 4s3S-4p3P° 1.5031 284 1.506 1.5 349 1.506 1.3 52 Chapter 3. IR Spectroscopy of SNe 1986G and 1987A Table 31.6 (continued). Epoch ^peak Intensity Identification ^reit (days) (pm) 10-io (/im) erg/s/cm 2 110 1.548 ~ 4 192 1.55 1.4 [Fe II] a*F9/2-a4D5/2 1.5335 255 1.55 1.2 + 284 1.55 1.7 [Co II] a 5F5-b3F4 1.5470 349 1.55 1.1 192 1.575 - 1 255 1.574 ~ 0.5 H I Brackett 15 1.5701 294 1.575 ~ 0.5 349 1.576 ~ 0.1 110 1.591 <"w> 8 192 1.594 4.5 H I Brackett 14 1.5881 255 1.592 2.5 + 284 1.595 2.0 Si I 4s1P°-4p1P 1.5888 349 1.597 0.9 110 1.61 ~ 9 [Fe II] a4F7/2-a4D3/2 1.5994 192 1.61 4.2 + 255 1.604 ~ 3 [Si I] 3Pi-1D2 (0.01F) 1.6068 284 1.608 2.4 •f 349 1.608 1.3 H I Brackett 13 1.6109 110 1.644 11 H I Brackett 12 1.6407 192 1.643 9.2 + 255 1.644 6.0 [Fe II] a4F9/2-a4D7/2 1.6440 284 1.648 5.6 + 349 1.648 4.3 [Si I] 3P2-1D2 (0.01F) 1.6454 Chapter 3. IR Spectroscopy of SNe 1986G and 1987A 53 Table 3 .6 (continued). Epoch ^peak Intensity Identification Arest O H o (days) ( p m ) ( p m ) erg/s/cm 2 110 1.683 ~ 5 192 1.683 6.0 H I Brackett 11 1.6806 255 1.681 3.9 + 284 1.686 3.1 [FeII]a Table 3.6 (continued). Epoch ^peak Intensity Identification ^re»t (days) (pm) 10-io (^m) erg/s/cm 2 110 2.17 16 192 2.172 9.5 255 2.171 6.4 H I Brackett j 2.1655 284 2.171 6.2 349 2.173 1.7 110 2.20 ~ 3 192 2.205 2.5 255 2.205 1.4 Na I 4s2S-4p2P° f 2.2056 284 2.21 1.1 [ 2.2084 349 2.212 0.5 110 2.26-$ > 50 192 2.264-J > 40 CO Av = 2 255 2.263-J > 20 4* 284 2.263-t > 15 CO+Au = 2 349 2.264-J > 4 192 3.054 5 284 3.045 1.7 H I Pfund e 3.0383 349 3.05 0.6 192 3.15 ~7 [Ni I] a3D3-a1D2 3.1191 284 3.14 ~2 + unidentified feature 349 3.13 ~ 1 (which fades by day 349) 192 3.318 5.5 255 t i 284 3.303 1.8 H I Pfund 6 3.2960 349 3.31 0.35 192 3.409 3.5 255 3.404 1 ? 284 3.395 1 349 3.390 0.5 Chapter 3. JR Spectroscopy of SNe 1986G and 1987A 55 Table 3.6 (continued). Epoch ^peak Intensity Identification Arest (days) H 10-io (H erg/s/cm 2 110 3.516 2.2 192 3.533 1 255 3.53 0.35 ? 284 3.526 0.3 349 3.53 0.17 110 3.740 6.5 192 3.757 5.5 255 3.746 2.5 H I Pfund 7 3.7395 284 3.738 2.4 349 3.746 1.0 110 ~ 3.89 ~ 2 192 3.88-3.93 2.0 255 3.88-3.92 1.0 CS Av = 2 284 3.86-3.91 0.5 349 3.88-3.91 0.25 110 4.056 18.5 192 4.072 16 H I Brackett a 4.0512 255 4.061 7.4 + 284 4.056 t SiO Av = 2 349 4.060 5.9 192 4.67 ~ 2 H I Pfund /3 4.6525 192 4.55-4.75} >16 CO Av = 1 (extends beyond both limits of spectral coverage) Notes for table 3.6: f Mg I feature heavily blended with Si I. } Partial spectral coverage of feature. The error in line intensities is typically ±10%. Presence of before the intensity indicates that the value is much less certain due to problems with blending, continuum placement or telluric absorption. 56 Chapter 3. IR Spectroscopy of SNe 1986G and 1987A Hydrogen yvj0>f. r U(U)aew speorroL a w ** gv^ _ C ^l I SS'v During the whole of its first year the infrared spectrum of SN 1987A was dominated by the presence of transitions from the Paschen, Brackett and Pfund series of hydrogen. In fact these are the first ever observations of transitions from these series in the spectra of a supernova. In the day 18 spectra all these lines show ‘P-Cygni’ profiles resulting from the scattering of radiation from the photosphere by the rapidly expanding envelope. We shall discuss the interpretation of these lines in chapter 4. Of particular interest, however, is the observation that the velocity of the blueshifted trough decreases as we go to higher series. In the classical scattering models for P-Cygni line profiles this velocity corresponds to the velocity of material at the photosphere. The lower velocities may be the result of lower optical depths for the transitions involving the upper levels of the hydrogen atom. However, the velocity of the Pf 7 trough is lower than that calculated for the surface of the photosphere. By fitting black-body spectra to the photometric data, Menzies et al. (1987) calculated that on day 18 the ‘photometric’ radius of the photosphere was 6 x l014 cm. For that date this corresponds to a velocity at the photosphere of 4000 km/s. However the Pf (3 line exhibited an absorption trough moving at 2600 km/s. This problem and its implications are discussed in the next chapter. In the later spectra the blueshifted absorptions characteristic of the P-Cygni profiles faded and the hydrogen exhibited pure emission features. The ratios of the intensities of the hydrogen lines can be seen in Table 3.7 where they are compared with theoretical Case B values from Hummer & Storey (1987). The spectrum loosely resembles that of a Case B recombination spectrum. This has also been pointed out by Oliva et al. (1988) and Catchpole et al. (1988).The observation of the Ly (3 pumped O I line at 1.13 pm suggests that the hydrogen recombination spectrum may in fact be case C. We must await theoretical predictions of such a spectrum to compare with our observations. The initial explosive heating of the envelope would not maintain the ionisation of hydrogen for such long periods. An additional source of energy is required. This source is probably the radioactive decay of 56Ni—*56Co—>56Fe. W opV^usi^i^j ^ \&L. 0 J 4A avo*V o x Chapter 3. IR Spectroscopy of SNe 1986G and 1987A 57 Table 3.7: Hydrogen line ratios in the spectra of SN 1987A relative to Paschen j3 Line Case B day 192 day 255 day 284 day 349 P P 1.000 1.00 1.00 1.00 1.00 Bra 0.506 0.18 0.17 t 0.45 Br7 0.175 0.11 0.15 0.25 0.13 BrlO 0.059 0.07 0.08 0.12 0.12 Pf/? 0.103 ~0.02 — — — Pf7 0.068 0.06 0.05 0.08 0.07 | Partial spectral coverage of feature. Helium Helium was not observed in the day 18 spectra, but by day 112 a deep absorption feature was present at 1.065 pm as was a weaker absorption feature at 2.04 pm. We interpret these features as the blueshifted absorption features in P-Cygni lines from the 1.083 ( 2s3S- 2p3P°) and 2.058 pm (2s1S-2p1P°) transitions in He I. We suggest that these lines formed as a result of resonant scattering from the enhanced 2s3S and 2s 1 S populations arising from the recombination of He II. This mechanism was also suggested by Elias et al. (1988) and Graham (1988) has shown that the X-rays and 7-rays emerging from the decay of the radioactive 56Ni are an adequate source of ionising radiation to maintain the required ratio of He II/He I. In the recombination of He II most electrons follow the triplet cascade and reach the metastable 2s3S level. The 2s3S to ground transition is forbidden (A = 1.3xl0 ~4 s-1) and therefore an enhanced population is established. Additionally the 2p 3P° to ground transition is also forbidden. Thus any electron which is collisionally or radiatively excited from the 2s3S to the 2p3P° level will almost certainly decay back to the 2s3S level emitting a 1.083 pm photon. The 2.058 pm transition is not as simple. Although the 2s1 S to ground transition is also forbidden and an enhanced population may be established^ that level, the upper level p p 1? 0) can decay to ground. Thus the relative strengths of the two 58 Chapter 3. IR Spectroscopy of SNe 1986G and 1987A features may be understood in terms of the fact that whereas almost all 1.083 pm photons which are scattered by the helium will result in another 1.083 pm photon not all 2.058 pm photons will have the same fate. Some will be converted to TJV photons depending on the optical depth of the 2p1P° to ground transition. The relative strengths of the absorption and emission features in the P-Cygni feature exhibited by the 1.083 pm line indicates that some of the emission is due not to resonant scattering but to recombination and collisional excitation. Oxygen On day 112 we observed a weak, broad feature at ~1.13 pm. By day 192 it had evolved into a strong, narrow (Vfwhm ~ 2100 km/s) emission line which persisted for the remainder of the first year. We identify this as the 1.1287 pm (3p3P-3d 3D°) transition of O I. This line forms part of the Bowen fluorescence cascade which is pumped by Ly /3 photons. Support for this identification is provided by the lack of a blueshifted absorption trough. However, the subsequent 8446 A transition remains undetected in the optical spectra, presumably due to strong blanketing by the Ca II absorption trough at about the same wavelength. McGregor (1988) has in fact proposed an alternative identification for the 1.13 pm feature to explain this discrepancy. He suggests that the [S I] transition at 1.1306 pm is responsible for this feature. However if this were the case the 1.082 pm [S I] line would also be present and have been X3.5 stronger. We therefore favour the O I identification. A notable aspect of the 1.13 pm feature is its exceptionally narrow width. If the expan sion of the ejecta were homologous we should expect such a narrow width to correspond to a distribution of the emitting material dose to the centre of the supernova. However, for the Bowen fluorescence mechanism to operate we require not only a laxge flux of Ly (3 photons but a coincidence in velocity space of the pumping hydrogen and pumped oxygen. In fact the velocities of the hydrogen and the oxygen can differ only up to 15 km/s before the pumping is switched off. Only regions where both conditions are met will be emitters of this transition. An immediate condusion from the presence of this line is that there has been substantial mixing in the ejecta of SN 1987A. Chapter 3. IR Spectroscopy of SNe 1986G and 1987A 59 Silicon From day 112 onwards a broad feature was present in the spectra between 1.16 and 1.22 pm. The long-wavelength part of this feature we identified as transitions from mul- tiplet 4 in silicon. These transitions are the strongest in the allowed infrared spectrum of silicon. If this identification were correct we would expect other optical transitions to be present. These transitions (around 3900 A ) are not seen but that part of the supernova’s spectrum suffers from severe line blanketing. Thus the Si I identification is not conclusive. Silicon also has forbidden transitions in the IR which it has been predicted (Fransson & Chevalier 1987) would be strong in the late time spectra. Although these lines lie in regions of the spectrum contaminated by emission from other species, nevertheless, their presence is apparent at certain epochs. The flux measured for the hydrogen Pa (3 line lies above the expected Case B values. We interpret this as the result of blended emission from [Si I] 1.099 pm. In the H window on day 192 all the emission at the wavelengths where emission from [Si I] would be expected can be accounted for by emission from [Fe II] and the Brackett series. However at later dates excess emission is found to be present in the 1.645 and 1.607 pm regions and we identify this excess as emission from [Si I]. Iron One of the most important results of the infrared spectroscopy of SN 1987A has been the unambiguous detection of emission from iron. From day 192 onwards the 1.257 pm [Fe II] line has been present in the spectra. The ‘sister’ line (same upper level) at 1.644 pm was also detected as were other [Fe II] lines in the H window. As mentioned above the H window lines were contaminated by Brackett and (sometimes) [Si I] emission. However since the 1.257 pm line is much less subject to blending we can infer the mass of the iron in the ejecta from its intensity provided we know the temperature and density of the thermal electron gas which is exciting the upper level of the line (c.f. Graham et dl. 1986). Oliva et dl. (1988) obtained an electron temperature of 4000 K from the intensity ratio of the 5530 A to 7160 A lines of [Fe II] around day 225, and Aitken et dl. (1988) using the [Ni II] 6.6 and 10.68 pm lines obtained 3250 K around day 260. In the high density limit and with a temperature of 4000 K our 1.257 pm line intensity on day 192 corresponds to 0.04 M®, in agreement with Oliva et dl. (1988). The mass of Fe II is at least an order of 60 Chapter 3. IR Spectroscopy of SNe 1986G and 1987A magnitude higher than the total iron mass which would have been ejected in the explosion (Weaver et al. 1978; Woosley 1988a) for typical LMC abundances, but is comparable to the 0.06 M© of iron (from the radioactive decay of 0.08 M© of 56Ni) which would have been present in the ejecta around that time. Thus one of the main goals of this work - to detect the decay products of the 56Ni in the spectra of supernovae - has been achieved. Cobalt The strongest [Co II] line should be at 1.547 pm. (a5Fs-b3F4), although it would be blended with both an [Fe II] line and several Brackett lines (which here overlap to form a weakly structured continuum). The cobalt/iron blend was probably present at all epochs from day 110, and we estimate that on day 349 the [Co II] line accounted for about 70% of it. The companion line at 1.634 pm (a5Fi~b3F2) would be ~ 1/3 as strong, but it was not seen because it lies in the blue wing of the much stronger [Si I]/[Fe II] feature. Using the transition probability computed by Nussbaumer k Storey (1988), and adopt ing the temperature of 3200 K derived in the previous section, we find a mass of Co II on day 349 of 6 x 10-3 M©, for optically thin emission. This result is sensitive to the adopted temperature — increasing Te to 4000 K halves the derived mass. The ionisation potentials of neutral cobalt and iron are similar, so cobalt would have been predominantly singly-ionised (Colgan k Hollenbach 1988). Thus the Co II mass we derive probably rep resents most of the cobalt present. This mass is at least two orders of magnitude greater than the mass of stable cobalt expected in the star prior to explosion (Rank et al. 1988). We conclude that the observed cobalt resulted predominantly from 56Ni decay. From the bolometric light curve we compute the mass of 56Co remaining on day 349 to be 3.55 X 10-3 M©, consistent with our spectrally derived mass estimates. Thus, on day 349 the observed intensities of iron and cobalt lines were consistent with optically thin emission at 3000 to 4000 K from the decay products of 0.08 M© of 56Ni. Nickel The strong [Ni I] transition at 3.119 pm (a3D3~a1D2) was present in the late-time spectra, blended with an unidentified feature at ~ 3.15 pm. The nickel feature became prominent Notes on derivation of Masses of Fe, Co and Ni. For a Boltzmann distribution of populations the intensity of anoptically thin emission line is given by: 1 Ngue Eu/kTAuihvui erg/s/cm 2 A'kD2 Q where N is the number of emitting atoms, gu is the degeneracy of the upper energy level, Eu is the energy of the upper level, T is the temperature, Q is the partition function, A is the radiative transition rate, v is the frequency of the transition and D is the distance to the supernova. For Fe II the atomic parameters are from Nussbaumer & Storey (1988). As the transitions are forbidden radiative transfer effects can be ignored. Multi-line Non-LTE model spectra of the forbidden iron and silicon lines indicate that the line ratios found in the spectra are consistent with Boltzmann distributions and therefore that the electron density is above the critical density 106 cm-3). The critical density at different temperatures was determined through the use of Non-LTE codes. Additionally the ionisation structure derived from line ratios such as Fe I/Fe II and the absence of Fe III lines in the K window together with the expected number density indicate that the electron density at the epochs considered is of order 108 cm-3 i.e. a factor of 100 above the critical density thus justifying the choice of Boltzmann distributions. The intensity of the line observed can be calculated from the spectrum by first removing the continuum under the line (see procedure described on page 47) and then integrating under the line. This procedure is atfrurate to ~10% excluding errors due to the intrinsic quality of the data. The initial mass of 56Ni is assumed to be 0.08-0.07 M® as inferred from the bolometric light curve of the supernova assuming that at late times the 56Ni decay powers the entire electromagnetic display from the supernova and there is no other contribution ( e.g. a pulsar) (Whitelock et al. 1988). The 1.547 /xm Co line is blended with the 1.53 /xm Fe II line. To extract a reliable flux for this line and therefore a mass we have used a line profile from a ‘clean’ line {e.g. 1.257 /xm) and blended the two lines present in the feature until a visual fit was obtained. The Co line was thus estimated at 70%±20% of the feature on day 349 (for detailed analysis of this feature see Varani et al. in preparation). The masses of Co and Fe are compared with the masses of 56Co and 56Fe that would be present in the ejecta allowing for the radioactive decay of 56Ni to 56Co and of 56Co to stable 56Fe. 60 & Chapter 3. IR Spectroscopy of SNe 1986G and 1987A 61 by day 349. Its presence was first reported by Witteborn et al. (1989) in spectra which they obtained on days 412-417. Subsequent spectra we have taken (day 494) show that the line retained essentially constant intensity as the local continuum faded (Allen et al. 1988). Using transition probabilities from Garstang (1964) we deduce the mass of neutral nickel on day 349 to have been about 6.5 x 10*"4 M© if the line was optically thin and assuming LTE at a temperature of 3200 K. Ni I has a similar ionisation potential to those of Fe I and Co I, and so we might expect a much larger mass to have been in the form of Ni n. Transitions due to [Ni II] occur in the J, H and K windows, but the potentially strongest of these all coincide with other prominent features. We therefore cannot, with any confidence, identify Ni II. At longer wavelengths, Aitken et al. (1988) have used their observations of the [Ni II] 10.68 pm line together with those by Rank et al. (1988) of the 6.6 pm line to obtain a total Ni II mass of 2.6 x 10-3 M© at about day 260. Given the differences in epoch, together with the uncertainties in the line intensities and atomic data, the inferred mass of Ni I on day 349 is consistent with most of the nickel having been singly-ionised. The 3.119 pm line might therefore have been expected to increase in prominence as Ni II recombined; this was indeed the behaviour reported by Allen et al. (1988). Clearly the nickel observed at these late epochs could not have been 56Ni (half-life 6.1 days), so must have comprised more stable isotopes. It is tempting to attribute an otherwise unidentified feature at 3.4 pm to the [Ni III] 3.394 pm (a1D2~a3Pi) transition. However, we would expect a comparable line at 3.8 pm from a 1D2~a3P2, and none is seen at that wavelength. Unidentified features Using the line identifications given above, no major features are unaccounted for in the day 18 spectrum. However in the later spectra there are several emission features which we have been unable to identify. These are at 1.16, 1.99, 2.07, 3.15, 3.40 and 3.53 pm. The 1.16 pm feature was blended with K I emission. The 1.99 pm line is probably due in part to Ca I. The 2.07 pm line is blended with the He I emission and the 3.15 ^m line is blended with the [Ni I] emission. The lines at 3.40 and 3.53 pm match well in wavelength the dust emission features in HD 97048 (Blades & Whittet 1980). However, they have 62 Chapter 3. IR Spectroscopy of SNe 1986G and 1987A retained almost constant equivalent width since day 110, and this would be unexpected behaviour for either dust condensing in the ejecta or for grains in the progenitor’s wind illuminated by the ultraviolet flash. The possibility that the 3.4 pm line is due to [Ni III] was dismissed above. While wavelength matches with known transitions can be made for most of these lines, all can probably be discounted because of the absence of much stronger emission expected at optical and other IR. wavelengths, although in some cases blends make it difficult to reach a firm conclusion. The lines therefore remain unidentified. Velocity behaviour We have measured the emission line widths at late times where the line was relatively strong and isolated viz. H (P /?, P 7, Br a, Br 7, Br 10), O I (1.13 pm), Mg I (1.503 pm) and [Fe II] (1.257 pm). The values, expressed as full-width half-maximum velocities (V fwhm) are given in Table 3.8. Between days 192 and 349 there was little significant change in the line widths. The hydrogen and helium lines exhibit similar widths, and yet, in unmixed models (e.g. Woosley’s (1988a) model 10H) the helium velocity is considerably less than that of the hydrogen. However, Woosley (1988a) points out that in a real supernova mixing of the H/He layers will occur both before and possibly also during the explosion itself. The insignificant difference between the observed hydrogen and helium velocities suggests that mixing did indeed occur (c/. Woosley’s (1988a,b) mixed model 10HM). None of the H or He line profiles exhibit evidence for the flat-top structure characteristic of stratification. The hydrogen and helium lines show faster moving material (Vfwhm ~ 3500 km/s) than do those of the heavier elements (Fe, Mg: Vfwhm ~ 2700 km/s; O: Vfwhm ~ 2100 km/s), implying that some degree of stratification was maintained, although the heavy element velocities reached more than twice those predicted by model 10H. Again, no evidence for the flat-top profiles is apparent. Of particular interest is that the [Fe II] line exhibits velocities, with respect to rest, of up to ~ 2500 km/s, which compares with a maximum velocity of ~ 1100 km/s for 56Ni in model 10H. Thus, the bubble produced by the 56Ni and 56Co decay may indeed have ‘popped’ (Woosley 1988a) with fingers of 56Co and other heavy elements mixing out into the H/He layers (c/. Woosley 1988b). Evidence of mixing has also been inferred from the line widths of the 17.93 pm and 25.99 pm Chapter 3. IR Spectroscopy of SNe 1986G and 1987A Table 3.8: FWHM of lines in the spectra of SN 1987A Feature Vfwhm (km/s) day 192 day 255 day 284 day 349 Mean P (3 3120 3120 2900 4030 3290 ± 250 P7 4400 3300 3500 3600 3700 ± 250 Bra 4140 3920 t 3950 4000 ± 70 Br7 3100 3400 2760 3260 3130 ± 140 Pf7 3450 3120 3800 3740 3530 ±160 BrlO 3240 3360 3180 3460 3310 ± 60 He I (1.083) 3400 4500 3200 3300 3600 ± 300 0 I (1.129) 1970 2020 2080 2190 2065 ± 45 Mg I (1.503) 2300 t 2840 2670 2600 ± 160 [Fe n] (1.257) 2640 2730 2720 2950 2760 ± 70 | Partial spectral coverage of feature. 64 Chapter 3. IR Spectroscopy of SNe 1986G and 1987A Table 3.9: Redshifts of emission line maxima in the spectra of SN 1987A Feature Redshift (km/s) day 192 day 284 day 349 Mean P P 660 740 780 730+40 Brj 850 880 1030 920+60 B ill 500 760 730 660+80 O 1(1.129) 1080 470 590 710+190 Mg 1 (1.503) 450 570 550 520+40 [Fe II] (1.257) 150 470 980 530+240 Mean 615+135 650+70 775+80 680+60 [Fe II] transitions (Erickson et dl. 1988; Moseley et al. 1988; Haas et al. 1988), of the [Ni II] 6.6 pm transition (Rank et al. 1988), and from the very early appearance of the [Co II] 10.52 pm line (Aitken et al. 1988). Mixing of 56Co into the outer layers has also been invoked to obtain a reasonable model fit to the bolometric light curve (Nomoto & Shigeyama 1988; Woosley 1988a,b), and to account for the early appearance of hard X- rays (Itoh et al. 1987) and 7-ray lines (Pinto & Woosley 1988) as well as the observed 7-ray line ratios (Matz et al. 1988). Table 3.9 gives the measured radial velocities of the more prominent lines in the J, H and K windows at several epochs. We repeatedly observed redshifts in the emission peaks, with values ranging from 150 to 1080 km/s. The wavelength uncertainty in these windows is only ~ ±100 km/s so the shifts are certainly real. The effect was observed in both allowed (H I, O I and Mg I) and forbidden ([Fe II]) transitions. Similar redshifts were seen in optical lines. There was a slight increase in redshift with time, and there is marginal evidence for a systematically larger redshift in the hydrogen lines. The mean of all listed velocities is +680 ± 60 km/s, but there were significant differences in redshifts both between species and between epochs. Balmer* series and [0 III] emission from circumstellar material (Wampler 1988) exhibited a redshift of +280 km/s. Presumably this is a measure of the intrinsic supernova centre of mass sp& A o^ k s j . o l l Chapter 3. IR Spectroscopy of SNe 1986G and 1987A 65 velocity. There is therefore a residual shift which ranges from -130 to +800 km/s, depending on the line and epoch considered. T f c of'\Q i\ cA p ^ tkj MCUj kLUUixl kj ^eJAc.fr6^ 3.3 Conclusions We have presented results from the infrared spectroscopy of SN 1986G and SN 1987A for early and late-time epochs. It is obvious why SN 1987A has revolutionised the field of infrared spectroscopy of supernovae. The detection of strong lines of [Fe II] in the late-time spectra of SN 1987A encourages us to search for these lines in the spectra of other supernovae. However, if a detection of infrared emission lines from an extragalactic supernova at late times is to be made we have to await the arrival of more sensitive instruments, such as two dimensional array spectrometers. 66 (10 7 Erg/s/cm2/Mm) Fx (10 7 Erg/s/cm2//^ni) Figure 3.11: March 1987 March 3.11: Figure Figure 3.10: March 1987 March 3.10: Figure Chapter 3. IR Spectroscopy of SNe 1986G and 1987A and 1986G SNe of Spectroscopy IR 3. Chapter H J window spectrum of SN 1987A. SNof spectrum window window spectrum of SN 1987A.SN of spectrum window Chapter 3. IR Spectroscopy of SNe 1986G and 1987A and 1986G SNe of Spectroscopy IR 3. Chapter (10 7 Erg/s/cm2/jum) Fx (10 7 Erg/s/cm2/Mm) Figure 3.12: March 1987 March 3.12: Figure Figure 3.13: March 1987 March 3.13: Figure lngt ^ ) (^m th g elen v a w K L wno pcrmo N 1987A. SN of spectrum window window spectrum of SN 1987A. SN of spectrum window 67 68 (10 7 Erg/s/cm2/yum) Fx (10 7 Erg/s/cmVjum) Figure 3.15: June 1987 June 3.15: Figure Figure 3.14: June 1987 June 3.14: Figure Chapter 3. IR Spectroscopy of SNe 1986G and 1987A and 1986G SNe of Spectroscopy IR 3. Chapter H J window spectrum of SN 1987A. SN of spectrum window window spectrum of SN 1987A.SN of spectrum window Chapter 3. IR Spectroscopy IR 3. Chapter (10 7 Erg/s/cm2//im) F* (10 7 Erg/s/cm2/^m) Figure 3.16: June 1987 June 3.16: Figure Figure 3.17: June 1987 June 3.17: Figure o f f o wavelength wavelength lngt ^ ) (^m th g elen v a w SNe 1986G and 1987A and 1986G SNe K L window spectrum of SN 1987A. SN of spectrum window window spectrum of SN 1987A. SNof spectrum window {/ jl m) 69 70 (10 7 Erg/s/cm2//2m) Fx (10 7 Erg/s/cm2/Mm) 11 . 1.3 1.2 1.1 1 Figure 3.19: Septemberl987 Septemberl987 3.19: Figure Figure 3.18: Septemberl987 Septemberl987 3.18: Figure aeegh (^m) wavelength Chapter 3. IR Spectroscopy of SNe 1986G and 1987A and 1986G SNe of Spectroscopy IR 3. Chapter H J wno pcrmo N 1987A. SN of spectrum window window spectrum of SN 1987A.SN of spectrum window pe . R Setocp Ne 96 ad 1987A and 1986G e SN f o Spectroscopy IR 3. apter h C (10 7 Erg/s/cm Z/^ m ) Fx (10 7 Erg/s/cm2//mi) Figure 3.20: Septemberl987 Septemberl987 3.20: Figure Figure 3.21: Septemberl987 Septemberl987 3.21: Figure vee h th g elen av w K L window spectrum of SN 1987A. SNof spectrum window window spectrum of SN 1987A. SN of spectrum window ( jll ) m 71 72 (10 7 Erg/s/cmV^ni) Fx (10 7 Erg/s/cm2/Mm) Figure 3.22: Septemberl987 Septemberl987 3.22: Figure Figure 3.23: Novemberl987 Novemberl987 3.23: Figure atr3 I pcrsoyofS 18G ad 1987A and 1986G e SN f o Spectroscopy IR 3. hapter C M J wno pcrmo N 1987A. SNof spectrum window window spectrum of SN 1987A. SNof spectrum window pe . R Setocp Ne 96 ad 1987A and 1986G e SN f o Spectroscopy IR 3. apter h C (10 7 Erg/s/cmVa™ ) Fx (10 ? Erg/s/cmV^^) Figure 3.24: Novemberl987 Novemberl987 3.24: Figure Figure 3.25: Novemberl987 Novemberl987 3.25: Figure H K window spectrum of SN 1987A. SN of spectrum window window spectrum of SN 1987A. SN of spectrum window 73 (10 7 Erg/s/cm2/yum) Fx (10 7 Erg/s/cm2//im) 74 Figure 3.26: Novemberl987 Novemberl987 3.26: Figure Figure 3.27: Decemberl987 Decemberl987 3.27: Figure 35 4 3.5 3 aeegh (yLzm) wavelength atr3 I pcrsoyofS 18G ad 1987A and. 1986G e SN f o Spectroscopy IR 3. hapter C J L window spectrum of SN 1987A.SN of spectrum window window spectrum of SN 1987A. SN of spectrum window pe . R Setocp Ne 96 ad 1987A and 1986G e SN f o Spectroscopy IR 3. apter h C (10 7 Erg/s/cmV^m) Fx (10 7 Erg/s/cmV^m) o I i11 i i i I i 1 i i i i Figure 3.28: Decemberl987 Decemberl987 3.28: Figure Figure 3.29: December1987 December1987 3.29: Figure . 16 . 1.8 1.7 1.6 1.5 aeegh (yum) wavelength lngt /m) (/im th g elen v a w H K window spectrum of SN 1987A. SN of spectrum window window spectrum of SN 1987A. SN of spectrum window _ i _ i _ i _ i _ L 75 76 (10 7 Erg/s/cm2//^111) Fx (10 7 Erg/s/cm2/Mm) Figure 3.30: Decemberl987 Decemberl987 3.30: Figure Figure 3.31: Februaryl989J Februaryl989J 3.31: Figure vee h (yum) th g elen av w aeegh (yum) wavelength atr3 I pcrsoyofS 18G n 1987A and 1986G e SN f o Spectroscopy IR 3. hapter C J L window spectrum of SN 1987A.SN of spectrum window window spectrum of SN 1987A. SN of spectrum window pe . R Setocp Ne 96 ad 1987A and 1986G e SN f o Spectroscopy IR 3. apter h C (10 7 Erg/s/cm2//im) Fx (10 7 Erg/s/cm2/yum) iue33: eray181 idwsetu fS 1987A. SN of spectrum window 198^17 February 3.32: Figure Figure 3.33: Februaryl98!Bi7 window spectrum of SN 1987A. SNof spectrum window Februaryl98!Bi7 3.33: Figure lngt th g elen v a w m) tm p ( 77 78 Chapter 3. IR Spectroscopy of SNe 1986G and 1987A Figure 3.34: February1980 L window spectrum of SN 1987A. wavelength Gum) Chapter 4 Interpretation of the Infrared spectra of supernovae In this chapter we shall develop models to interpret the infrared spectra of supernovae. Two models will be presented. Through the early-time spectral model we shall determine the population structure of the hydrogen in the ejecta of SN 1987A. This is done by fitting theoretical P-Cygni line profiles to our data. Through the late-time spectral model we shall determine a self-consistent ionisation structure, electron density and temperature for the ejecta of a Type la supernova and produce some synthetic spectra to interpret the spectra of SN 1986G. 4.1 Early time spectra In the infrared spectra of SN 1987A and SN 1986G a strong continuum is present at early times. On this continuum P-Cygni lines axe superimposed. In SN 1986G, a Type la supernova, these lines arise from metals whereas in SN 1987A, a Type II, the lines arise predominantly from hydrogen. It is evident that a complete model of the early time spectra should produce not only the continuum observed but also the lines. For a more complete understanding of the nature of the early supernova the conditions under which the continuum and the lines form should be addressed. Here we will discuss the theoretical models of the continuum formation separately from our models of the line formation at 79 80 Chapter 4. Interpretation of the IR spectra of SNe these eaxly epochs. 4.1.1 Theory of continuum formation Studies of continuum formation in supernovae have concentrated on Type II events ( e.g. Shaviv et al. 1985; Wagoner 1981). This is due to the fact that most work in this field has been instigated with the goal of utilising supernovae as distance indicators. Type I events are used as ‘standard candles’ to obtain distance estimates t.e. their spectroscopic qualities are not used. Consequently their continua have not been extensively modelled. Type II supernovae can be used as distance indicators through the use of the Baade-Wesselink method. In this case the conditions at the photosphere are of considerable importance. How does the photosphere of a supernova evolve? During the first 2-3 days the pho tosphere is locked into the steep outer boundary of the expanding fireball (Hoflich 1987, Dopita 1988). The effective temperature decreases rapidly through adiabatic cooling. Around 7-14 days after the explosion owing to the recombination of the hydrogen and the geometric dilution of the material the photosphere slows down in spatial co-ordinates and recedes in mass co-ordinates. In a low density fully ionised plasma the dominant source of opacity is electron scattering taking place in a high velocity gradient environment. It is important to note that the region where photons axe last emitted from is where Taba = 1- This is the deepest layer in the supernova which we could possibly see. It is easiest to understand the effects of scattering on the photons from the continuum formation region in terms of the parameter £„. kaba (v = u«cat"f" kaba This parameter gives the probability that photon is thermalized per scattering event. We can therefore define a thermalization depth based on a random walk for the photon through the scattering regions rther (Mihalas 1978). Thus a photon that is emitted at a large optical depth is more likely to be thermalized than one emitted further out. Given that the absorptive opacity kabB varies as ^“^the thermalization depth will vary as v. The surface of last scattering is therefore formed at lower total optical depths as we move to longer wavelengths. For example Hoflich et al. (1986) calculated that whereas the continuum at 5000 A forms at a total optical depth of ~32, at 10000 A the corresponding value for the total optical depth is ~18. One of the more important results Chapter 4. Interpretation of the IR spectra of SNe 81 of this effect is that the temperature of the emitting region will be higher than the effective temperature which would correctly determine the bolometric luminosity of the supernova. This has profound effects on the use of supernovae as distance indicators. If we use the colour temperature of the spectrum to determine the luminosity of the supernova as a black-body we would be overestimating the effective temperature of the flux emitted and therefore we would overestimate the total luminosity. The determination of the distance to Type II supernovae therefore has to take into account the effects of scattering so as not to underestimate the Hubble constant. Additionally the temperature of the emitting region is wavelength dependent. However this variation in the temperature is relatively small (~15% between 3000 and 5000 A ) (Hoflich et dl. 1986). The effects of scattering on the line profiles may explain why the absorption trough of the infrared lines lies at a velocity that is lower than that inferred from the photometric observations for the photosphere in the spectra of SN 1987A taken on day 18 after its explosion (Hoflich private communication). In models of continuum formation ( e.g. Shaviv et dl. 1985) the equation of radiation transfer is solved in a spherically symmetric atmosphere taking into account the free- free, bound-free, bound-bound and Thomson opacities. A steep density gradient and homologous expansion are assumed. These models assumed LTE populations but this assumption has been dropped by Hoflich (1987) in his calculations. In addition in Hoflich’s models the line formation is also calculated. He finds that although the continuum is formed under LTE conditions large deviations from the LTE occupation numbers are found above the photosphere. Hauschildt et dl. (1989) have calculated model continua for the near infrared spectral region. They have modelled the spectrum of SN 1980K (a Type II) very successfully including the somewhat depressed K window continuum. Such a depression is also present in the spectra of SN 1987A from day 18 after the explosion although this may be the result of errors in the absolute fluxing of the spectrum. Clearly such models need to be applied to our spectra to determine the conditions at the continuum formation regions. 4.1.2 Theory of P-Cygni profile formation Although as discussed above the conditions near the photosphere of a supernova are not simple, we have attempted some modelling of the line formation in the early supernova, 82 Chapter 4. Interpretation of the IR spectra of SNe assuming a single well defined photosphere. Similar models have been applied to the optical spectra of supernovae (Kirshner & Kwan 1974, Branch et al. 1981, Hempe 1981). These models have been based on the theoretical calculations of Castor (1970) and/or Lucy (1971) for an expanding envelope above a sharp photosphere. Although more comprehensive spectral codes exist ( e.g. Hoflich 1987) no results from these models for the early time infrared spectra of SN 1987A have been published to date. The derivations of the formulae used here can be found in Castor (1970). We have assumed a homologous expansion and the escape probability approximation (Sobolev 1960) is used. In this, owing to the presence of the steep velocity gradient, when a photon is emitted it only travels a short distance before going out of resonance with its parent transition. The probability, /?, of escape from the region of emission is: 1 — e-1 0 = (4.1) where r is the optical depth for scattering (Castor 1970). In such an atmosphere the con tinuum photons from the photosphere of appropriate wavelength are resonantly scattered by atomic transitions in the expanding gas above the photosphere. Since the emission from these atoms is isotropic and since the photospheric disc obscures some of the atmosphere on the far side, the resulting lines have ‘P-Cygni’ type profiles i.e. emission peak at the local rest wavelength and a blueshifted absorption trough. 5 1 de, pfofVfib. In the coordinate system used (see Fig. 4.1) the line of sight coordinate is z and the coordinate p is defined such that p2 -f z2 = r2. The shape of a P-Cygni line profile can then be shown to be given by the equation : Fv - F c 1 f°° S((p2 + z2y /2) = i r (1 - exp{—r(r) }) 2 p dp Tl h 1 Fe —- (1 - exp{-r(r)y})2pdp rc Jo _1_ rrc S((p2 + *g)i/2) (exp{-r(r)y} - exp{r(r)})2p dp r2 Jo Ic (4.2) where Ic is the flux from the photosphere, 5 is the source function and rc is the radius of the photosphere (Castor 1970; Mihalas 1978.) The y term is a step function which takes into account the effects of the optically thick photosphere and therefore the occultation of the far side of the atmosphere and the ‘absorption’ from the near side, zo in a homologous expansion define^a locus of values of p at which the atoms are observed to emit Chapter 4. Interpretation of the IR spectra of SNe 83 Figure 4.1: Supernova co-ordinate system © at the same frequency. The source function for a two level atom in a homologous expansion can be shown to be (Castor 1970; Mihalas 1978) : (1 -€)PW IC + €B„ (4.3) (1 - e)(3 + e where is the geometrical dilution factor, and € = e' €'+ 1 with c* = Cut (1 - exp(—hu/kT))/Aui T is the temperature of the electron gas. e can be understood as a photon destruction coefficient being the ratio of the collisional de-excitations to the total de-excitations from 84 Chapter 4. Interpretation of the IR spectra of SNe Figure 4.2: Optical depth effects on the line profile. Profiles labelled according to tq the upper level. The optical depth for a photon in resonance with a transition from a lower level / to an upper level u in a homologously expanding atmosphere is: 2 El &L flu r(r) =^r(5/)lu-a-me ' ' u0V(r) cr (4.4) where vq is the rest frequency of the transition, V(r) is the velocity of the emitting region and the other symbols have their usual meanings. We have developed a code to evaluate equation 4.2 in a homologously expanding at mosphere. Results from the code with = 2rc indicating the effects of optical depth and different density power laws are shown in Figs 4.2 and 4.3. Where two lines overlap we have used the formula of Castor & Lamers (1971) for the calculation of a doublet line profile. The profile of a doublet can be approximated by FD(v) = F r ( v ) + [FB(v) — l] expf-r^dul)] + Q R(v) (4.5) with Q = / [FB(v) - 1]{1 - exp[-rR(\v\)]} dv J-Voo Chapter 4. Interpretation of the IR spectra of SNe 85 Figure 4.3: Density power law effects on the line profile. Profiles labelled according to n. and fVoo *00 = [^(M ) - l]/2 / [Ffi(|t>|) - 1] dv where F is the emission relative to the continuum and v = — 1) . The superscripts R and B refer to the relative shifts (red or blue) of each line being blended. The first term in equation 4.5 is the red line profile which is unaffected by the presence of the blue component. The second term is the contribution of the blue component which is reduced by a factor exp(—rR) due to scattering by the red component. The third term is the radiation from the blue component which is scattered and redistributed over frequency by the red component. The factor Q is the amount of scattered radiation and the function R(v) describes the redistribution of this radiation by the red component. The redistribution is assumed to be the same as the original distribution of the flux by the red component namely i 7'fl(|u|) — 1. The effects of blending two lines can be seen in Fig. 4.4. It is important to note that in the short wavelength region where the two lines do not overlap v < Uqo , FR(v) = 1, rR(|v|) = 0 and R(v) = 0; so FD = FB which is the undistributed part of the blue profile. If the separation of two lines is greater than 2Voo then FR(v) = 1 86 Chapter 4. Interpretation of the IR spectra of SNe Figure 4.4: Blending of two P Cygni lines. in the wavelength region of the blue profile and FB(v) = 1 and Q = 0 in the range of the red profile. Thus in this case the blending formula correctly returns two undisturbed line profiles. When calculating a blend of more than two lines we first calculate each line indepen dently. We then start at the blue end of our spectrum and blend with the first line to the red. This then forms the new blue line and it in turn is blended with the next red line. Thus a complete model spectrum may built up. 4.1.3 Modelling the early time spectra of SN 1987A We have used the theory described above to model the infrared spectra of SN 1987A taken on day 18 after the explosion. The observed spectra have been described in chapter 3. They show a strong near-thermal continuum with P-Cygni lines superimposed. Since nearly all the early-time features originated from transitions of hydrogen we have considered only this element. As discussed in chapter 3 the flux in the emission component of these lines Chapter 4. Interpretation of the IR spectra of SNe 87 exceeds that ‘taken out’ in the absorption trough. This is probably due to the large volume over which the emission due to collisional excitation is also active. In the source function we used for our models we assumed the photosphere emitted a black-body spectrum at the temperature derived in chapter 3 and that the thermal electron gas above the photosphere could be described by the same temperature. The factor € can be calculated from atomic data given the temperature and electron density and is generally Reasonable fits to the profiles are achieved, although the quality of fit varies from line to line. In particular Pa 7 is clearly contaminated by other lines to the blue. The departure coefficient for levels 7 and 8 are not consistent with the values derived for the other energy levels. However, neither Br 7 nor Pf 7 which originate from these levels are particularly well fitted in our spectra. Clearly a single thermal population structure for 88 Chapter 4. Interpretation of the IR spectra of SNe Figure 4.5: Model spectrum for J window compared with data. E N OE W CIO w r- lO .< the entire nebula is ruled out. The aim of this modelling was to establish the populating mechanism of the hydrogen. We interpret the leveling off of the departure coefficients as we move to higher principal quantum numbers as evidence for collisional excitation being the dominant populating mechanism. Collisional excitation from an electron gas at ~5000 K will be more efficient where energy levels are closely packed such as the the higher energy levels of hydrogen and will tend to establish relative LTE populations. C hapter 4. In terpretation o f th e IR spectra o f SNe f o spectra IR e th f o terpretation In 4. hapter C (10 7 Erg/s/cm2/Mm) Fx (10 7 Erg/s/cm2//im) Figure 4.6: Model spectrum for for spectrum Model 4.6: Figure Figure 4.7: Model spectrum for for spectrum Model 4.7: Figure aeegh (^m) wavelength H K window compared with data. with compared window window compared with data. with compared window 89 90 departure coefficient FA (10 “7 Erg/s/cm/^m) 2 Figure 4.8: Model spectrum for for spectrum Model 4.8: Figure Figure 4.9: Hydrogen departure coefficients. departure Hydrogen 4.9: Figure atr4 I epeain ofteI setaofS e SN f o spectra IR the f o terpretation In 4. hapter C L window compared witli data. witli compared window Chapter 4. Interpretation of the IR spectra of SNe 91 4.2 Late time spectra In this section we shall describe a model for the late time spectra of supernovae. We shall compare the results of the model with the spectra of SN 1986G described in chapter 3. In this model we have assumed that the radioactive decay of56Co powers the electromagnetic display at the epochs concerned. The model traces the deposition of the 7-rays produced in the radioactive decay. These 7-rays Compton scatter accelerating ‘primary’ electrons to ~1 MeV energies. The primary electrons then ionise and excite the ejecta and also deposit some of their energy into the quasi-thermal electron gas. The ionisation structure of the ejecta is then determined by the interactions of the primary and thermal electrons with the ejecta. The structure of the model is based on similar calculations by Meyerott (1980), Axelrod (1980) and Fransson & Chevalier (1989). Here the thermal and ionisation balance is solved for a single species plasma only. A fully comprehensive model should include all relevant species. Most of the formulation of the code is such that it may be expanded for use in a multi-species plasma. 4.2.1 The 7 rays The radioactive decay of 56Co generates 7-rays with discrete energies of typically 1 MeV, plus a continuous energy spectrum of positrons with a maximum energy of 1.46 MeV. 4% of the energy of the decay is carried by the positrons (Colgate & McKee 1969). The positrons may (Axelrod 1980) or may not (Colgate et al 1980) deposit their energies in the ejecta. Axelrod (1980) has pointed out that even for weak magnetic field the cyclotron radius of a 1 MeV positron is a factor of 10 7 smaller than the typical radius of a supernova. Therefore unless the magnetic field is radially combed as suggested by Colgate et al. (1980) the positrons will deposit their energy in the ejecta. If the positrons remain in the ejecta they will eventually annihilate but only after having deposited their kinetic energy (Arnett 1979). We follow the energy deposition of individual 7-rays through a Monte-Carlo technique developed by Cashwell & Everett (1959). In our code the 7-rays Compton scatter off electrons accelerating them to MeV energies. The code calculates not only the deposition fraction of the 7-rays but also the distribution of the energies of the primary electrons. The radial dependence of the deposition can also be calculated. The effective opacity for 92 Chapter 4. Interpretation of the IR spectra of SNe Figure 4.10: Deposition fractions for 7-rays vs column density in the ejecta the 7-rays k7 is taken to be 0.033 cm2/g (Axelrod 1980). Although the code is written so that the ejecta can be split into any number of zones with different densities for simplicity we have calculated the deposition of the 7-rays from a central source in a uniform density envelope. Fransson & Chevalier (1989) have shown this to be a reasonable approximation. The deposition fractions are plotted against column density in Fig 4.10. 4.2.2 Non-thermal electron energy deposition The primary electrons deposit their energy in the ejecta through interactions with the electron gas or through ionisations and excitations of atoms. The deposition of their energy can be dealt with as a local problem as their mean free path is of order 1% of that of the 7-rays (Fransson & Chevalier 1989). Additionally as mentioned earlier in the presence of magnetic fields this mean free path will become even shorter. Therefore the deposition of primary electron energy into the ejecta can be taken to equal that of the 7-rays into electron energy. We axe interested in the percentage of the primary electron energy which goes respectively into ionisations, excitations and the thermal electron gas. Chapter 4. Interpretation of the IR spectra of SNe 93 In any plasma powered by energetic electrons secondary electrons axe produced when the primaries ionise the ejecta. These secondary electrons do not necessarily have low energies and can themselves ionise and excite the plasma producing in turn further secondary electrons. This cascade has been discussed in the context of hydrogen dominated plasmas (e.g. Shull 1979). In our model we follow Shull’s (1979) method. In order to take the secondaries into account we start the calculation of the relative energy deposition at low energies. We then consider the deposition for a slightly higher energy. For each energy any secondary electrons produced, which by definition have lower energies that the primaries, will have their relative depositions already calculated. To solve this problem we need expressions for the cross-sections for electron-electron interactions, and the ionisation and excitation cross-sections for the energy range Emin to Emax where Emin is the starting energy which will define the thermalization energy for the electron gas and Emax ~ IMeV. The cross-section for electron-electron interactions resulting in an energy loss A E can be shown to be (Habing & Goldsmith 1971): cTee(AE) = 40 7re4 (Ink) (—.05, ) 1 E 2 cm2 (4.6) where / = (AE/E) and values for InAjcan:an(+ kbe found in Spitzer (1962).Spitzer (1962). The ionisation cross-section used are from Lotz (1967) and have the form: . _ Aaln(E/x) '° C XE cm (4.7) Values for the the parameters £,a and Xcan¥ also be found in Lotz (1967). The distribution of the secondary electron energies, P(Eo,E), has been measured by Opal et al. (1971) and can be approximately described by: P^=TT(W (4'8) where A is weakly dependent on Eo, the primary energy, and J is a constant characteristic of each system. Because of the identity of the outgoing electrons in these measurements the convention is generally made that E refers to the less energetic electron and ranges from 0 to (Eo — 7)/2 where I is the ionisation potential of the atom. In the Monte-Caxlo code when an ionisation occurs the resulting secondary electron energy is chosen at random weighted by the distribution shown in equation 4.8. The results of our code are parameterized in terms of Xc which is the ratio of the electron number density to the total number density and is determined by the ionisation equilibrium in the plasma. The excitation and ionisation cross-sections depend on the £\S*\ ’’0*$ c ( th i'i3MIS£U,0”'>e ° V&wh'oi (K f 5 h u J r ru *f*IO 94 Chapter 4. Interpretation of the IR spectra of SNe Figure 4.11: Division of primary electron energy. 1 0001 &°{Ne/N&) 1 1 atomic species present in the ejecta. It is possible to calculate the fractional deposition of energy in a multi species plasma through the use of this code. For simplicity we have modelled the relative deposition fractions in an oxygen plasma. Oxygen is present in all types of supernova and has reasonably well defined cross-sections. Although a more detailed calculation should include all species present in the ejecta it should be pointed out that the results of this Monte-Carlo calculation, in the high ionisation limit, agree with the semi-analytical calculations of Axelrod (1980) for iron. The excitation cross-sections for O I are from Kazaks et al. (1972), Sawada & Ganas (1973), Thomas & Nesbet (1975) and Smith (1976). The fact that we ignore excitations of other ionisation stages does not affect our results adversely as we shall see that excitations to not form a major deposition path for the primary electrons for the conditions observed to be present in the supernova. Three ionisation stages of oxygen were included in the calculations. The results from the code can be seen in Fig. 4.11. Note that as Xe increases the fraction of energy going into excitations drops whereas that going into the thermal electron gas increases. This is a result of increased electron density. Also note that the electron-electron cross-section has a steeper energy gradient Chapter 4. Interpretation of the IR spectra of SNe 95 than the ionisation cross-section (see equations 4.6, 4.7) and therefore ionisations are the dominant interaction at high energies. The primary energy lost in an ionisation is dictated by the energy of the secondary electron produced. Therefore the deposition fractions are not sensitive to Emax provided it is higher than ~10 keV. This is in agreement with Shull (1979) and is important since it means that the precise spectrum of the primary electrons (and the 7-rays) is not important in our calculations. 4.2.3 The thermal and ionisation balance Above we calculated the fractions of the radioactive energy going into the thermal electron gas and into ionisations. Note that we cannot determine either the ionisation or the thermal balance independently as photons from the recombination of the ionised species will also be able to ionise the gas and the recombination coefficients are temperature dependent. We have therefore developed an iterative process along the lines of Axelrod (1980) which calculates a self-consistent temperature, electron density and ionisation for a plasma. The calculations are based on the assumption that such a balance is possible. Examination of the typical timescales for the processes present vindicate this assumption (Axelrod 1980). The thermal balance We have determined that the main recipient of the energy from the 7-rays at high ionisa- f iQwly£*Kil) tions will be the thermal electron gash In fact the results indicate that almost 90% of the energy goes to the thermal electrons with 10% being deposited in the form of ionisations. We have modified the source function describing the heating of the thermal electron gas by the positrons (Axelrod 1980) to include the effects of the 7-rays. s — fther X ( L(T, Ne) = ^2 n* erg/s/atom (4.10) « j where n, is the population of energy level i and the other terms have their usual meanings. We therefore need to establish the population of each energy level. This can be done by finding the solution to the simultaneous equations determined by the populating and depopulating mechanisms ;Results from such a Non-LTE calculation for a 25 level Fe III atom at a fixed electron density can be seen in Fig. 4.12. Note that as the temperature drops the cooling curve flattens before it starts decreasing again. This is due to the fact that the ground level energy levels can be collisionally populated over a broader range of temperatures than the the higher energy level transitions. This effect may result in the infrared catastrophe as first discussed by Axelrod (1980). As the nebula cools the heating of the thermal gas drops due to the exponential term in equation 4.9 whereas the cooling efficiency will remain constant. Therefore the temperature decline may accelerate resulting in only the fine-structure lines originating within the ground level transitions being observed. As these occur predominantly in the far-infrared the emphasis in the spectrum will also shift to that spectral region. To pin down the temperature of the electron gas we need to know not only the populations of the levels in each atom but the ionisation structure of the nebula as well. The ionisation balance Atoms in the ejecta are ionised by the energetic electrons and by UV recombination photons. When calculating the ionisation balance we need to take both these effects Chapter 4. Interpretation of the IR spectra of SNe 97 Figure 4.12: Cooling curve for FeUI into account. Tlie ionisation balance equation therefore can be written as (Axelrod 1980): - 7i/t “ ^2pt^aj+iNefj+i + (1 ~ Pt,t)<*t+i(T)Neft+1 = 0 (4.11) j > * where /, is the fractional population of ion i and 7, is the ionisation efficiency of the 7-rays for ion i : S where S is the source function from equation 4.9 modified for ionisations and W is the work per ion pair for ion i taken from Meyerott (1980). Nc is the electron density and a,(T) = 3.0 x 10-13(t)2(T/104)"3/4 cm3s_1 is the recombination coefficient to all states of ion i (Allen 1973). If we consider only the recombinations to ground then that coefficient becomes (Allen 1973): af(T) = 1.0 x 10"13(i)2( r /104)-1/2 cm3s_1 In equation 4.11 P,tJ is the probability that a photon from the recombination to ion j results in the ionisation of ion i. In calculating the values of PXyJ we have followed the 98 Chapter 4. Interpretation of the IR spectra of SNe formalism of Axelrod (1980). Only photons from recombinations to the ground state of atoms of greater or equal ionisation potentials than the one considered are considered able to ionise that ion. Given the photo-ionisation cross-sections and the abundances of the relevant ions the total optical depth for the recombination photons may be determined. PM is then simply the ratio of the optical depth for UV photons that will ionise ion i over the total optical depth for the recombination photons. Since not all recombinations are to ground a free parameter exists in this formalism which reduces the probability of a recombination photon being recycled. This recycling parameter is difficult to determine other than to require that it is > a9/a (Axelrod 1980). The procedure to determine a self-consistent electron temperature, density and ion isation structure is as follows. We estimate the temperature, total number density of the ejecta and an initial ionisation structure. The ionisation structure and total num ber density determine the electron density. This together with the assumed temperature then allows us to calculate the recombination coefficients. We then solve the simultaneous equations defined by equation 4.11 for each ion and the condition that N £ / . = 1 i=0 The results then determine a new ionisation structure. The cooling of the nebula due to collisional excitations is then calculated as described above. We then compare this cooling to the heating source function. Iterating around this cycle varying only the temperature until the cooling and heating rates are balanced then determines a self-consistent temper ature and ionisation structure for the plasma. At each stage depending on the ionisation structure the deposition fractions into ionisations and the thermal electron gas are varied according to the results obtained for the deposition of the primary electrons. The only input parameters therefore are the composition of the ejecta, the time since the explosion and the total number density. The results One of the ‘standard’ models of a Type la supernova explosions is the deflagration model W7 of Nomoto et al. (1984). Branch et al. (1985) have successfully fitted the early optical spectra of SN 1981B (a Type la) with a synthetic spectrum using a supernova composition and expansion velocities based on W7. In W7 the supernova produces ap- C h apter 4. Interpretation of the IR spectra of SNe 99 Table 4.1: Results of self-consistent model for ionisation and thermal balance equations Epoch Number recycling Fel Fell FelH FelV Nc T density fraction (days) cm -3 cm-3 K 350 3.8x10s .5 .01 .16 .47 .36 8.3x10s 4510 450 1.8x10s .5 .01 .21 .48 .30 3.7x10s 3870 550 9.7X104 .5 .02 .33 .47 .18 1.8x10s 2770 650 5.9X104 .5 .08 .58 .30 .04 7.7X104 850 700 4.7X104 .5 .14 .64 .21 .01 5.2X104 510 proximately 0.6 M© of 56Ni and approximately 0.1 M© of 54Fe with an maximum velocity of ~ 9000 km/s. In W7 the iron group elements dominate the inner 7000 km/s of the supernova. We followed the procedure described above to model the late time conditions in an iron plasma including 4 ionisation stages including a total of 135 energy levels. The relative deposition fractions into ionisations and the thermal electron gas derived from the oxygen plasma are similar to those derived by Axelrod (1980) for an iron plasma using a continuous slowing down approximation and therefore applicable to the iron plasma as well. We present results from the code for different epochs for 0.7 M© of iron expanding at 7000 km/s in Table 4.1. The collision strengths and A values used in the determination of the cooling were from Grevesse et dl. (1971) for Fe I, Nussbaumer & Storey (1980; 1988) for Fe II and Garstang et al. (1978) for Fe III. Where collision cross-sections were not available we used the empirical approximation described in Axelrod (1980). Fe IV does not have any ground state forbidden transitions and therefore does not contribute signif icantly to the cooling. The photo-ionisation cross-sections were from Reilman & Mason (1978). The results are parameterized in terms of the density of iron in the ejecta. The ionisation structure was such that the neglect of excitations was justified (see above and table 4.1). It is important to note that though the heating of the nebula is determined by the amount of 56Ni produced in the explosion, the cooling is effected by all the isotopes of iron present in the ejecta. The final results are sensitive to the recombination photon recycling fraction assumed. For example at an epoch of 350 days changing the recycling fraction from 0 to 0.7 results in the population of Fe II decreasing from 23% to 13% and the temperature increasing from 4300 K to 4700 K. 100 Chapter 4. Interpretation of the IR spectra of SNe 4.2.4 The final spectrum Given the ionisation and thermal structure of the ejecta we should now be able to calculate a model spectrum. In a homologous expansion the location of points giving rise to emission at a certain frequency v from a transition with rest frequency vq is a disc. For a forbidden line where line scattering may be neglected given the emissivity as a function of density it is easy to calculate to emission from each disc and therefore the line profile. In the uniform density case the line profile can be shown to be parabolic. In addition when comparing the theoretical line profiles with those observed the effects of electron scattering may need to be considered. Ignoring Compton effects which at IR frequencies are minimal the change in frequency of a scattered photon is simply: v= c —(cosd i — cosB 2) where 6\ and 02 are the angles the incoming and outgoing photons make with the bulk velocity of the material and v is the velocity. The cross section in a coherent scattering of unpolarized light off an unbound charge is (Haxwit 1973) : d(T 1 . 6 9 , —dfl = 2-(— v me2 7)2( J v i + cos2 e) ' where 0 is the angle between the incoming and outgoing photon. Thus the combination of the Doppler shift and the preferential forward (or backward) scattering result in a net redshift of the photons emitted by the supernova. We have modified the Monte Carlo code that calculates the Compton scattering of the 7-rays to calculate this effect on line profiles resulting from different density gradients. The obvious result of these calculations is the presence of a redward wing in the line profiles. Depending on the distribution of the emitting material the peak of the line profile may also exhibit a shift in wavelength. A sample line profile produced by this code can be found in Fig. 4.13. 4.2.5 Comparison with observations In the use of the code described in this section to interpret the spectra of SN 1986G we have assumed a uniform density for the iron. As mentioned above the line profile for this case is parabolic. We have used the results of the Non-LTE cooling calculations to determine the spectrum of the resulting supernova in the H window and convolved it with such a line profile. In Figure. 4.14 the model spectrum is compared with the spectrum Chapter 4. Interpretation of the IR spectra of SNe 101 Figure 4.13: Effects of scattering on line profile. of SN 1986G taken in April 1987 at the AAT. The model spectrum was smoothed to simulate the binning process described in chapter 3. The parameters used for the model spectrum were 0.034 M© of Fe II, a temperature of 4500 K, an expansion velocity of 7000 km/s and 0.45 magnitudes of extinction in the H atmospheric window (Phillips et dl. 1987). The distance to NGC 5128 was taken to be 5 Mpc (see chapter 3). Given the uncertainty in the distance to NGC 5128 and uncertainty in the ionisation structure (see above) the spectrum observed at the AAT in April 1987 is consistent with the spectrum predicted by our model for the deflagration supernova W7 of Nomoto et dl. 1984. Note that given the low significance of the feature in the observed spectra (see chapter 3) we are only attempting to place a limit on the amount of Fe II that could result in the observed feature if it were real. It should also be pointed out that given these same uncertainties the CTIO data on SN 1986G are also consistent with model W7. A more detailed treatment of the recombination radiation may reduce the uncertainty in the ionisation structure of the gas and therefore allow more stringent limits to be placed on the deflagration models. Oof tocxllll d o e * 'V cot p & L ( L i ©~i *&■ £ \c ow o m \V "hk_ W> HOCSlZJp X iV “ftc i om! s\lUcil)r£ . 102 Chapter 4. Interpretation of the IR spectra of SNe Figure 4.14: Model spectrum and AAT April 1987 SN 1986G spectrum^ lo/\*hVua£ p£L LL^ ooUcto-o /^ocUL p^- 15 E 7) 10 l oS \tuO U0) -< 0 1.5 1.6 1.7 1.8 wavelength (jum) 4.3 Conclusions The two models presented here are too simple to give a complete picture of the conditions in the ejecta of a supernova. However, through their use a better understanding of the dominant effects which determine those conditions may be attained. A source function that takes into account Non-LTE effects is clearly needed for the complete interpretation of the early time spectra. For the late time spectra we need to take into account the recycling of the recombination photons using radiation transport codes to determine the fate of the strongest recombination lines from each ion involved. Before this code can be applied to late time spectra of SN 1987A additional modifications will be needed. In particular charge transfer effects need to be included in the equation of the ionisation balance (eq. 4.11). In its present formulation the code is only applicable to a single species plasma. Chapter 5 Molecules in Super novae In this chapter we shall discuss the unique observations of molecular emission seen in the spectra of SN 1987A. Since all molecular emission observed in the spectra of that object has been from diatomic molecules we shall give a brief description of the nature of the spectra originating from such molecules. We shall also discuss the application of the theory in modelling of the observed emission and the results and implications of these models. 5.1 The infrared spectra of diatomic molecules The theory of the spectra of diatomic molecules is discussed in Herzberg (1950). In the brief summary presented here only aspects of the theory relevant to the later modelling of the spectra are presented. 5.1.1 Energy level structure The energy of a particular level in a diatomic molecule can be approximated by the sum of three parts. The electronic, vibrational and rotational energies each of which is approximately two orders of magnitude smaller than the last one. 103 104 Chapter 5. Molecules in SNe Figure 5.1: Potential curve for electronic energy level Electronic energy As two atoms approach each other the valence electrons of each experience the attractive potential of the other nucleus. As the separation of the nuclei decreases repulsive forces begin to dominate giving rise to a potential of the form shown in Fig. 5.1. The situation for a molecule is similar to that for an atom, where a given electron configuration gives rise to several different energy levels characterized by their resultant angular momenta. There is, however, a significant difference due to the presence of an electric field along the internuclear axis. The sum of the individual orbital angular momenta of the electrons quantized along this axis give a resultant A which takes values 0, 1, 2, 3... designated E ,n ,A __ The individual electron spins form a resultant S as in atoms. If S is coupled to the internuclear axis then the axial component of S is designated E and couples to A to form a resultant 12 = |A + E|. The electronic states are further complicated by symmetries of the electron orbitals but these are beyond the scope of this work. Chapter 5. Molecules in SNe 105 Vibrational energy When two nuclei with zero kinetic energy are in a potential such as the one shown in Fig. 5.1 they will have separation Ro, while if they have kinetic energy Ev they can oscillate between Ri and i? 2- The potential in Fig. 5.1 can be approximated by a potential of the form: V = Dc{ 1 - exf[-P(R - J?o)]}2 (5.1) where V is the energy referred to the minimum of the curve, De is the depth of the well and /3 is a constant. Expanding equation 5.1 for small displacements we get: V = Dcp \R - JJo)2(l - ^ R ~ Ro) + • • •) The first term gives a parabolic potential which when used in Schrodinger’s equation gives the following possible values of vibrational energy Ev = hvQ(v -f 1/2) where v takes integral values 0, 1, 2,... and the energy levels are equally spaced. In practice when the jt3fr-\ o p "VUL ; potential, oiQ, Viewing a diatomic molecule as two nuclei linked by a rigid bar we may make an analogy with the classical rigid rotator and use the reduced mass fi and the moment of inertia I(= /ztq) as the parameters defining the rotational energy of the system which then takes discrete values: h2J(J + 1) Er = 8n2I where J is the angular momentum of the system. However a rigid rotator cannot exhibit vibrations and therefore we need to include non-rigidity terms. The rotational energy then takes the form: F(J) = BVJ(J + 1) - DVJ2(J + l )2 cm 1 (5.3) 106 C hapter 5. Molecules in SNe Figure 5.2: Hund’s coupling cases. case (a ) case (b ) N where Bv and Dv depend on the moment of inertia and ue. The angular momentum considered so far comes entirely from rotation of the nuclei about each other. However, the total angular momentum is the resultant of contributions from coupling of the rotational, spin and orbital angular momenta. Depending on whether or not the spin angular momentum is coupled to the internuclear axis we distinguish two coupling cases known as Hund’s cases (a) and (b). The vector diagrams in Fig. 5.2 are illustrative of the two coupling cases. In case (a) J = ft + N where N is the angular momentum from nuclear rotation. In case (b) the spin S is not coupled to the internuclear axis therefore ft is not defined and J = K -|-S = (A + N)-fS.In case (a) the energy levels are the same as those described by equation 5.3 but shifted by a constant factor. In this coupling case levels with J< ft will be missing. For case (b) the energy levels are described by equation 5.3 with quantum number K replacing J with an additional additive term describing the effects of spin splitting. Chapter 5. Molecules in SNe 107 The total energy Since Ee > Ev >> Er we can write the total energy of each level as the sum of each individual term. So combining equations 5.2 and 5.3 we get: E — Ee + Ev + Er = Ee + + 1/ 2) - uexe(v + 1/2)2 + uicyc{v + 1/ 2)3 +BVJ(J + 1) - DVJ2(J + l )2 cm"1 (5.4) However, as mentioned above additional terms need to be considered in equation 5.4 depending on which of Hund’s coupling cases applies. 5.1.2 Selection rules and band structure The selection rules for rotational-vibrational transitions within a given electronic level for a diatomic molecule are dependent on the electronic state in which these transitions are occurring. In general however the strongest lines will be the results of transitions obeying the following rules: Aw = ± 1, (5.5) and A J = ±1^ O (5*6) ■ Thttfct+C-o purely rotational transitions (Aw = 0) occurring predom inantly in the millimetre waveband Qijrotational transitions combined with a change in vibrational energy level found in the near and mid-infrared. Note that every vibrational transition is associated with a change in rotational quantum number therefore the transi tions from vibrational state v to v' can be split into the P and R branches where A J = — 1 and +1 respectively. These branches extend either side of the ‘band origin’ where the J = 0 —► 0 transition would occur if it were allowed, with the R branch being at the short wavelength side. Due to the fact that the rotational energy levels become more widely spaced with increasing rotational quantum number and the spacing becomes smaller with increasing vibrational quantum number the R branch transitions reach a lower limit in wavelength at a particular value of J. This effect results in the formation of a sharp cutoff at one edge of the spectrum known as a ‘band head’. The bands are labelled according to 108 Chapter 5. Molecules in SNe the change in vibrational quantum number with the Av = 1 known as fundamental bands, Av = 2 as the first overtone and so on. Since the spacing between a vibrational energy level and its nearest neighbour is weakly dependent on vibrational quantum number the Av = 1 transitions will be closely packed with the v = 1 —> 0 being at the short end of the band and similarly for the Av = 2 bands. 5.1.3 Radiative rates The radiative transition rate for any transition between levels u and / is given by: _ 3 hg 1 1 (5.7) where vui is the frequency of the emitted radiation, gu is the degeneracy of the upper level and Ru/ is the transition moment for the two energy levels given by: R u' = f 4>"uMiPidT where ipi is the eigenfunction of the lower energy level and M is the .. moment of the system. It is possible to represent the wavefunction as the product of a vibrational term and a rotational term ip(0, ^ = IpeMr which allows us to calculate the transition moment as the product of a radial and an angular term. It is therefore possible to split the transition moment into two parts. Sw'JJ' = \RT \2SjJ' (5.8) where R™' is the overlap integral of the vibrational and electronic wavefunctions and S j j i , the Honl-London factor, is the square of the overlap integral of the rotational wave- functions. Kovacs (1969) has calculated Honl-London factors for transitions of diatomic molecules. No similar tables for R”v' exist but theoretical calculations and experimental measurements for some molecules can be found in the literature. As the Sjjt term is independent of vibrational quantum number the spectra of different bands will differ in overall intensity but not in shape. As the selection rule for |Av| > 1 results from the anharmonic terms in the potential these bands will be weaker than the fundamental. Chapter 5. Molecules in SNe 109 5.2 Carbon monoxide emission in SN 1987A 5.2.1 The discovery of emission bands As a number of observers have been studying SN 1987A there does not exist a unique first observation of the emission bands of Carbon Monoxide (CO1) in the literature. First over tone emission bands of the ground state of CO were identified in the spectra of SN 1987A around day 150 after the explosion by Catchpole & Glass (1987), McGregor & Hyland, (1987), Oliva et dl. (1987) and the Ames Research Centre & Cohen (1987). The first overtone bands of CO lie in the K atmospheric window around 2.3 pm. and are present in our spectra presented in chapter 3 as early as day 112 after the explosion (see Fig. 3.16.) As mentioned above the fundamental bands are potentially much stronger than the first overtone bands and for CO are centred around 4.6 pm in the M atmospheric window. A photometric excess was observed in this window from day 100 and has persisted to the present date. Spectra of this window taken by us on day 192 (see Fig. 3.22) and by Oliva et al. (1988) around day 225 are consistent with emission from the fundamental bands of CO. KAO spectra of this region (Dwek 1988b) taken around day 260 clearly show the fundamental band. The observed intensity ratio of the fundamental to the first overtone bands is however an order of magnitude lower than might be expected. However consid eration of the energy levels and the oscillator strengths for these bands and the conditions in the ejecta suggests that the CO may be optically thick to the Av = 1 photons while remaining optically thin to the Av = 2 (Oliva et al. 1987). The second overtone bands lie in the H window but as they are expected to have about 1.4% of the intensity of the first overtone bands (Bouanich & Brodbeck 1974) it is not surprising that they cannot be clearly identified in our spectra. In addition to the CO emission Petuchowski et al. (1989) suggested that emission from CO+ was also present in the K window, corresponding to the previously unidentified feature at 2.26 pm. This is difficult to verify since the fundamental and other overtone transitions are either in poorly covered parts of the spectrum or very weak. However, successful modelling, discussed later, of the observed spectra by invoking CO and CO+ increases the significance of the CO+ identification. 1We shall refer to 12C160 as CO from here on and any references to other isotopes will be explicit. 110 C h apter 5. M olecules In SN e 5.2.2 The location of the emitting CO In order to interpret the CO spectra, we must first ascertain its location. Specifically, did the CO lie in the progenitor wind, or did it form in the supernova ejecta ? Chevalier (1988) has argued that observed narrow UV lines originate from a swept-up shell of gas produced by the interaction of the winds from the blue supergiant progenitor and the previous red supergiant. He estimates the radius of the interaction zone to be > 2.4 x 1017 cm. Any CO lying in the shell might be excited by the initial UV flash. However, the wind velocity would be of order a few tens or hundreds of km s-1, resulting in sharp features at the band heads. However, in the spectra of the first overtone bands (see Figs 5.3-5.7) sharp band heads are not seen. Could this simply be the result of the low resolution (A/A A ~ 1300) which was used when the spectra were taken? Thev = 2 —► 0 band head occurs at rotational quantum number ~ 54 which means that up to 60 strong lines are present between the band origin and the band head which are only 0.05 pm apart. In order to look for such narrow features, in March 1988 we took spectra of the first overtone bands of CO with FIGS in third order which at 2.3 pm gives resolutions of (A/AA ~ 2200). At these resolutions we did not observe a consistent set of narrow features indicating that the CO emission was intrinsically broad. Note that this part of the K atmospheric window is contaminated by many narrow telluric features and therefore some isolated spikes do occur. Moreover, the spectral shape should have remained constant for the travel time of the UV flash through the shell. In contrast, the observed CO spectrum is changing continuously. Shock excitation by fast moving ejecta of CO lying in the shell is also implausible: at the highest velocity seen, ~ 30,000 km/s (Hanuschik & Dachs 1987), the travel time to the shell exceeds 2.5 years. We would in any case expect the shock to cause dissociation of CO. We conclude that the CO forms part of the ejecta. 5.2.3 Modelling the emission from CO We calculated a model vibrational-rotational spectrum for the ground state of CO, 13CieO and CO+. The energy levels were calculated using equation 5.4 using parameters for CO from Young (1969), for 13C160 from Benedict et dl. (1962) and for CO+ from Herzberg (1950). Note that the ground state of CO is a *£ in which case there is no point in distinguishing between Hund’s cases (a) and (b). CO+ has a 2E ground state (Hund’s Chapter 5. Molecules in SNe 111 case (b)) and has rotational energy levels defined by quantum number K rather than J. In 2E —2 E transitions the rotational selection rules 5.5, 5.6 apply with quantum number K substituted for J and with the additional selection rule A J = 0, ±1 where J = K ± However since AK = 0 is forbidden the resulting spectrum does not have a central band, arising from the A J = 0 selection rule, but weak satellite lines next to the P and R branch lines. Since our lines are intrinsically broad these effects are not obvious in the spectra but for completeness they have been included in the calculations. The radiative transition rates for the two isotopes of CO were calculated using equation 5.7 with values for the transition matrix from Chakerian & Tipping (1983). For CO+ we used equation 5.8 with values for R™' from Rosmus & Werner (1982) scaled to match the more recent calculations of Marian et dl. (1989) and values for Sjji from Kovacs (1969). It is important to note that the ground state of CO+ is a doublet which results in two R and P branches be < present in the spectra. The AJ = 0 satellite lines are very weak and are only included for completeness. We assumed a Boltzmann population for the energy levels. The line intensity for the transition from level u to level / is then given by: tEu/kT Iui = ------huulAul where N is the number of molecules emitting, Eu the upper energy level and gu the degeneracy of that level. Q is the partition function of the molecule which given the description of the energy levels as a sum of terms can easily be shown to be Q ~ QeQvQr where Qe, Qv and Qr are the electronic, vibrational and rotational partition functions respectively. The final spectrum was calculated by convolving the line spectrum with a parabolic line profile representing the uniform density optically thin case (see chapter 4). However in our spectra isolated lines tend to exhibit approximately Gaussian line profiles. Adoption of such line profiles does not affect the model spectra significantly due to the large number of lines being blended (on average 200 lines per band and up to 15 bands for each molecule). H the CO was confined to a shell in the ejecta then the line profile would be flat-topped. Spectra using such a line profile do not fit the data well. In particular the band heads can not be clearly separated. In calculating the absolute flux from the CO and CO+ we have taken the distance to the LMC to be 50 kpc. The fitting procedure for our models was as follows. First we determined the underly ing continuum using the method described in chapter 3. We then used the ratio of the CO 112 Chapter 5. Molecules In SNe Table 5.1: CO model parameters Epoch Continuum Temperature Velocity CO Mass CO+ Mass days Erg/cm2//mi K km/s M0 M0 192 2.38x10-® 3000 2000 1.4X10"5 3.6X10-® 255 1.00x10"® 1800 1200 3.85X10” 5 1.85x10"® 284 0.70x10"® 1600 1200 9.3x10-® 1.85X10-® 349 0.27x10"® < 1300 1200 > 4 x 10~® > 2 X 10~® band bead intensities to determine the temperature corresponding to the Boltzmann pop ulations. We assumed a common temperature for the CO and CO+ and used the relatively isolated v = 2 —* 0 CO+ band head to determine the width of the parabolic line profile. The mass of the C0+ was determined by matching the intensity of the same isolated feature to the model intensity. The mass of the CO was then determined by matching the combined CO and CO+ model spectra to the observed spectrum. The model spectra are compared with the data Figs 5.3-5.7 and the parameters used are listed in Table 5.1. The uncertainty in the derived temperature is of order 200 K. The uncertainty in the velocity for parabolic line profile case is of order 200 km/s. In Fig. 5.5 we compare a model spec trum with no CO+ to the data taken on day 255 after the explosion. Comparing Figs 5.4 and 5.5 it is obvious that an improvement in the fit is obtained through the inclusion of CO+ in our models, although the improvement is not dramatic. The model spectra in Figs 5.4 and 5.5 have the same velocity and temperature parameters but slightly differing masses of CO. Reasonable fits to the data were achieved without the inclusion of 13C160. We found that the addition of emission from this isotope produced degradation of the fits for 12C/13C less than 10. This lower limit is similar to those found in novae (Ferland et al. 1979) and consistent with the solar ratio of 80 (Cameron 1968). It is immediately obvious from Fig. 5.7 that the model fails to fit the day 349 spectrum. In this spectrum the band heads originating from levels higher than v = 2 are not seen. This evolution is consistent with the overall change in the spectral shape of the bands indicating a cooling of the ejecta. For any specific rotational-vibrational transition, indi vidual component pairs of the P and R branches originate from the same upper level (such Chapter 5. Molecules in SNe 113 Figure 5.3: Supernova (dotted line) and model (continuous line) spectra for day 192. N os utlfl W CO IO .< tin wavelength (^m) as {R(0) and P(2)} or {R (l) and P(3)} and so forth) and therefore have fixed intensity ratios. Moreover, the ratio is not strongly dependent on rotational quantum number. It follows that for a specific band, say v = 2 —» 0the ratio of the P and R branches will be insensitive to the ambient temperature. Since by day 349 the v = 3 —> 1 branch had practically disappeared the ratio of the 2.29 pm. feature to the 2.34 pm feature should have been dose to the ratio of the R to P branches. However the models to not produce fits for the observed spectrum. Either emission from an as yet unidentified feature is present at 2.26 pm or else there is absorption occurring at 2.34 pm. The only candidate suggested so far for the additional emission at 2.3 pm is the Ar II 2.313pm line (Whitelock et al. 1988) which requires some, as yet unspedfied, pumping mechanism as its excitation potential is 189415 cm-1. It is unlikely that the required absorption would be present in an emission line spectrum such as the one observed on day 349. 114 Chapter 5. Molecules in SNe Figure 5.4: Supernova (dotted line) and model (continuous line) spectra for day 255. B N B o U-t wavelength (//m) 5.2.4 Interpretation of the CO emission For days 192, 255 and 284 the quality of the model fits is good. Two immediate results are obvious. The temperature describing the populations in the molecules is lower than that inferred for the electron gas exciting the forbidden lines of iron (see chapter 3) and the expansion velocities inferred from the line width are lower than those inferred from the atomic line widths. These low expansion velocities indicate that some stratification has been maintained in the ejecta or that the CO only forms deep in the ejecta. The most plausible populating mechanism that would result in a Boltzmann distribution is colli- sional excitation as also suggested by McGregor & Hyland (1987). However, two other mechanisms are almost certainly present: infrared pumping by the optically thick funda mental transitions (t\e. v = 0 —» 1 and 1 —► 2 are converted into a single overtone photon v = 2 —*■ 0) and UV pumping of the A*H electronic level in CO by 1600 A photons fol lowed by radiative decays to excited vibrational levels (Frank-Condon principle). All these processes and their effects on the emergent spectrum have been considered in detail by Scoville et ol. (1980) and Krotkov et dl. (1980). They find that the vibrational populations Chapter 5. Molecules in SNe 115 Figure 5.5: Model spectrum with no contribution by CO +. wavelength (/j,m) of levels connected by optically thin transitions can be described by a single temperature even when the levels axe not thermalized. This results in an underestimate of the tempera ture of the exciting gas when photon trapping becomes important. The infrared pumping mechanism makes use of the fact that vibrational energy levels are almost equally spaced and therefore a v = 1 —► 0 photon can excite v — 1 —» 2,v = 2 —>3 and so on. Since the fundamental transitions are optically thick this excitation mechanism will enhance the populations of the vibrational energy levels. However given the evolution of the spectrum of the first overtone bands and the strength of the fundamental bands in the M window (see above) it is unlikely that this is a principal populating mechanism. Krotkov et al. (1980) find that if the UV pumping mechanism were the primary populating mechanism the spectrum would exhibit bands originating from vibrational levels higher than v = 4 which is not consistent with our observations. In addition there is no evidence for a strong ~ 1600 A pumping flux. However, if the CO were formed in an excited electronic state or through recombination of CO+ the subsequent cascades could enhance the populations of the upper vibrational energy levels. We conclude that Non-LTE populations are probably present and that collisional excitation is the dominant populating mechanism. 116 Chapter 5. Molecules in SNe Figure 5.6: Supernova (dotted line) and model (continuous line) spectra for day 284. 6 3. S W o wavelength (//m) 5.3 Other molecules in SN 1987A CO is the most tightly bound molecule with a dissociation energy of 11.09 eV (Douglas & M0ller 1955). CO+ has a dissociation energy of 8.33 eV making it the most tightly bound ionised species. When searching for emission from other molecules in the spectra of SN 1987A the rationale was to look for other tightly bound diatomics. CS and SiO with dissociation energies of 7.6 and 7.4 respectively were considered realistic candidates (Dwek 1988c). SiO fundamental emission was present in the spectra of SN 1987A in the 8.1 pm band around day 165 (Aitken 1988). In our spectra the first overtone bands of SiO lie at the same wavelength as Brackett a and therefore cannot be clearly identified. However, a broad flat-topped feature at ~ 3.9 pm is consistent with emission from the first overtone of CS. To date, molecular emission from other species has not been identified in the spectra of SN 1987A. Chapter 5. Molecules in SNe 117 Figure 5.7: Supernova (dotted line) and model (continuous line) spectra for day 349. wavelength (//m) 5.4 Molecular formation in SN 1987A It is improbable that any molecule in the progenitor atmosphere would survive the explo sion of a supernova. We therefore need to consider the various mechanisms for making CO using the constituents of the supernova after the explosion. The formation of molecules in the ejecta of SN 1987A has been discussed by Petuchowski et dl. (1989) and Latter & Black (1989). They consider three-body association, radiative association, radiative and dissociative recombination, charge exchange, electron impact excitation and ionisation, photodissociation and photoionisation processes. Their calculations indicate that enough CO can be produced via these processes to explain the masses inferred by our models. The formation of CO+ is not as yet clear. Lepp & Dalgamo (private communication) in their calculations make very little CO+ in conffhfwith the results of Petuchowski et al. (1989). Note that model spectra with no CO+ provide good fits to the data longwaxds of 2.28 pm (see Fig 5.5 and Spyromilio et al 1988). Therefore an alternative identifi cation of the 2.26 pm feature would be acceptable to the author. In order to maintain the population of CO and CO+, it is necessary to ‘shield’ it from the UV flux resulting 118 Chapter 5. Molecules in SNe from the two-photon decays of the metastable helium 2*S level (Lepp & Dalgarno, private communication). This requires that the ejecta are not microscopically mixed. In addition Petuchowski et al. need to constrain the electron density such that the fractional electron density in the immediate vicinity of the molecular material is < .1% which is an order of magnitude lower than the typical value throughout the ejecta derived by Moseley et al. (1989) from continuum measurements. This tends to support a scenario in which the ejecta are clumpy. A particularly interesting result of these calculations has been the identifica tion (Dwek 1988b) of rotational-vibrational transitions within the a3ic excited electronic state at 2.95 pm. This state is the lowest electronic state in CO and is metastable. It is populated by recombinations of CO+ to the triplet energy levels of CO. These transitions are not seen in our spectra possibly due to problems in the cancellation of the numerous telluric features present in that spectral region. Chapter 6 Conclusions and further work What have we learned about supernovae through studying their infrared spectra? The early-time infrared spectra of SN 1987A have allowed us to determine the population structure of the hydrogen in the ejecta going up to levels with principal quantum number 15. Compared to the optical which is limited due to line blanketing and blending this gives us a unique picture of the supernova ejecta. For the use of supernovae as distance indicators it is crucial that the conditions in both the continuum and line formation regions are known. Our work in modelling the early-time infrared spectra has indicated large deviations from Boltzmann distributions for the populations of the hydrogen energy levels. This is in agreement with the results of Hoflich for the optical region. It is obvious that careful consideration of the line formation mechanisms is needed when using the Baade-Wesselink method in determining the distance to supernovae. The velocity of the trough in the P-Cygni line profile does not necessarily give the velocity of expansion of the photosphere. The late-time spectra of SN 1987A provide a large variety of interesting features. The presence and evolution of the oxygen 1.129 /im resonance line may allow us to determine the evolution of the optical depth of the Lyman (3 line, which in turn affects recombina tion spectrum of hydrogen and as such is of great importance. The ratio of the helium 10830 A and 20580 A lines is a powerful diagnostic of the conditions in the ejecta and should be used as such. The detection of emission from radioactive Co in the H window (Varani, Meikle, Spyromilio & Allen in preparation) should allow us to place limits on the abundance of 57Co in the ejecta. Unfortunately SN 1987A could not resolve the question 119 120 Chapter 6. Conclusions over the identification of the 1.64 pm feature in the spectra of SN 1983N (see Graham et al (1986) and Oliva (1987)). The author’s opinion is that by comparison with SN 1987A and considering the explosion models of Ensman & Woosley (1988) the feature in SN 1983N was probably a blend of Si I and Fe II. Our late time spectral code falls short of its objectives due to the approximate treat ment of the radiation field. However useful insights to the nature of the late-time supernova have been gained. Neither Axelrod nor Meyerott considered the stable isotopes of Fe as coolants. This makes a dramatic difference to the temperature and ionisation structure derived. Our spectra of SN 1986G, in particular the spectrum taken at CTIO on the 16th of March 1987, may place strong limits on the amount of iron in Type la supernovae. In the future we would like to be able to derive the abundances of the elements ex hibiting transitions in our spectra. This however cannot be done without understanding how the spectrum is formed and the effects of interactions between different elements. The code described in chapter 4 needs to be expanded. As mentioned in chapter 4 most of the code is in a form that will allow such development. The transfer of the UV radiation field will have to be dealt with in the optically thick limit of the Sobolev approximation. The transitions of most species in the supernova will have to be included and processes such as charge transfer need to be considered. Good atomic data for these processes are needed before such a code is developed. In addition the effects of formation of molecules and dust will also have to be included. Density inhomogeneities axe almost certainly present in the ejecta of SN 1987A. To date no spectral code known to the author includes ‘clumping’ even though the ionisation and heating of the ejecta may be dominated by the conditions in these clumps. Clearly work remains to be done in the late-time spectral modelling field. One of the most exciting results of SN 1987A has been the discovery of molecular emission. However it should not have surprised the observers since molecular emission from CO displaying a very similar spectrum to that seen by us was seen in nova NQ Vulpeculae by Ferland et al as early as 1979. What is surprising is how little work has been done on CO in novae since then. In SN 1987A it would be interesting to use a Non- LTE code including all the effects of recombination, pumping and collisional processes to compute the level populations. 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The first year. W.P.S. Meikle, D.A. Allen, J. Spyromilio & G.-F. Varani, Mon. Not. R. astr. Soc., 238, 193, (1989). 3. Spectral line profiles of Iron and Nickel in Supernova 1987A. Evidence for a frag mented Nickel bubble. J. Spyromilio, W.P.S. Meikle & D.A. Allen, Mon. Not. R. astr. Soc., (in press). 4. Helium abundance and 0/tfm etiy in the wind from the precursor to SN 1987A. / D.A. Allen, W.P.S. Meikle & J. Spyromilio, Nature, (submitted). 129 Spectral line profiles of Iron and Nickel in Supernova 1987A. Evidence for a fragmented Nickel Bubble. J. Spyromilio and W.P.S. Meikle Blackett Laboratory, Imperial College of Science, Technology and Medicine, Prince Consort Road, London SW7 2BZ David A. Allen Anglo-Australian Observatory, PO Box 296, Epping, NSW 2121, Australia. 130 Spectral line profiles In SN 1987A 131 Summary We present line profiles of iron-group elements seen in near-infrared spectra of SN 1987A. The profiles are strikingly different from the observed transitions of other elements, but quite similar to those observed in the mid-infrared transitions of the iron-group elements. We are able to rule out blends, optical depth effects, and electron scattering as the source of the distinctive profiles. We propose that the iron-group elements have a unique spatial distribution in the ejecta, and that the profiles were due to emission from a low-density inner region (the nickel bubble) which has fragmented into high-velocity bullets. Dramatic changes in the profiles of these lines were seen about 650 days after the explosion. We cannot attribute these changes entirely to the formation of dust, as has been claimed for similar alterations in optical line profiles. Instead we argue that the abrupt change of the iron-group line profiles is primarily the result of faster cooling of the ejecta on the fax hemisphere from earth. 132 Spectral line profiles in SN 1987A 1. Introduction An important consequence of 56Ni decay in the interior of SN 1987A is the formation at an early phase of a low-density ‘nickel bubble’ (e.g. Woosley 1988). The boundary between this bubble and the denser overlying material would be subject to Rayleigh- Taylor instabilities which could result in fingers or ‘bullets’ of 56Co and other iron-group elements penetrating the outer layers. The detailed behaviour of the nickel bubble is poorly understood. The size of any bullets produced, their velocities and density contrast axe unknown. Indirect support for nickel bubble penetration of the outer layers is provided by the strong evidence for mixing in the ejecta of SN 1987A. Mixing of 56Co into the outer layers has been invoked to obtain a reasonable model fit to the bolometric light curve (Nomoto & Shigeyama 1988; Woosley 1988), and to account for the early appearance of hard X-rays (Itoh et dl. 1987) and 7-ray lines (Pinto & Woosley 1988) as well as the observed 7-ray line ratios (Matz et dl. 1988). Whitelock et al. (1988) estimate that 6 X 1048 ergs was deposited in the bubble, or about 40% of the total 56Ni decay energy. Nevertheless, direct confirmation of the bubble phenomenon is desirable through observations that have the necessary penetrative power and spatial resolution to allow a detailed examination of its behaviour. Rayleigh-Taylor fragmentation of the ejecta might reveal itself as a departure from spherical symmetry in the spatial distribution of the emission. Speckle interferometry (Karovska et al. 1988) and spectropolarimetry (Cropper et al. 1988) have already shown that the supernova appeared elongated. However, this was possibly due to scattering in an asymmetric envelope rather than to any intrinsic elongation of the ejected heavy element distribution (Cropper et al. 1988). In contrast, high-resolution spectroscopy has the potential to provide direct information about the nickel bubble. Not only can the nucleosynthesised products both of the progenitor and of the explosion be studied via abundance measurements, but also the spatial distribution of the ejected material can be determined. This is possible because the expansion becomes increasingly homologous (v ~ r) with time. Provided that the optical depth is low, velocity structure in Doppler- broadened line profiles can be used as a diagnostic of the physical location of the material (Fransson & Chevalier 1989). Spectral line profiles in SN 1987A 133 Optical emission line profiles have already provided evidence for departure from spheri cal symmetry. Satellite features which appeared at early times in lines of hydrogen (Blanco et dl. 1987; Phillips 1988; Larson et al. 1987,1988) and sodium (Hanuschik et al. 1988) have been interpreted as evidence for significant deviations from a spherical distribution of the emitting regions (Phillips et al. 1989). At about age 350 days the [O I] 6300 and 6363 A lines exhibited a series of ripples on the profiles that matched in velocity space (Stathakis & Cannon 1988). A plausible explanation of these features is the existence of inhomogeneities in the ejecta. At mid-infrared wavelengths (5-25 pm), the mass of iron inferred from forbidden lines observed around day 260 was a factor of three less than that derived from the light curve (Dwek 1988; Catchpole et al. 1988). Had the iron been uniformly distributed, the ejecta should have been optically thin in these lines. A plausible explanation for this discrepancy is that most of the iron was concealed in optically thick clumps (Dwek 1988). Spectroscopy at A/A A >1000 in the 0.7 to 5 pm wavelength region has the greatest potential for providing a probe of the slow moving (v <1000 km/s) ejecta, and allowing us to see through the supernova. At late times almost all infrared lines are collisionally excited by the thermal electron gas. The low temperature (T<5000K) of the electron gas means that relatively high excitation lines (E>10000cm-1), such as the near-infrared lines of the iron group elements, should be optically thin even in the inner, high density regions of the supernova. At other wavelengths, observations are handicapped by large continuum/line opacities, severe line blending, low signal/noise or poorer spectral resolution. 2. Observations We have monitored the spectral behaviour of SN 1987A at 1-4 pm using the infrared spectrometer FIGS (Bailey et al. 1988) attached to the Anglo-Australian Telescope. The techniques used in the acquisition and reduction of the spectra have been described in detail by Meikle et al. (1989a), and a comprehensive description of the data obtained during the second year is in preparation (Meikle et al. 1989b). Here we report only the observations of the emission line profiles. 134 Spectral line profiles in SN 1987A To investigate line asymmetries we selected strong and apparently isolated emission lines. Fig. 1 shows profiles of He I, [Fe II], Pa/?, and [Ni I] plotted in velocity space. The first three were obtained on day 494 and the nickel line on day 553. The spectral resolution A/A A was 1000—1200 for all lines except [Ni I], for which a resolution of about 2500 was used. This corresponds to a resolution of 3-4 pixels in all spectra. However, sharp features in the [Ni I] profile should be interpreted with caution as the raw data in this spectral region contain many telluric absorption features. The error bars, where shown, are 1 o and do not include uncertainties in the absolute flux. The wavelength calibration, derived from accompanying observations of a xenon comparison lamp, is accurate to 1-2 pixels (about 100 km/s). Inspection of Fig. 1 immediately shows some emission line profiles to be asymmetric about the velocity of the LMC rest frame. The [Fe II] and [Ni I] profiles have an inflection near the LMC velocity, but peak at about -flOOO km/s (+700 km/s in the LMC rest frame). A less pronounced peak at the same velocity is also present in the He I profile; blends on its long-wavelength edge with Pa 7, and [Si I] cannot account for this peak. The mid-infrared lines of the iron group elements around day 600 show similarly asymmetric line profiles to those seen in the near infrared. These include [Ni II] 6.6 pm, [Ni I] 11.3 pm and [Fe II] 17.9 pm (E.F. Erickson private communication to M. Cohen; R.J. Boyle, private communication). In contrast to these irregular profiles, the contemporary profile of Pa /? is quite symmetric. It is also clear that in the spectra taken on day 377 the Mg 1 1.503 pm line does not share the line profile of the [Fe II] line (Fig. 2). The temporal behaviour of the [Fe II] line profile for ages 192—735 days is shown in Fig. 3. On day 192 it had a simple flat-topped profile. The asymmetry described above had developed by day 284, and persisted for nearly one year. Between days 574 and 696 a striking change occurred: the redshifted maximum disappeared and gave way to a blueshifted peak at about -500 km/s. The evolution of the [Ni I] line profile is poorly known. The line became prominent by about day 400 (Allen et al. 1988), but only on day 553 was a high-resolution spectrum obtained. However, low-resolution spectra indicate an asymmetric profile with a redshifted maximum persisting from day 377 until at least day 574. Neither the hydrogen nor the magnesium lines developed irregular profiles in the period 192 to 735 days. Spectral line profiles in SN 1987A 135 3. Discussion 3.1 CAUSE OF LINE ASYMMETRIES An asymmetric line profile can arise through blending, optical depth effects, electron scattering from an expanding envelope, or asymmetry in the line-emitting region itself. In the present case blending can be ruled out by the similarity of all well observed iron and nickel lines in both the near- and mid-infrared regions, as well as by the improbability of sufficiently strong, unanticipated emission lines in these regions. We rule out the possibility that the iron-group line profiles resulted from the effect of optical depth in a spherically symmetric explosion. While some of the low lying fine structure transitions of the iron group elements (wavelengths 5-25 pm) were probably optically thick (Dwek 1988), the lines in the 1-2 pm region would have had an optical depth a factor of 103 lower, making them optically thin. Yet, these lines all share the same line profile. Electron scattering produces both an extended red wing and a redshift of the line centroid, even in spherically symmetric models (Fransson & Chevalier 1989; Woosley et al. 1989; Witteborn et al. 1989). However, symmetric models cannot generate the ripples and peaks that we noted (in section 1) were seen in optical lines of hydrogen, sodium and oxygen. In addition, there is evidence that the electron scattering optical depth was actually quite low. The profile of the mid-infrared line of [Ar II] (6.983 pm) was modelled by Witteborn et al. (1989) using spherically symmetric electron scattering. They found an optical depth r ~ 0.4 which, while accounting for the low- amplitude red wing to the line, is insufficient to explain the detailed line asymmetry we record. The steep gradient redward of the iron and nickel profile peaks (Fig. 1) is further evidence for a low electron scattering opacity. An additional argument against electron scattering in the present case is the constancy of the profiles. As noted above, the detailed shape of the [Fe II] line remained essentially unchanged for almost a year. During this period the electron scattering opacity must have fallen by a large factor due to both expansion and recombination, and had it dominated the line profile we would expect to have seen a marked reduction in the asymmetry. We are thus encouraged to attribute the line asymmetries primarily to emission from regions where the density and/or electron temperature depart significantly from spherical 136 Spectral line profiles in SN 1987A symmetry. The asymmetrical density/ temperature distribution in the ejecta could be a bulk property of the supernova. However, we favour the hypothesis that it arose from an asymmetric distribution of discrete high-velocity fingers of iron-group elements penetrating the outer layers i.e. we believe that the line profiles constitute evidence for ‘popping’ of the nickel bubble due to Rayleigh-Taylor instability. In support of this, we draw attention to the inflections near the LMC rest wavelength exhibited by both the iron and nickel lines (Fig. 1). A natural interpretation of these is that there is a paucity of material travelling at low velocities, and hence that the innermost regions of the ejecta have lower density. This appears to be a direct observation of the proposed low-density nickel bubble. The presence of a similarly irregular profile in the He I line suggests that at least some of the helium was co-extensive with the iron group elements. Yet, the lack of asymmetry in the magnesium profile seems to rule out simple mixing of the mantle as an explanation. We therefore suggest that some of the He I emission arose not from the outer layers, but from helium created during the photo-dissociation of the iron core and which escaped absorption into the forming neutron star. Such core-helium might be expected to have a similar distribution to that of the iron-group elements. 3.2 THE CHANGE IN LINE PROFILE Finally, we consider the dramatic change in line profile between days 574 and 696. At much the same time, Danziger et dl. (1989) reported the appearance of asymmetry in optical lines with peaks blueshifted by a similar amount. Similar changes are seen in the optical data of the AAO archive (R. Stathakis, private communication). Danziger et al. attributed the effect to the absorption by dust forming in the metal-rich ejecta. Support for this interpretation is provided by the AAO archive data: Mg I] A4571 shows much greater asymmetry than [O I] A6300 to a degree compatible with a standard interstellar reddening law. If the formation of dust alone accounts for the observed change in the optical profiles, then an increase in visual extinction of about 2 mag is required in the line of sight to the emitting gas (receding at 2000 km/sec) on the far side of the supernova. Even if the dust lies in clumps, therefore, those clumps must cover at least 80% of the path across the ejecta. Consequently, 40% or more of the optical luminosity of the supernova should now Spectral line profiles in SN 1987A 137 be appearing as thermal emission from the dnst. Using our own data and recent optical and infrared photometry from the South African Astronomical Observatory (Whitelock et al. 1989) we can place limits on the thermal emission from such dust. There is no indication that a thermal continuum characterised by temperatures around 1000 K developed after day 600. Indeed, at the very wavelengths where such a component would have been most prominent ( 2—4 pm) SN 1987A faded most rapidly. A thermal component of the necessary intensity, but which is consistent with the near-IR data, could therefore only have existed if the dust had been cooler than about 700 K in December 1988, or cooler than 500 K in late February 1989. The observed flux at longer wavelengths may permit the presence of such a cool compo nent. By October 1988 the bulk of the supernova luminosity appeared longward of 5 pm. This long wavelength radiation has been attributed to an infrared echo coursing through material shed in the red supergiant phase, (Roche et al. 1989). This interpretation is supported by the observations that most of the 10-pm emission was both extended and displaced by ~1 light year from the supernova (P. Roche, C. Smith & D. Aitken, private communication). Thus, the thermal radiation from dust forming in the ejecta could only have been a minor contributor to this radiation. Nevertheless, the measurement uncer tainties are such that there could have existed an unresolved, undisplaced fraction large enough to account for the flux lost through absorption at short wavelengths. The optical line profiles began to change at about the epoch of our day 574IR observa tions. Between days 574 and 696, most of the lines in our waveband exhibited drifts to the blue. The shift in the non-iron group lines is compatible with absorption of the red wings by dust having approximately interstellar characteristics and fitted to the optical Mg I] and [O I] lines. However the iron group lines in our waveband (1.257,1.533, 1.644 pm of [Fe II] and 3.119 pm [Ni I]) underwent a blueward shift of their peaks about five times greater than that of the other IR lines. This cannot have been primarily due to dust. For example, a visual extinction exceeding 5 magnitudes would be required to account for the change in the 1.257-/zm line by day 696, assuming an interstellar reddening law. This is far too high a figure to be compatible with the extinction implied by the optical and non-iron group IR data. An additional effect in the iron group lines therefore appears to have dominated. We propose that the profiles were modified by an intrinsic fading of the 138 Spectral line profiles in SN 1987A redder emission. In support of this hypothesis is the change that occurred before day 574 (Fig. 3). The strong red peak to the [Fe II] line seen, for example, on day 494 had already weakened by day 574, and further weakening of the red emission may have continued. The line emission is a function of temperature and density, so a change of either in one part of the ejecta would modify the line profile. The more likely parameter to change abruptly is temperature. Part of the changing line profile we see can be attributed to a more rapid cooling of some of the ejecta on the far side of the supernova. Rapid cooling preferentially occurs in the densest regions, where the emission is most intense. It would therefore most clearly manifest itself as a change of line profile in those lines which are strongly asymmetric. This is consistent with what we observe. We conclude that dust condensation is unlikely to have been a major factor in the change of profile of the iron group lines. Acknowledgments We thank Patricia Whitelock for providing unpublished data and their interpretation, and Raylee Stathakis for access to the AAO archive and for stimulating discussions. We also thank Sean Colgan, Philip Pinto and Stanford Woosley for helpful discussions. Spectral line profiles in SN 1987A 139 References Allen, D., Spyromilio, J., Meikle, P. & Varani, G., 1988. Int. ostr. Union Circ., 4623. 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Supernova 1987A in the Large Magellanic Cloud. Fourth George Mason Fall Workshop in Astrophysics , p.16, eds Kafatos, M. & Michalitsianos, A.G., 140 Spectral line profiles in SN 1987A Cambridge University. Phillips, M.M. & Heathcote, S.R., 1989. Pubis, ostr. Soc. Pocif., in press. Pinto, P.A. & Woosley, S.E., 1988. Nature, 333, 534. Roche, P.F., Aitken, D.K., Smith, C.H. & James, S.D., 1989. Nature, 337, 533. Stathakis, R. & Cannon, R.D., 1988. Anglo-Australian Observatory Newsletter No. 45. Whitelock, P. A., et al ., 1988. Mon. Not. R. astr. Soc., 234, 5 p. Whitelock, P.A., et al., 1989, in preparation. Witteborn, F.C., Bregman, J.D., Wooden, D.H., Pinto, P.A., Rank, D.M., Woosley, S.E. & Cohen, M., 1989. Astrophys. J., 338, L9. Woosley, S.E., 1988. Astrophys. J., 330, 218. Woosley, S.E., Pinto, P.A. & Weaver, T.A., 1989. Proc. astr. Soc. Aust., 7, 355. Spectral line profiles in SN 1987A 141 Figure captions • Figure 1. [Fe II] 1.257, Paschen /3 1.282 and He I 1.083 pm lines from day 494 and the [Ni X\ 3.119 pm line from day 553 in velocity space with respect to their rest wavelengths. The velocity of the rest frame of the supernova is also shown. To avoid overlap the profiles have been normalized in intensity and displaced vertically. The horizontal lines at the left indicate the zero flux levels. • Figure 2. [Fe II] 1.257 and Mg 1 1.503 pm line profiles from the day 377 spectra. • Figure 3. Evolution of the [Fe II] 1.257 pm line profile from day 192 to 735. normalized F 142 N 97 AT FIGS AAT 1987A SN pcrlln rflsi S 1987A SN in profiles line Spectral (0 normalized F 1987A SN in profiles line pectral S N 97 AT FIGS AAT 1987A SN [l) 143 Spectral line profiles in SN 1987A 144 SN 1987A AAT FIGS (3) Helium abundance and asymmetry in the wind from the precursor to SN 1987A D. A. Allen*, W. P. S. Meiklet & J. Spyromiliot * Anglo-Australian Observatory, PO Box 296, Epping, New South Wales 2121, Australia t Blackett Laboratory, Imperial College of Science, Technology and Medicine, Prince Con sort Road, London SW7 2BZ The initial UV flash from the surface of SN 1987A ionized circumstellar gas to produce narrow emission lines which were first seen at UV wavelengths1 where the supernova’s photospheric emission quickly faded. Later, optical lines were also reported2. The UV lines indicated densities Ne ~ 104 cm-3, far greater than expected for the wind shed by the precursor star, Sk-69°202, in either its blue supergiant (BSG) or preceding red supergiant (RSG) phase3. However, the interaction of the two winds would produce a shell of the req uisite density4'5. From observations of the infrared helium triplet at 1083 nm we demonstrate that the RSG wind was strongly asymmetric, and we verify predictions that its helium abundance was high. In the wind interaction 4,5 masses as high as a few Mq can be compressed to densities exceeding 104 cm-3 as a shell of radius about one light yr (ly ~ 1018 cm) and thickness 1015 cm. Gas interior and exterior to this shell is of such low density that it may be ignored. The intial UV flash is predicted to have released io56-57 photons isotropically in an hour or less6'7, sufficient to ionize almost 1 solar mass of gas. Confirmation of these figures was provided by the UV line ratios during the first 400 days (Ref. 3). At such densities recombination timescales are months or years. The intensity of an emission line at any point in the nebula evolves with time as the populations of its parent ions and of free electrons change due to recombination. A simple description of the changes was given by Chevalier4. The mass of ionized gas seen in a spherical shell of radius r ly centred on the supernova increases monotonically, due to light travel effects, until a time 2r yr after the UV flash, when the entire shell is seen. A steady increase was indeed observed in the early evolution 145 146 Asymmetry in the wind of SN 1987A of some of the UV emission lines. The angular diameter of the emitting material also evolves. Initially we see only gas lying immediately in front of the supernova. By time r yr we see material within the entire frontal hemisphere of the shell, subtending a radius r/D, where D is the distance to the supernova. Thereafter the rear hemisphere also contributes, but the angular size is constant (except for shell expansion). Almost two years after the eruption of SN 1987A, narrow emission lines were seen extending to a radius of about 2.5 arcsec (Ref. 8), from which a minimum value of r = 2 ly is inferred. However, we find that the shell is not spherical. Asymmetry in SN 1987A has already been found by several means, but mostly on smaller scales than we now report. In 1987 speckle observations 9,10 briefly revealed a companion (the ‘mystery spot’) about twenty light days from the supernova at position angle about 195°. At much the same PA spectropolarimetry 11 indicated an elongation or asymmetry of the supernova. When the ejecta themselves became resolvable 12 they too were extended along PA 20 — 200°. More recently8, velocity gradients were seen in the circumstellar gas; the greatest recession velocity is at about PA 190°. We have undertaken a neax-infraxed spectroscopic study of SN 1987A using the Anglo- Australian Telescope13. Since about mid 1988 we have recorded a narrow component to the 1083 nm He I line, which we demonstrate below must arise in the circumstellax shell. The evolution of its intensity is given in Table 1; in each case we measured the integrated flux through a 3.5 arcsec square aperture centred on SN 1987A. The naxrow component was always superimposed on a broad emission blend of He I (1083 nm) and Paschen 7 (1094 nm). Figure 1 shows two recent spectra. Table 1 also includes intensities for naxrow optical lines of relevance; they were taken from the AAT supernova archive (Stathakis, private communication) and have been adequately corrected for emission from diffuse nebulosity in the region. In 1989 Feb we explored the spatial distribution of the naxrow He I emission feature. We set the grating so that one of the 16 detectors in the spectrometer (FIGS) received radiation from the naxrow component while several straddled the broad component. By repeated raster scanning of the telescope we thus created simultaneous maps of the two components. The focal plane was sampled on a square grid with 0.7 arcsec spacing through a 1.4 arcsec squaxe aperture. Figure 2 shows grey-scale representations of the maps. The broad component is an unresolved point source, as expected since it is believed to arise in the supernova. Its contribution was removed from the observed flux in the naxrow Asymmetry in the wind of SN 1987A 147 component by scaling from a fully-sampled spectrum. The resulting image shows the distribution of the narrow component alone. It is clearly displaced from the supernova, by 0.8 ± 0.3 arcsec at PA 200 ± 15 deg, indicating a circumstellar origin. Approximately half of the total intensity of the narrow component arises in the displaced condensation, whose diameter is less than about 1.5 arcsec (1.2 ly). We also made north-south line scans through SN 1987A, using the same spectral and spatial sampling. Figure 3 reproduces spectra at different declinations from approximately 200 coadded scans. A similar displacement in declination is seen. Following our announce ment of this displacement14, Karovska et al.15 reported a source seen in a 35 nm continuum bandpass at the same location from 1988 Nov to 1989 Jan. Since, however, this feature is unconfirmed in direct imaging 16,17 we do not consider it further. The fluxes in optical lines (Table 1) were measured in a comparable focal plane aperture to ours, but the lack of exact match both spatially and in date of observation introduces some uncertainty into the analysis we now perform. The rapid fall in the ratio of 469 nm He II to H (3 indicates Ne ~ 105. The 1083 nm He I line arises both directly by recom bination of ionized helium and from collisional excitation of the metastable 23S level. At Ne ~ 105 the latter dominates and the line is stengthened relative to the recombination case by more than an order of magnitude18. Three arguments suggest that the condensation lies behind the supernova, and has therefore been viewed for a shorter time than the supernova itself. First, the circumstel lar gas at this PA is receding8. Second, a few days before our observation in February, Roche, Aitken & Smith (private communication) found the dust to lie in a compact region displaced from the supernova and, to within the errors, coincident with the helium conden sation. This emission is dominated by grains illuminated by the supernova near maximum light, about three months later than the UV flash19. The dust and gas emission differ not only in that delay but also because the dust cools instantaneously, so that its emission maps out only the present location of energy from the light-curve peak. Considerations of geometry and temperature indicate that the dust echo lies behind the supernova19’20. Third, the high density shows that the condensation did not provide the material emitting the UV lines during the first 1.1 yr. There is little evidence for variation of the 1083 nm line since it became spectrally resolvable from the broad component. Because the condensation lay near the edge of our 148 Asymmetry in the wind of SN 1987A entrance aperture, which was always centred on the supernova, the measured intensity would depend heavily on the pointing and tracking stability. For Ne ~ 105 we predict (below) almost constant line emission per unit volume from 3 months to 1 year after initial ionization. We conclude that the mass of emitting gas has not increased significantly during the period of observation, so that by March 1989 we had been viewing the same condensation for at least 0.8 yr. If we adopt a viewing time of 1 yr, the condensation lies 0.8 ly from the supernova. In the plane of the sky the shell extends at least 2 ly from the supernova and has lower density. Such an asymmetry cannot be produced by the second (BSG) wind, for if it expanded into an isotropic shell it could only accumulate high densities at larger radii. Thus it is either a manifestation of the RSG mass loss, or is caused by a dense interstellar cloud that has successfully resisted the momentum flux of the RSG wind. The latter view was taken by Felten & Dwek20. A bipolar configuration is an interesting alternative, but a symmetrical bipolar form is incompatible with our data since the frontal lobe would now be seen at age 2 yr, yielding a 1083 nm intensity 60% that of the rear lobe. A northerly extension to the emission is indeed seen in Figs. 2 & 3, but at lower intensity. Moreover, any such frontal lobe must be less dense than the rear if the density is to be compatible with the UV line ratios. We have constructed a simple model of the compact He I source comprising a slowly- moving cloud of hydrogen and helium only. We assume that it was fully ionized by the UV flash, and that light travel time across it is unimportant. We have allowed for re combination, collisional effects and ionization of neutral atoms by photons released in the recombination of He++. We take the He I line to be optically thin and use atomic pa rameters (from Refs 18, 21-25) for an electron temperature of 20 000 K, based upon a recent intensity ratio of the narrow 436.3, 495.9 and 500.7 nm [O III] lines (Stathakis pri vate communication). A higher temperature was found initially26, but cooling would have been rapid. Our only free parameters are Ne, and the relative proportions of the two ele ments N(H.e)/N(H). We correct for interstellar reddening27,28, and include contributions of He II to the Balmer lines. Our model satisfactorily reproduces all values in Table 1, except the high He II strength on 1988 Oct 28, if the number density of hydrogen is 6 x l04 cm-3 and JV(He)/iV(H) = 0.35, a figure about three times the solar value. We can fit the He II line rather better if we have been viewing the condensation for a shorter time. The undoubted presence of density Asymmetry in the wind of SN 1987A 149 inhomogeneities in the shell favours emission of Ha, because at low density the collisional excitation of the 1083 nm line disappears. Thus the abundance is itself a lower limit. We regard the figure as provisional until a more thorough analysis becomes possible of the geometry, and data at different wavelengths can be obtained with better spatial matching. It is unlikely that an interstellar cloud with so high a helium abundance would lie within 1 ly of the supernova. Evolutionary models 29 indicate that the last material shed in the RSG phase should have had N(Re)/N(R) ~ 0.6, so we interpret the condensation as part of an asymmetric RSG outflow. From the He I line intensity in a 3.5 arcsec square aperture we derive a total mass of ionized gas ~ 0.016 M®, an order of magnitude smaller than that derived from the UV lines3 for material in front of the supernova. This is also much less than the mass inferred from the dust emission29, suggesting that the UV flash was unable to penetrate the cloud completely. The cloud subtends up to 2 sr at the supernova, so the total number of photons capable of ionizing hydrogen or helium and radiated isotropically in the UV flash must have been at least 5 x 1056, in agreement with predictions. 150 Asymmetry in the wind of SN 1987A 1. Wamsteker, W., Gilmozzi, R. k Cassatella, A. Int. astr. Union Circ. 4410 (1987). 2. Wampler, E. J. Int. astr. Union Circ. 4541 (1988). 3. Fransson, C., Cassatella, A., Gilmozzi, R., Kirshner, R. P., Panagia, N., Sonneborn, G. k Wamsteker, W. Astrophys.J. 336, 429-441 (1989). 4. Chevalier, R. A. Nature 332, 514-516 (1988). 5. Chevalier, R. A. ESO Workshop on the supernova 1987A, Conf. Workshop 26, Proc. 481-494 (1987). 6 . Dopita, M. A., Meatheringham, S. J., Nulsen, P. k Wood, P. R. Astrophys. J. 322, L85-L89 (1987). 7. Lundqvist, P. & Fransson, C. ESO Workshop on the supernova 1987A, Conf. Workshop Proc. 26, 495-502 (1987). 8 . Wood, P. R. k Faulkner, D. J. Int. astr. Union Circ. 4739 (1989). 9. Nisensen, P., Papaliolios, C., Karovska, M. k Noyes, R. Astrophys.J. 320, L15-L18 (1987). 1 0 . Meikle, W. P. S., Matcher, S. J. & Morgan, B. L. Nature 329, 608-611 (1987). 11. Cropper, M. S., Bailey, J. A., McCowage, J. C., Cannon, R. D., Couch, W. J., Walsh, J. R., Straede, J. O. k Freeman, F. F. Mon. Not.R. astr. Soc. 231, 695-722 (1988). 12. Papaliolios, C., Karovska, M., Koechlin, L., Nisenson, P., Standley, C. k Heathcote, S. R. Nature 338, 565-566 (1989). 13. Meikle, W. P. S., Allen, D. A., Spyromilio, J. k Varani, G-F. Mon. Not.R. astr. Soc. 238, 193-223 (1989). 14. Allen, D. A., Meikle, W. P. S. k Spyromilio, J. Int. astr. Union Circ. 4747 (1989). 15. Karovska, M., Nisenson, P., Papaliolios, C., Standley, C. k Heathcote, S. R. Int. astr. Union Circ. 4749 k 4752 (1989). 16. Crotts, A. k Kunkel, W. E. Int. astr. Union Circ. 4753 (1989). 17. Heathcote, S. R., Suntzeff, N. B. k Walker, A. R. Int. astr. Union Circ. 4753 (1989). 18. Osterbrock, D. E. Astrophysics of gaseous nebulae, Freeman, San Francisco (1974). 19. Roche, P. F., Aitken, D. K., Smith, C. H. k James, S. D. Nature 337, 533-535 (1989). 2 0 . Felten, J. E. k Dwek, E. Nature 339, 123-125 (1989). 21. Jacobs, V. L. k Burke, P. G., J. Phys. B} 5, 2272-2281 (1972). 22. Bell, K. L., Kingston, A. E. k Taylor, I. R. J. Phys. B 6 , 2271-2279 (1973). 23. Hata, J. k Grant, I. P. J. Phys. B 14, 2111-2124 (1981). 24. Clegg, R. E. S. Mon. Not.R. astr. Soc. 229, 31P-39P (1987). 25. Clegg, R. E. S. k Harrington, J. P., Mon. Not.R. astr. Soc. in press (1989). 26. Wampler, E.. J. k Richichi, A. The Messenger 52, 14-16 (1988). 27. Dopita, M. A. Proc. astr. Soc. Aust. 7, 344-351 (1988). 28. West, R. M., Lauberts, A., J0 rgensen, H. E. k Schuster, H-E. Astr. Astrophys. 177, L1-L2 (1987). Asymmetry in the wind of SN 1987A 151 29. Maeder, A. ESO Workshop on the supernova 1981A, Conf. Workshop 26, Proc. 251-269 (1987). ACKNOWLEDGEMENTS We thank Raylee Stathakis for providing data from the AAT archive, and Robin Clegg, Peter Storey and Richard Learner for valuable discussions. Table 1. Observing log and line intensities (optical lines from AAT archive) Date Observed intensity in 10-13 erg cm -2 s-1 He I 1083 nm Ha 656 nm H/? 486 nm He II 469 nm 1988 May 1 13 ±6 1988 Jul 1 12 ± 5 1988 Sep 19 15 ±3 1988 Oct 28 1.1 ±0.4 1988 Dec 22 20 ± 5 1989 Jan 19 16 ±2 1989 Feb 12 0.7 ±0.2 0.4 ± 0.1 1989 Feb 27 13 ±2 1989 Mar 19 2.2 ±0.2 0.7 ±0.1 0.2 ± 0.05 Quoted errors do not include an additional uncertainty arising because the optical and infrared instruments accepted diffferent portions of the extended emission. 152 Asymmetry in the wind of SN 1987A Figure legends Fig. 1. Spectra of SN 1987A taken in 1988 September and 1989 January. The broad component is helium and hydrogen emission from the supernova ejecta. The narrow component is helium in the circumstellar shell. Fig. 2. Grey scale representations of maps of the broad (left) and narrow (right) compo nents seen in Fig. 1. The pixels are 0.7 arcsec square. Both maps have been automatically scaled to cover the full tonal range. Fig. 3. Six undersampled spectra from the declination scans through SN 1987A, taken at 0.7 arcsec spacing. The narrow component peaks in the second and third spectra from the bottom (south), and the broad component in the third and fourth. Asymmetry in the wind of SN 1987A 153 H« I in 1988 StfiJ and 1989 Jan Daclioalton slicta through SN I987A in Ha I lina I Ha in I987A SN through slicta Daclioalton Jan 1989 StfiJand 1988 in IH« 154 Asymmetry In the wind ofSN 1987A