Science with the Square Kilometer Array

Home , 3C 319, Flare star, Iota Tauri, Malmquist bias, Megamaser, SN 1181

edited by: A.R. Taylor and R. Braun

March 1999 Cover image: The Hubble Deep Field Courtesy of R. Williams and the HDF Team (ST ScI) and NASA. Contents

Executive Summary 6

1 Introduction 10 1.1 ANextGenerationRadioObservatory ...... 10 1.2 The Square Kilometre Array Concept ...... 12 1.3 Instrumental Sensitivity ...... 15 1.4 Contributors...... 18

2 Formation and Evolution of Galaxies 20 2.1 TheDawnofGalaxies ...... 20 2.1.1 21-cm Emission and Absorption Mechanisms ...... 22 2.1.2 PreheatingtheIGM ...... 24 2.1.3 Scenarios: SKA Imaging of Cosmological H I ...... 25 2.2 LargeScale Structure and GalaxyEvolution ...... 28 2.2.1 A Deep SKA H I Pencil Beam Survey ...... 29 2.2.2 Large scale structure studies from a shallow, wide area survey 31 2.2.3 The Lyα forest seen in the 21-cm H I line...... 32 2.2.4 HighRedshiftCO...... 33 2.3 DeepContinuumFields...... 38 2.3.1 ExtragalacticRadioSources ...... 38 2.3.2 The SubmicroJansky Sky ...... 40 2.4 Probing Dark Matter with Gravitational Lensing ...... 42 2.5 ActivityinGalacticNuclei ...... 46 2.5.1 The SKA and Active Galactic Nuclei ...... 47 2.5.2 Sensitivity of the SKA in VLBI Arrays ...... 52 2.6 Circum-nuclearMegaMasers ...... 53 2.6.1 H2Omegamasers ...... 54 2.6.2 OHMegamasers...... 55 2.6.3 FormaldehydeMegamasers...... 55 2.6.4 The Impact of the SKA on Megamaser Studies ...... 56 2.7 TheStarburstPhenomenon ...... 57 2.7.1 TheimportanceofStarbursts ...... 58 2.7.2 CurrentRadioStudies ...... 58 2.7.3 The Potential of SKA for Starburst Studies ...... 61

3 4 CONTENTS

2.8 InterstellarProcesses ...... 63 2.8.1 HII Regions: High Resolution Imaging of Thermal Emission . . 64 2.8.2 Centimetre Wavelength Molecular Probes of the ISM . . . .. 66 2.8.3 SupernovaRemnants ...... 68 2.8.4 TheOriginofCosmicRays ...... 74 2.8.5 Interstellar Plasma Turbulence ...... 76 2.8.6 RecombinationLines ...... 77 2.9 MagneticFields...... 78 2.9.1 RotationMeasureSynthesis ...... 78 2.9.2 Polarization Studies of the Interstellar Medium in the Galaxy andinNearbyExternalGalaxies ...... 80

3 Formation and Evolution of Stars 86 3.1 ContinuumRadioEmissionfromStars ...... 86 3.2 ImagingtheSurfacesofStars ...... 90 3.2.1 RedGiantsandSupergiantStars ...... 90 3.2.2 Complementarity to Planned Optical-IR Interferometers ... 97 3.3 StarFormation ...... 98 3.3.1 ProtostellarCores...... 99 3.3.2 ProtostellarJets ...... 102 3.3.3 Uncovering the Evolutionary Sequence ...... 105 3.3.4 Magnetic Fields in Frotostellar Objects ...... 107 3.4 CoolStarAstronomy ...... 109 3.4.1 TheRadioSun ...... 111 3.4.2 Observing Solar Analogs at Radio Wavelengths ...... 114 3.4.3 WherearethemanyotherRadioSuns?...... 115 3.4.4 ProspectsfortheSKA ...... 117 3.4.5 FlaresandMicroﬂares ...... 118 3.4.6 Summary of Scientiﬁc Objectives ...... 120 3.5 Imaging of Circumstellar Phenomena ...... 121 3.6 StellarAstrometry ...... 121 3.7 Supernovae ...... 122 3.7.1 RadioSupernovae...... 123 3.7.2 New Observations Possible with the SKA ...... 127 3.7.3 SummaryandConclusions ...... 129 3.8 TheRadioAfter-GlowsofGamma-rayBursts ...... 132 3.9 Pulsars...... 135 3.9.1 PulsarSearcheswiththeSKA ...... 136 3.9.2 PulsarTimingwiththeSKA ...... 137 3.9.3 Radio Pulsar Timing and General Relativity ...... 138 3.9.4 Summary ...... 139 CONTENTS 5

4 Solar System Science 141 4.1 Thermal Emission from Small Solar System Bodies ...... 141 4.1.1 Asteroids ...... 142 4.1.2 PlanetarySatellites ...... 145 4.1.3 KuiperBeltObjects ...... 146 4.2 RadarImagingofNearEarthAsteroids...... 149 4.3 The Atmosphere and Magnetosphere of Jupiter ...... 152 4.4 CometStudies...... 153 4.5 SolarRadar ...... 153 4.6 CoronalScattering ...... 154

5 Formation and Evolution of Life 155 5.1 DetectionofExtrasolarPlanets ...... 155 5.2 Pre-BioticChemistry ...... 157 5.3 The Search for Extraterrestrial Intelligence ...... 157

Bibliography 161 List of Tables

1.1 SKADesignGoals ...... 14 1.2 Instrumental Sensitivity per Polarisation in 8 hours ...... 16

2.1 Detectable H I Masses for an SKA Deep Pencil Beam Survey . . . . . 30 2.2 CODetectionsatHighRedshift ...... 35

3.1 Some Current and Planned Optical-IR Interferometers ...... 97

4.1 SKA Speciﬁcations for Imaging Small Solar System Bodies ...... 141

5.1 Parameters of Selected SETI Surveys ...... 159

6 List of Figures

1.1 ComparisonofFieldsofView ...... 12 1.2 TheSquareKilometreArray...... 13 1.3 SpectralLineSensitivity ...... 16 1.4 ContinuumSensitivity ...... 17

2.1 Detectability of fluctuations in H I brightness temperature...... 27 2.2 H I emission from the region surrounding a QSO source ...... 27 2.3 The Local H I MassFunction...... 30 2.4 Properties of Galaxy Redshift Surveys ...... 32 2.5 Observable CO Transitions as a Function of Redshift ...... 34 2.6 CO Luminosity as a Function of Redshift ...... 36 2.7 CO Flux Density as a Function of Redshift ...... 36 2.8 Signal to Noise Ratio as a Function of Redshift ...... 37 2.9 Simulated Continuum Observation of a region the size and shape of theHubbleDeepField ...... 41 2.10 Instrumental Distortion of HST Field Stars ...... 44 2.11 Smoothed Distortion Field toward CL 1358+62 ...... 45 2.12 Dependence of peak VLA power on distance for 374 UGC galaxies . . 48 2.13 H I absorptionin3C236...... 49 2.14 DeepVLBIImageof3C84 ...... 50 2.15 Compact Component Luminosity as a Function of Redshift ...... 52 2.16 ThecenterofNGC4258 ...... 54 2.17 A λ6cmMERLIN/VLAimageofM82 ...... 60 2.18 ExpansionofaYoungSNRinM82 ...... 61 2.19 Brightness Temperature – Angular Size Space ...... 65 2.20 RotationMeasureofPulsarJ0214+4232 ...... 79 2.21 Effects of the Diffuse Faraday Screen in our Galaxy ...... 82

3.1 TheRadioHRDiagram ...... 87 3.2 Distribution of Stellar Radio Luminosities ...... 88 3.3 DetectionofMiraPhotospheres ...... 89 3.4 Non-thermallyEmittingStars ...... 90 3.5 HST and Interferometer Images of Mira and Betelgeuse ...... 91 3.6 Image and Model of UV emission from Betelgeuse ...... 92

7 8 LIST OF FIGURES

3.7 Flux Density and Size for a Mira-like Red Giant Star versus Distance andFrequency...... 93 3.8 Flux Density and Size for a Red Supergiant Star versus Distance and Frequency ...... 94 3.9 VLAImageoftheBetelguese ...... 96 3.10 Angular Resolution of the SKA and Planned Optical-IR Inteferometers 98 3.11 Dust and Ionized Wind from the Protostellar Object HL Tau..... 99 3.12 OpacityofaSphericalDustCloud ...... 101 3.13 Collimated Outﬂows in Young Stellar Objects ...... 103 3.14 A λ3.6cmImageofDouble? JetsinL1551...... 104 3.15 Proper Motion of the Ionized Jet of HH80-81...... 104 3.16L1498 ...... 106 3.17 TMC1, Rosetta Stone forstar formation?...... 106 3.18TheRadioSun ...... 112 3.19 DynamicSpectrumofaSolarRadioBurst ...... 113 3.20 Radiospectraoftheactive dMestarUVCet...... 115 3.21 Dynamic Spectrum of a Coherent Radio Burst from the dMe Star AD Leo...... 116 3.22 Expected sensitivity range of the SKA ...... 119 3.23 Light curves at multiple radio frequencies for several extragalactic SN 125 3.24 Light curves at multiple radio frequencies for SN 1993J ...... 126 3.25 Peak 6 cm luminosity, L6 cm peak of RSNe vs. time...... 128 3.26 Peak RSNe λ6cmﬂuxversusredshift...... 130 3.27 The radio luminosity vs. age for a number of young radio supernovae 131 3.28 The broad-band radio/IR/optical/X-ray spectrum of GRB980703 . . 133

4.1 HSTImageoftheVesta ...... 143 4.2 Light Curve of Vesta at λ2mm...... 144 4.3 Kuiper Belt like Phenomena Around Nearby Stars ...... 147 4.4 RadarImagesofAsteroid4769Castalia...... 150 4.5 RadarImagesofAsteroid4179Toutatis ...... 151

5.1 Angular Motion of the Sun-Jupiter System ...... 156 5.2 Predicted and Observed Abundance of Carbon-chain Species ..... 158 Executive Summary

The Square Kilometre Array New developments in all fields of astronomy have brought the current generation of astronomers to the brink of probing the origin and evolution of the Universe as a whole. Planning for the next generation of facilitiies leads to the conclusion that a revolutionary new instrument at radio wavelengths is needed, one with an effective collecting area more than 30 times greater than the largest telescope ever built. Such a telescope will reveal the dawn of galaxy formation, as well as a plethora of other new discoveries in all fields of astronomy. Vigorous technological developments in computing and radio frequency devices make it possible for such a telescope to be built within the next decade, and the international radio astronomical community is proposing that such a telescope, with a million square metres of collecting area, be the next major radio telescope to be built. The project has acquired the appellation, the Square Kilometre Array (SKA). The driving ambition for this new facility, indeed of the next generation of astronomers, is no less than to chart a complete history of time. This imperative demands a frequency coverage from several meters to about one centimetre, with a sensitivity of 100 times that of the Very Large Array. Using new technologies, the SKA will become the premier imaging instrument of its generation in any wavelength region. With a spatial resolution better than the Hubble Telescope, a field of view larger than the full Moon, and the ability to simultaneously image a wide range of red shift (as many objects at high redshift in one long integration as the whole Las Campanas redshift survey of galaxies!), the SKA will be a discovery instrument to rival the NGST. Simply put, the goals of SKA are to

probe the structure and kinematics of the Universe before the dawn of galaxies • to understand the physics of the early Universe and how galaxies arose. To chart the formation and evolution of galaxies from the epoch of formation. • To Measure the evolution of the properties of galaxies, including dark matter haloes, trace the star formation history of the Universe, and explore the origin of cosmic magnetic ﬁelds and their role in galaxy evolution. To understand key astrophysical processes relating to the process of star forma- • tion and the physical and chemical evolution of galaxies by studies of the local Universe. To trace the physical mechanisms that give rise to planetary systems, to un- • derstand the evolution of our own solar system, and to engage in deﬁnitive experiments to answer the question, “Are we alone?”. To detect long-period gravitational waves, conduct exhastive tests of general • relativity, and explore the properties of nuclear matter within neutron stars.

9 10 Executive Summary

The SKA can be expected both to drive the agenda for, and complement in essential ways, much of the program of the Next Generation Space Telescope (NGST), which would give a similarly large increase in sensitivity over the Hubble Space Tele- scope in the near infrared waveband, where thermal starlight from the earliest objects is most readily visible.

An International Project The time frame in which the Square Kilometre Array is needed to complement other next generation instruments will be in the years around 2010. In September 1993 the International Union of Radio Science (URSI) established the Large Telescope Working Group to begin a worldwide effort to development the scientific goals and technical specifications for a next generation radio observatory. Subsequent meetings of the working group have provided a forum for discussing the technical research required and for mobilizing a broad scientific community to cooperate in achieving this common goal. The project is rapidly gathering momentum in the international arena. In 1997, eight institutions from six countries (Australia, Canada, China, India, the Netherlands. and the U.S.A) signed a “Memorandum of Agreement to Cooperate in a Technology Study Program Leading to a Future Very Large Radio Telescope”. Interest has also been expressed by groups of scientists in several other countries. The Square Kilometre Array international planning meetings are occurring at a rapid pace. The workshop in Sydney, Australia in December 1997 was followed rapidly by meetings in Calgary, Canada (July 1998) and Green Bank, U.S.A. (October 1998). In April 1999 two more meetings on the scientific goals and technical challenges will take place in Amsterdam and Dwingeloo, the Netherlands. The SKA will be an interferometric array with maximum baseline of several hundred to a thousand kilometres. As currently envisaged, each station of the array will be a radio telescope with aperture of about 200 meters - a factor of two larger in diameter than the largest existing fully steerable radio telescope. Wide field of view will be achieved, despite large antenna diameters, by constructing stations with arrays of of smaller sub-elements or by use of phased-array, multi-beam receivers on large reflectors. To realize a Square Kilometre Array at reasonable cost, a new means must be developed to construct very large aperture radio telescopes at a small fraction of the cost of conventional technology. Research and development activities are underway at several international centres. Solutions under study include adaptive array technology, “smart” antennas, large arrays of low-cost parabolic antennas, and novel concepts for very large, single-aperture antennas. Plans are well developed for construction of different prototype telescopes within the next several years. Convergence on a technological concept(s) for the SKA is expected by around the middle of the next decade. For historical reasons, radio astronomy has no international vehicle (such as an ESO or CERN) to promote co-ordination of activities among countries and to organize large multi-national facilities. As an ad hoc forum, the Organization for Economic Executive Summary 11

Co-operation and Development (OECD) has established a Working Group on Radio Astronomy in order to inventory plans for future large facilities around the world, and to consider whether specific actions by governments are necessary to make future large telescopes possible. Astronomers and funding agency officials from fifteen countries have attended Working Group meetings, and have in their Final Report to governments identified the mm arrays and SKA as the main international mega- projects being discussed for development during the coming 10-20 years in most of the participating countries. The Working Group identified one area that requires high level government in- volvement even now if the planned large investments are to yield maximum scientific returns. That is, SKA will not only have 100x the sensitivity of current instruments to celestial sources but also to man-made interference. And to survey in redshift will require access to large portions of the radio spectrum. The Working Group concluded that these crucial needs are probably unachievable within the science system alone. It therefore falls to governments at a high level to initiate steps to evolve the current regulatory regime such that by 2010 it will be possible, somewhere on Earth, to observe with the required sensitivity and bandwidths. The OECD delegations have accepted this challenge, and are recommending to the tri-ennial summit meeting of OECD science ministers in June 1999 that a special task force be formed to formu- late and carry out appropriate measures to ensure that the desired observations are possible when SKA gets built. Initial thoughts are that this will require establishing one or more internationally recognized interference free zones in unpopulated areas of the world.

Chapter 1

Introduction

1.1 A Next Generation Radio Observatory

The generation of radio telescopes constructed in the decades following the Second World War revolutionized our understanding of the Cosmos. The discovery of interstellar atomic hydrogen, the discovery of the 3 K microwave radiation, the discovery of pulsars, and the discovery of radio galaxies and quasars astonished the astronomical community and electrified the public. Other major discoveries included the discovery of interstellar molecules and molecular clouds, which opened the field of star and solar system formation. The Hulse-Taylor binary pulsar led to the discovery of gravitational radiation. The discovery of powerful interstellar masers was compelling in its own right, and led to the most accurate way of calibrating the extragalactic distance scale, of fundamental importance to cosmology. Each of these discoveries followed on the heels of new technological developments in antennas, receivers, timing circuits, and interferometry, among others. New developments in all fields of astronomy have brought us to the brink of understanding the origin and evolution of the Universe as a whole, but to make new progress at radio wavelengths a revolutionary new instrument is needed, one with an effective collecting area more than 30 times greater than the largest telescope ever built. Such a telescope will reveal the dawn of galaxy formation, as well as a plethora of other new discoveries in all fields of astronomy. Technological developments make it possible for such a telescope to be built within the next decade, and the international radio astronomical community is proposing that such a telescope, with a square kilometre of collecting area, be the next major radio telescope to be built.The project has acquired the appellation, the Square Kilometre Array. Additional new urgency is provided by recent planning in the space research community for a Next Generation Space Telescope (NGST), which would give a similarly large increase in sensitivity over the Hubble Space Telescope in the near infrared waveband, where thermal starlight from the earliest objects is most readily visible. Our investigations of the local Universe have shown us that proper understanding of the discoveries to be made by a NGST and other advanced new instruments cannot

13 14 CHAPTER 1. INTRODUCTION

be developed without a radio telescope operating at similarly improved sensitivity. The time frame during which a new radio facility is needed to complement other planned instruments will be in the years around 2010. In September 1993 the In- ternational Union of Radio Science (URSI) established the Large Telescope Working Group to begin a world-wide effort to develop the scientific goals and technical specifications for this next generation radio observatory. The seven subsequent meetings of this working group have provided a forum for discussing the technical research required and for mobilizing a broad scientific community to cooperate in achieving this common goal. The current document has grown from these discussions and represents the first international effort to document some of the many science goals that will be addressed by this facility. As will be evident even to the most casual reader, the range of important applications is as broad as astronomy itself, extending from the orbits of near Earth asteroids to the first structures in the infant universe. These topics organize themselves naturally into the Chapters of the document, dealing with the formation and evolution of galaxies, of stars, the solar system and finally of life itself. Before considering these applications we begin by outlining the technical capabilities of the envisioned facility; the spectral coverage, sensitivity, resolution and field of view which determine its utility. One particularly noteworthy aspect of the facility deserves special mention at the outset. This will be the world’s premier astronomical imaging instrument. No other existing or planned instrument in any wavelength regime can provide simultaneously:a spatial resolution better than the Hubble Space Telescope (< 0.1 arcsec), a field of view significantly larger than the full moon( 1 square degree), a spectral ∼ coverage of more than 50% (ν/∆ν < 2) a spectral resolution sufficient for kinematic studies (ν/dν > 104) and all at a sensitivity which is about 100 times that which is now available. By comparison, the largest optical integral field units which are now being considered for construction on 8-m telescopes would only provide a field of view of perhaps 10 arcseconds on a side with 10% spectral coverage at a comparable resolution, while the next generation of millimetre arrays is envisaged to provide a field of about 40 arcseconds with perhaps 10% spectral coverage (10 GHz at ν = 100 GHz). The field of view of the SKA is illustrated in comparison to the HST deep field image in Fig. 1.1.The implications of this wide-field spectral imaging capability are truly profound. The sky will become accessible in a way that is difficult for us now to fully appreciate. In many scientific areas we can already make confident predictions of the very substantial impact that this quantum leap in capability will enable. At the same time we recognize that the most profound discoveries will by their nature be impossible to predict. History has taught us humility in this respect. Every leap in sensitivity has revealed entirely unexpected object classes and physical phenomena.The fundamental problems of today are often not resolved, but instead dissolve into a new set of fundamental questions for tomorrow. But therein lies the beauty of science; the 1.2. THE SQUARE KILOMETRE ARRAY CONCEPT 15

Figure 1.1: Comparison of the fields of view of the SKA with size of the Hubble Deep Field and the field of view of the MMA. By combining interferometry and phased- array receiver technology, the SKA will image a field of view of 1◦ at λ21 cm with angular resolution of 0.1′′. potential for discovery has no bounds.

1.2 The Square Kilometre Array Concept

Extensive discussion of the envisioned science drivers and of the evolving technical possibilities has led to a concept for theSquare Kilometre Array and a set of design goals. The SKA will be an interferometric array of individual antenna stations, synthesizing an aperture with diameter of approximately 1000km. An example conﬁg- uration for the array is illustrated in Fig. 1.2.The 106 square metres of collecting area is distributed over30 interferometric stations. Each station is a 200 metre diameter telescope. Approximately 80% of the array is contained within a centrally-condensed inner array to provide ultra-high brightness sensitivity at arc-second scale resolution for studies of the faint spectral line signatures of structures in the early Universe. The outrigger stations provide a ten-fold increase in angular resolution to allow high resolution imaging of faint emission from the interstellar media of distant galaxies. At an SKA technical workshop in Sydney, Australia in December 1997,the goals 16 CHAPTER 1. INTRODUCTION

Figure 1.2: An example conﬁguration for the Square Kilometre Array. The 106 m2 of collecting area is spread over a synthetic aperture 1000 km in diameter. Approximately 80% of the collecting area is contained withi∼n a centrally condensed inner region yielding extremely high surface brightness sensitivity at arcsecond-scale resolution. The outrigger antennas provide a higher angular resolution mode. There are 30 individual antenna stations – each a radio telescope with 200 metre diameter. 1.2. THE SQUARE KILOMETRE ARRAY CONCEPT 17

Parameter Design Goal A /T 2 104 m2/K eff sys × Total Frequency Range 0.03 – 20 GHz Imaging Field of View 1 square deg. @ 1.4 GHz Number of Instantaneous Pencil Beams 100 Maximum Primary Beam Separation low frequency 100 deg. high frequency 1 deg. @ 1.4 GHz Number of Spatial Pixels 108 Angular Resolution 0.1 arcsec @ 1.4 GHz Surface Brightness Sensitivity 1 K @ 0.1 arcsec (continuum) Instantaneous Bandwidth 0.5 + ν/5 GHz Number of Spectral Channels 104 Number of Simultaneous Frequency Bands 2 Clean Beam Dynamic Range 106 @ 1.4 GHz Polarization Purity 40 dB − Table 1.1: SKA Design Goals for the basic system parameters were established. These are summarized in Ta- ble 1.1. The array will blend basic interferometric techniques with multi-element, phase dreceiving systems. Consequently, some of the terminology used,particularly that related to the beam-forming hierarchy,is non-intuitive. For the sake of clarity a short technical explanation for all of the listed parameters is provided below.

A /T : The eﬀective collecting area divided by the system temperature. • eff sys This may be a function of frequency.

Total Frequency Range: The total frequency tuning range of the instrument. • This may be divided into sub-ranges with diﬀerent antenna technologies, and it is not necessarily contiguous.

Imaging Field-of-View: The instantaneous, contiguous solid-angle area of the • sky that can be imaged, given a suﬃciently capable correlator. This area will be a function of frequency.

Number of Instantaneous Pencil Beams: The number of “phased array” pencil • beams that can be placed simultaneously within the Imaging Field-of-View for point source observations such as pulsars, stars (including SETI), and VLBI.

Maximum Primary Beam Separation: This speciﬁcation assumes that the facil- • ity will have at least two levels of beam forming. Signals from small antennas (dipoles, small dishes, etc.) are combined to form an array element primary beam, and signals from array elements can be combined in a correlator to make 18 CHAPTER 1. INTRODUCTION

a map within the primary beam or combined directly to form one or many pencil beams within the primary beam. More than one primary beam could be formed within the pattern of the small antennas. The Maximum Primary Beam Separation speciﬁes how far apart these primary beams can be formed simultaneously.

Number of Spatial Pixels: The number of spatial resolution elements in a map • synthesized within the Imaging Field-of-View.

Angular Resolution: The maximum angular resolution of the array as deter- • mined by its largest linear extent (longest baseline).

Surface Brightness Sensitivity: The minimum detectable (5σ) continuum sur- • face brightness for a speciﬁed resolution, e.g., 1K @ 0.1 arcsec. This may be a function of frequency.

Instantaneous Bandwidth: The widest contiguous frequency range that may be • observed simultaneously given enough correlator or other processing capability. Typically this means the widest selectable IF ﬁlter bandwidth before the digitizer.

Number of Spectral Channels: The number of independent frequency samples • from the array after all signal processing.

Number of Simultaneous Frequency Bands: The number of widely spaced fre- • quency ranges that may be observed simultaneously. For example, a stellar ﬂare study might want to observe at 1.4 and 5.0 GHz at the same time, each with instantaneous bandwidths of 0.3 GHz. Clean Beam Dynamic Range: The best intensity dynamic range that may be • obtained in a fully processed synthesized map, as limited by unknown errors in the array or its environment.

Polarization Purity: The error in Q, U, and V Stokes parameters as a fraction • of I for a strong radio source after all data processing, as limited by unknown errors in the array or its environment.

1.3 Instrumental Sensitivity

The system sensitivities which follow from the design goals for both broad-band continuum and moderate resolution spectral line applications (ν/dν = 104, corresponding to 30 km s−1) are summarized in Table 1.2 for a single polarization channel of an observation with an assumed 8 hour duration. These sensitivities are contrasted with those of many existing and planned facilities in Figs. 1.3 and 1.4, together with simulated spectra of the spiral galaxy M101 as it would appear at red-shifts of 0.5, 2, 8 and 32 under the assumption of no 1.3. INSTRUMENTAL SENSITIVITY 19

Freq. Total Band. Aeff /Tsys Cont. rms Line rms (MHz) (MHz) (m2/K) (nanoJy) (µJy) 40 20 500 5140 364 80 40 3 103 610 43. 160 80 2×104 64 4.6 × 320 160 2 104 45 3.2 640 320 2×104 32 2.3 1280 640 2×104 23 1.6 × 2560 1280 2 104 18 1.1 5120 1500 2×104 15 0.80 × 10240 2500 2 104 11 0.57 20480 4500 1×104 17 0.80 ×

Table 1.2: Instrumental Sensitivity per Polarisation in 8 hours

Figure 1.3: Simulated spectra of the spiral galaxy M101 are shown for frequencies between about 108 and 1014 Hz after being red-shifted to z = 0.5, 2, 8 and 32, under the assumption of no spectral evolution. Instrumental sensitivities (1σ) of existing and planned instruments are overlaid for spectral line observations. Spectral line IDs for some of the major emission lines are indicated at z = 0. 20 CHAPTER 1. INTRODUCTION

Figure 1.4: Simulated spectra of the spiral galaxy M101 as in Fig. 1.3 but with the instrumental sensitivity of broad-band continuum observations overlaid. 1.4. CONTRIBUTORS 21

spectral evolution. The striking result illustrated in Fig. 1.3 is that even a normal 10 galaxy like M101 with an atomic gas mass,MHI = 2 10 M⊙ and molecular gas 9 × mass,M 2 = 3 10 M , could be detected efficiently(recall the large field of view H × ⊙ and large instantaneous spectral coverage) via its emission lines at arbitrary red-shift. The HI line can be tracked to red-shifts beyond 4, while the CO lines also become accessible at red-shifts of 4 and larger. Similarly, the continuum emission of normal galaxies can be detected out to ba- sically arbitrary distances (as shown in Fig. 1.4).The non-thermal continuum will be seen at low frequencies to red-shifts beyond 8, while the thermal emission from ionized gas dominates the emission at high-frequencies.The bright far-infrared dust continuum peak becomes red-shifted into the high frequency band of the facility at even larger red-shifts. Whatever the epoch was of the first star and galaxy formation in the universe, we will be able to study it in detail with this facility.This fact serves to illustrate the magnitude of the scientific revolution before us. Similar impacts are expected in all the current areas of research, but as mentioned previously, the most profound discoveries will likely be made in directions for which we haven’t yet even formulated the right questions.

1.4 Contributors

A large number of people have contributed to the eﬀorts of the URSI Large Tele- scope Working Group. This document was prepared partially based on presenta- tions at the SKA Science Workshop held in Calgary, Canada in July 1998, and partially on other contributions. Sections of the text, relevant background information, and illustrations were made available by: T. van Albada, K. Anantharamaiah, W. Baan, M. Bailes, N. Bartel, R. Beck, J. Bell, L. Blitz, R. Braun, F. Briggs, G. de Bruyn, B. Burke, H. Butcher, L. Cram, P. Dewdney, J. Dickel, S. Dougherty, J. Dreher, N. Duric, R. Ekers, W. Erickson, D. Frail, M. Franx, L. Gurvits, A. Green, M. Guedel, M. van Haarlem, H. Hoekstra, A. Hopkins, J.M. van der Hulst, C. Jack- son, D. Jauncey, N. Kassim, K. Kellermann, L. Knee, L. Koopmans, K. Kuijken, T. Kuiper, T. Landecker, J. Lim, R. Manchester, A. Meiksin, M. Montes, R. Norris, T. O’Brien, D. O’Neill, S. Ostro, I. de Pater, A. Pedlar, G. Pettingill, R. Redman, L. Rodriguez, R. Schilizzi, E. Seaquist, G. Squires, L. Stavelely-Smith, R. Strom, J. Tarter, R. Taylor, B. Thomas, J. Ulvestad, B. Wallace, K. Weiler, S. Van Dyk, R. van de Weygaert, R. Wielebinski and P. Wilkinson. The document has been edited by A. R. Taylor and R. Braun.

Chapter 2

Formation and Evolution of Galaxies

2.1 The Dawn of Galaxies: Searching for the Epoch of First Light

Prior to the epoch of full reionization, the intergalactic medium and gravitationally collapsed systems will be detectable in 21-cm radiation. Physical mechanisms that would produce a 21-cm signature are Lyα coupling of the hydrogen spin temperature to the kinetic temperature of the gas resulting from the radiation by an early generation of stars, preheating by soft x-rays from collapsing dark matter halos, and preheating by ambient Lyα photons. A patchwork of either 21-cm emission, or absorption against the Cosmic Microwave Background, will result. The Square Kilometre Array offers the prospect of measuring this signature, and so detecting the transitional epoch from a dark universe to one with light. The development of structure in the Universe was well advanced at early times. Quasars have been detected nearly to a redshift of z = 5, and the most distant galaxies to even greater redshifts (Dey et al., 1998). The spectra of high redshift QSOs have additionally shown that the Intergalactic Medium (IGM) itself had undergone an extensive development of nonlinear structures at early times as well, as revealed by the Lyα forest. Still unknown, however, is the nature of energetic processes at these early times. While numerical simulations have shown that the IGM is expected to fragment into structures at early times in Cold Dark Matter (CDM) dominated cosmologies (Zhang et al., 1998), and even into early galaxies (Governato et al., 1997), the simulations are much less able to predict the efficiency with which gravitationally collapsed objects will emit radiation. Although QSO’s may account for the photoionizing UV background at high redshifts (Meiksin & Madau 1993; Haardt & Madau 1996), it is less clear that they were responsible for the original reionization of the IGM. Similarly, although z > 5 galaxies have been detected, the epoch during which the first generation of stars formed is still poorly constrained (Madau et al., 1998; Hughes et al., 1998). Although IR observations will permit even higher redshift

23 24 CHAPTER 2. FORMATION AND EVOLUTION OF GALAXIES galaxies and QSOs to be observed, detections become increasingly difficult because of the diminution in surface brightness due to cosmological expansion. This difficulty calls for alternative means to be found for discovering the nature and period of the first major generation of energy-producing sources, the epoch of First Light. Probing this era is critical for answers to the following questions:

When did the ﬁrst stars form? At some redshift z > 10 in a CDM Universe, • objects will start to form that are suﬃciently massive to lead to the production of primordial star clusters. If the conditions under which such objects can collapse and cool occur everywhere, then they are likely to be distributed homo- geneously, and the transition from a neutral to fully ionized IGM can be very rapid (when percolation of the regions around the ionizing sources occurs).

What are the first energy sources? In the event an inadequate fraction of • photoionizing radiation can escape the regions of early star formation to fully ionize the IGM, reionization is likely to be a much more inhomogeneous process dominated by QSOs. Very luminous QSOs will generate giant expanding HII regions, that are only limited in size by the lifetime of the QSO. The clustering properties of the QSO’s determine the epoch of breakthrough of the expanding HII bubbles; for a more clustered distribution full reionization will be delayed in the low-density regions and will therefore occur later than for a homogeneous distribution of sources. However, since accretion is a much more efficient mechanism than nuclear fusion (per unit mass used in the process), considerably fewer QSO’s would be required to fully ionize the IGM, than would be the case for a population of primordial stars. Nonetheless, the first stars may still reveal their presence due to the effect of their continuum radiation on the spin temperature of the neutral IGM.

What is the size distribution of mass perturbations? The free electrons re- • quired to reionize the Universe decrease the amplitude of fluctuations in the microwave background radiation, but only on scales below the size of the horizon at reionization ( 10 degrees). Different CDM models that do not violate the constraints set by∼ COBE measurements, and that produce approximately the correct number of galaxies at low redshift, can differ appreciably in the amplitude of fluctuations on small scales. The details of the process of reionization 6 9 lie fully in the collapse of objects in the 10 10 M⊙ range. Reionization probes precisely this mass range, and therefore forms− an excellent way of discriminating between the competing cosmological models.

How do collapsing objects evolve? Reionization itself plays an important role • in the collapse and cooling of subsequent generations of objects. For example, the ambient radiation after reionization may delay the collapse of galaxies and inhibit star formation. 2.1. THE DAWN OF GALAXIES 25

The Square Kilometre Array could reveal the ﬁrst epochs of energy generation. The means by which the sources are revealed is through their impact on the surrounding neutral IGM and the resulting emission or absorption of 21-cm radiation.

2.1.1 21-cm Emission and Absorption Mechanisms The Spin Temperature

The emission or absorption of 21-cm radiation from a neutral IGM is governed by the spin temperature TS of the hydrogen, deﬁned by

n1 = 3 exp[ T∗/TS], (2.1) n0 −

where n and n are the singlet and triplet n = 1 hyperfine levels, T hν /k = 0 1 ∗ ≡ 10 B 0.07 K, where ν10 is the frequency of the 21-cm transition, h is Planck’s constant, and kB is Boltzmann’s constant. In the presence of only the Cosmic Microwave Back- ground (CMB) radiation, the spin temperature will be the same as the temperature of the CMB, and no emission or absorption relative to the CMB will be detectable. A mechanism is required that decouples the two temperatures. This may be achieved by coupling the spin temperature to the kinetic temperature of the gas itself. Two mechanisms are available, collisions between hydrogen atoms (Purcell & Field 1956) and scattering by Lyα photons (Wouthuysen 1952; Field 1958). The collision-induced coupling between the spin and kinetic temperatures is dominated by the spin-exchange process between the colliding hydrogen atoms. The rate, however, is too small for realistic IGM densities at the redshifts of interest, although collisions may be important in dense regions, δρ/ρ > 30[(1 + z)/10]−2 (Madau, Meiksin, & Rees 1997). Instead the dominant∼ coupling mechanism is likely to be Lyα scattering through the Wouthuysen-Field effect. This process mixes the hyperfine levels of neutral hydrogen in its ground state via an intermediate transition to the 2p state. An atom initially in the n = 1 singlet state may absorb a Lyα photon that puts it in an n =2 state, allowing it to return to the triplet n = 1 state by a spontaneous decay. At this point, the astute student of quantum mechanics will ask how is it possible for electric dipole radiation (Lyα photons) to induce a spin transition? The key is spin-orbit coupling: it’s the total angular momentum F = I + J that counts. (Here I is the proton spin and J is the total electron angular momentum, J = S + L.) There are four hyperfine states involved, the n = 1 singlet 0S1/2 and triplet 1S1/2 states (the notation is F LJ ), and the two triplet n = 2 states 1P1/2 and 1P3/2. The selection rule ∆F =0, 1 permits the transitions 0S1/2 1P1/2, 1P3/2 and 1P1/2, 1P3/2 1S1/2, and so effectively S S occurs via one→ of the n = 2 states. → 0 1/2 → 1 1/2 When the IGM is highly opaque to the scattering of Lyα photons, as it is when still neutral, the large number of scatterings of Lyα photons in an ambient radiation field will ensure a Boltzmann distribution for the photon energies near the Lyα frequency, with a temperature given by the kinetic temperature TK of the IGM (Field 1959). In 26 CHAPTER 2. FORMATION AND EVOLUTION OF GALAXIES

this case, the spin temperature of the neutral hydrogen becomes1

TCMB + yαTK TS = , (2.2) 1+ yα where TCMB = 2.73(1 + z) K is the temperature of the CMB (Mather et al., 1994), and P10 T∗ yα (2.3) ≡ A10 TK is the Lyα pumping efficiency. Here, A =2.9 10−15 s−1 is the spontaneous decay 10 × rate of the hyperfine transition of atomic hydrogen, P10 is the indirect de-excitation rate of the triplet via absorption of a Lyα photon to the n = 2 level, and T T was S ≫ ∗ assumed. In the absence of Lyα pumping the spin temperature goes to equilibrium with the 21-cm background radiation field on a timescale T /(T A ) 5 104 yr, ∗ CMB 10 ≈ × and neutral intergalactic hydrogen will produce neither an absorption nor emission signature. If yα is large, TS TK , signifying equilibrium with the matter. A consideration of the net transition→ rates between the various hyperfine n = 1 and n =2 levels above shows that the 1 0 transition rate via Lyα scattering is related to → the total rate Pα by P10 = 4Pα/27 (Field 1958). This relation and equation (2.2) are derived in the Appendix. In the limit T T , the fractional deviation in a K ≫ CMB steady state of the spin temperature from the temperature of the CMB is

T T T −1 S − CMB 1+ CMB . (2.4) TS ≈ " yαTK #

There exists then a critical value of P which, if greatly exceeded, would drive T α S → TK . This thermalization rate is (Madau et al., 1997)

27A10TCMB −12 −1 1+ z Pth (5.3 10 s ) . (2.5) ≡ 4T∗ ≈ × 7

21-cm Emission Eﬃciency To illustrate the basic principle of the proposed observations, consider a region of neutral material with spin temperature TS = TCMB, having angular size on the sky which is large compared to a beamwidth, and6 radial velocity extent due to the Hubble expansion which is larger than the bandwidth. Its intergalactic optical depth at 21(1 + z) cm along the line of sight,

3 3 2 3c h nHI(0)A10 1.5 −2.9 −1 TCMB ΩIGMh50 1/2 τ(z)= 3 2 (1 + z) 10 h50 (1 + z) , (2.6) 32πH0kBT∗ TS ≈ TS 0.05 ! 1In the presence of a radio source, the antenna temperature of the radio emission should be added to TCMB in equation 2.2. The radio emission may make an important contribution in the vicinity of a radio-loud quasar (Bahcall & Ekers 1969), and would itself permit the IGM to be detected in 21-cm radiation. 2.1. THE DAWN OF GALAXIES 27

will typically be much less than unity. The experiment envisaged consists of two measurements, separated in either angle or frequency, such that one measurement, the ﬁducial, detects no line feature, either because there is no H I or because T T , S ≈ CMB and the second at TS = TCMB. Since the brightness temperature through the IGM −τ 6 −τ is Tb = TCMBe + TS(1 e ), the diﬀerential antenna temperature observed at the Earth between this region− and the CMB will be

2 1/2 −1 −τ −1 ΩIGMh50 1+ z δT = (1+ z) (TS TCMB)(1 e ) (0.011 K)h50 η, (2.7) − − ≈ 0.05 ! 9 where the 21-cm radiation eﬃciency is deﬁned as

TS TCMB η xHI − . (2.8) ≡ TS Here x refers to the neutral fraction of the hydrogen in the region for which T = HI S 6 TCMB. As long as TS is much larger than TCMB (hence if there has been signiﬁcant preheating of the intergalactic gas), η xHI, and the IGM can be observed in emission at a level which is independent of the exact→ value of T . By contrast, when T T S CMB ≫ S (negligible preheating), the diﬀerential antenna temperature appears, in absorption, a factor TCMB/TS larger than in emission, and it becomes relatively easier to detect intergalactic∼ neutral hydrogen (Scott & Rees 1990).

2.1.2 Preheating the IGM

The role of the spin temperature is manifest in eq. (2.8): when TS < TCMB the IGM absorbs 21-cm radiation from the CMB, while for TS > TCMB the IGM emits 21-cm radiation in excess of the CMB. In the absence of decoupling mechanisms, TS = TCMB. The presence of Lyα photons with suﬃcient intensity will thus enable the IGM to be “seen.” The adiabatic expansion of the Universe will generally bring the kinetic temperature of the IGM well below the temperature of the CMB. Coupling TS to TK will permit the IGM to be detectable in absorption. If there are sources of radiation that heat the IGM, however, it may be possible instead to detect the IGM in emission. Possible heating sources are soft x-rays from an early generation of QSOs or thermal bremsstrahlung emission produced by the ionized gas in the collapsed halos of young galaxies. In CDM-dominated cosmologies, the latter may be in suﬃcient number to heat the IGM above the CMB temperature by z 7 (Madau et al., 1997). While photons just shortward of the photoelectric edge≈ are absorbed at the ionization front generated by a QSO source, photons of much shorter wavelength will be able to propagate to much greater distances. Most of the photoelectric heating of the IGM by a QSO is accomplished by soft x-rays. The time required for the radiation at the light front to heat the intergalactic gas to a temperature above that of the CMB is typically 10% of the Hubble time. The HII region produced by a QSO will therefore be preceded by a warming front. Note that, as the X-ray-heated bubbles around QSOs will survive as fossils even after the quasar has died, several generations Ng of 28 CHAPTER 2. FORMATION AND EVOLUTION OF GALAXIES

quasars may actually be responsible for preheating the entire IGM. For a typical QSO age of t 3 108N −1/3 yr, the required QSO comoving space density to heat the Q ≈ × g entire IGM to a temperature above that of the CMB by z 6 is 10−10 Mpc−3N −1. ≈ ∼ g By comparison, the comoving space density of bright QSOs at z = 4 is 100Ng times larger (Warren, Hewett, & Osmer 1994). If all bright galaxies undergo∼ a quasar phase, QSOs must have a very short lifetime, and N 100. Soft X-rays from a few g ∼ bright QSO sources could then prevent collapsing structures, such as protoclusters while still in the linear regime, from being detected in 21-cm absorption against the CMB. An additional heating source is the Lyα photon scattering itself. The average relative change in a Lyα photon’s energy E after having been scattered by a hydrogen atom at rest is ∆E hνα −8 = 2 10 , (2.9) h E i −mHc ≈ − where mH is the mass of the hydrogen atom. (It should be noted that this is an approximation valid only for hνα kTK. In the opposite limit, energy will ﬂow from the atoms to the photons.) Through≫ recoil, energy is transferred from photons to atoms at a rate ∆E E˙ = hν P . (2.10) α −h E i α α

where Pα is the Lyα scattering rate per H atom. In the case of excitation at the thermalization rate Pth, equation (2.10) becomes

2 ˙ 27 (hνα) A10TCBR −1 1+ z Eth = 2 (220 KGyr ) , (2.11) 4 mHc T∗ ≈ 7 (Madau et al., 1997). The characteristic timescale for heating the medium above the CMB temperature via Lyα resonant scattering at this rate is

2 2 mHc ν10 −1 8 ∆theat = 2 A10 10 yr, (2.12) 9 hνα ≈ about 20% of the Hubble time at z 8. The result is a ﬁnite interval of time during which Lyα photons couple the spin≈ temperature to the kinetic temperature of the IGM before heating the IGM above the CMB temperature. If Lyα sources turned on at redshifts zα < 10, this interval would present a window in redshift space near z 8 that would enable∼ a large fraction of intergalactic gas to be observable at 160 MHz≈ ∼ in absorption against the CMB, and so isolate the epoch of First Light.

2.1.3 Scenarios: SKA Imaging of Cosmological H I Two cosmological models are considered, a tilted Cold Dark Matter model (tCDM) −1 −1 with Ω0 = 1, H0 = 50 kms Mpc , and σ8 = 0.55, designed to match both CMB measurements on large scales and the constraint on amplitude imposed by galaxy cluster abundances on small, and a ﬂat open CDM model (OCDM) with Ω0 = 0.3, 2.1. THE DAWN OF GALAXIES 29

−1 −1 ΩΛ = 0.7, H0 = 70 kms Mpc , and σ8 = 1.1, which similarly matches both con- 2 straints. The baryon density in both models is assumed to be ΩBh = 0.024, where −1 −1 h = H0/100 kms Mpc . We consider a scenario in which sources of Lyα photons are in sufficient abundance throughout the universe to couple the spin temperature to the kinetic temperature of the IGM everywhere. We further suppose the IGM has been preheated to a temperature well above that of the CMB, either by the same Lyα photons responsible for the coupling or by soft x-ray sources. In this case, the IGM will emit at a rate independent of TS (the hyperfine levels will be occupied according to their statistical weights, n1/n0 = 3). Because of structure in the IGM, the emission will not be uniform. In Fig. 2.1, we show the range of density fluctuations that would be detectable in a single beam by a Square Kilometre Array (SKA) interferometer at 160 MHz (z = 8), as a function of beam size and frequency band width, and for several assumed integration times. The detection thresholds are scaled according to a continuum rms in a 80 MHz band at 160 MHz of 64 nJy over an 8 hour integration time (Table 1.2). The experiment is based on taking differences between beams. The dashed lines show curves of constant rms antenna temperature fluctuation within the IGM, 2 1/2 (δTa) = σρT¯a, where σρ is the rms relative density fluctuation of the IGM for a hvolumei corresponding to a given bandwidth (∆ν/ν = ∆z/[1 + z]) and angular size ∆θ. Here, T¯a is the mean antenna temperature from the IGM at redshifted 21-cm. At a fixed bandwidth, the antenna temperature fluctuation increases with decreasing angular scale because σρ increases with decreasing linear scale. Because the detected flux is proportional to the solid angle of the beam, the detected signal decreases with decreasing angular scale until it falls below the detection threshold, indicated by the solid lines. The size of the detectable IGM fluctuations differs greatly for the two models. This is because the growth of density fluctuations ceases early on in an open universe, so that the fluctuations on a given angular scale are much larger in the OCDM model than in tCDM at high redshift. In a second scenario, the spin temperature is again coupled to the IGM kinetic temperature everywhere, but the IGM has not had time to heat above the CMB temperature. We then consider the emission signature resulting from a QSO soon after it turns on as the medium surrounding it is heated by soft x-rays from the QSO. The experiment in this case is done by differencing beams pointed in regions around the QSO with one pointed through the QSO HII region, where both emission and absorption are absent. An image of the resulting emission for the tCDM model is shown in Fig. 2.2. As the warming front produced by the QSO expands, a growing amount of the surrounding IGM is revealed. Note that, although the QSO was placed in the corner of the simulation volume, the figure can equally be viewed as the emission due to heating by a beam of soft x-rays from the QSO with an opening angle of 90◦. Thus imaging the gas surrounding a QSO in 21-cm emission would provide a direct means of measuring the opening angle of QSO emission. It should also be noted that a region of a given fixed density fluctuation will not always yield the same fluctuation in 21-cm emission. This is because of the 30 CHAPTER 2. FORMATION AND EVOLUTION OF GALAXIES

Figure 2.1: Detectability of fluctuations in 21-cm brightness temperature (10−3 < 2 1/2 −2 (δTa) < 10 ), from the IGM by the SKA as a function of beam width and hbandwidth.i The dashed lines show the size of the rms fluctuations expected in a volume defined by the angle and frequency widths. The solid lines show the threshold below which the fluctuations are undetectable for integration times of 10, 100, and 500 hours.

Figure 2.2: H I emission from the region surrounding a QSO source (lower right) at z=8.5 revealed once the region is heated above the temperature of the Cosmic Microwave Background by soft x-rays from the QSO. The image is 17′ on a side. The model was computed using Hydra (Couchman, Thomas, & Pearce 1995) 2.2. LARGE SCALE STRUCTURE AND GALAXY EVOLUTION 31

dependence of the spin temperature on the temperature of the IGM. Only when the IGM temperature much exceeds that of the CMB will the 21-cm emission be independent of the IGM temperature, according to equation (2.8). In general, the fluctuations in brightness temperature will depend on both the density fluctuations of the IGM and the temperature fluctuations, which in turn depend on the ages and distribution of the sources. A knowledge of the cosmological density fluctuation spectrum, as may be measured by future CMB missions like MAP and Planck, will then enable the statistical distribution of the sources that heat the IGM, whether QSOs as here or Lyα photons from early stars, to be established using measurements of the fluctuations in the 21-cm sky at high redshift. The calculations in this section were done by Avery Meiksin in collaboration with Paolo Tozzi and Piero Madau.

2.2 Large Scale Structure and Galaxy Evolution

The fascinating observations obtained with the HST, in particular the ongoing analysis of the few thousand galaxies in the Hubble Deep Field shows that there is already significant evolution detectable in the comoving star formation rate density by looking back to a redshift of 1. (Madau, 1998). Already looking back between z = 0.5 and z = 1 (3 4 h−1Gyr) there should be a noticeable increase of a factor 2 3 in SFR density.− The SFR density appears to peak around z = 1.5 with the most− vigorous evolution between z = 1 and z = 3. This analysis is largely based on optical photometry and spectroscopy of galaxies in the HDF. Little is known, however, about the evolution of the H I in galaxies out to redshifts of 1 and beyond, because present day instruments lack the sensitivity and resolution to directly measure the H I in galaxies at these redshifts. Damped Lyα studies (Lanzetta et al. 1995) indicate that the comoving H I mass density is roughly 5 10 the present beyond z = 1 and out to z = 3. This is also the period during− which× metal-rich gaseous halos appear, confirming that this is an era of strong evolution, where it is imperative to have good insight into the evolution of the H I content of galaxies. The SKA can measure the H I in galaxies back to redshifts of z 3 and will revolutionize this area of research. The great potential of SKA is that,≈ unlike optical surveys, it will be able to find galaxies independent of effects of extinction and color using the H I , with the additional advantage that once an object has been found the H I line provides an accurate redshift at the same time. To fully exploit the scientific potential one needs both the H I and optical information. The latter will be coming forward from planned surveys such as the Sloan Digital Sky Survey (SDSS, Gunn and Weinberg 1995) and projects with the Hubble Space Telescope, such as the Hubble Deep Field. The SKA will probe a piece of parameter space, i.e. the neutral gas content, which is absolutely required for understanding galaxy evolution and can only be probed at radio wavelengths. To fully demonstrate the potential of a large radio telescope facility such as the SKA three specific examples of studies involving measurement of the H I 21-cm line 32 CHAPTER 2. FORMATION AND EVOLUTION OF GALAXIES

in distant galaxies will be described here: (i) a deep “pencil beam” survey of an area of one square degree out to redshift of z > 4; (ii) a shallower survey of an area of 1000 square degrees to a limiting redshift of about 1; and (iii) a search for low column density intergalactic H I emission to try to map out the structure of the H I responsible for the Lyα forest lines. Before discussing these it should be emphasised that observing the 21-cm H I line in emission has special requirements for the geometry of SKA. H I studies require high surface brightness sensitivity rather than sub-arcsecond resolution since the brightness temperatures of the emission are at most several tens of Kelvin. Thus there always is a delicate tradeoﬀ between resolution and surface brightness sensitivity. For example: to reach a brightness temperature limit of 0.5 Kelvin in 12 hours of integration time (roughly corresponding to a column density limit of a few 1019 atoms cm−2) one requires a resolution of 1–3 arcsec for redshifts below z = 2. This implies that one needs to have most of the collecting area in baselines below 100 km. In contrast to this: continuum emission usually has much higher brightness temperatures and can therefore be observed at much higher resolution.

2.2.1 A Deep SKA H I Pencil Beam Survey

Let us consider a 360 hour integration on a single ﬁeld of one square degree. For H I studies one requires high surface brightness sensitivity rather than sub-arcsecond resolution. A 1 km2 array with baselines up to 100 km will provide a resolution of 1” at 610 MHz (z =1.3) corresponding to a linear size of 4.4 kpc (for H0 = 100 km/s/Mpc and Ω=1.0). Such an instrument will be able to detect L∗ galaxies (which typically 9 −2 have H I masses of 3.5x10 h M⊙) out to redshifts of z = 3, i.e. beyond the redshift range where the universe shows considerable evolution. Using the H I mass function of Zwaan et al. (1997) shown in Fig. 2.3 and assuming no evolution of the H I properties of galaxies with redshift one can calculate how many galaxies one would expect to detect in a one square degree ﬁeld of view per redshift interval. Table 2.1 gives a brief summary.

Out to a redshift of 2 the resolution will be suﬃcient to resolve a fair fraction of the galaxies. This implies that one can obtain rotation curves, mass distributions, gas fractions for some 105 galaxies between now and 5 h−1Gyr ago. One would be in a unique position to trace the evolution of the ISM in galaxies over a substantial fraction of the age of the universe, from the era of strongest evolution and star formation activity until the present. In addition one would learn whether and how the evolution depends on the dark matter content and environment. Low surface brightness galaxies for example appear to be rather unevolved, happen to avoid the denser regions in the universe and probably have low density dark matter halos (de Blok and McGaugh 1997). This notion, however, is based on only a small number of well studied objects and clearly a survey like that described here is required to ﬁrmly establish such relationships. 2.2. LARGE SCALE STRUCTURE AND GALAXY EVOLUTION 33

Figure 2.3: Lower panel: The distribution of H I masses of the detected galaxies. The error bars are given by Poisson statistics. Upper panel: The thin line is the sensitivity of the survey. The measured H I mass function per decade is given by the points. The thick line is a Schechter luminosity function with the parameters given in the upper right corner. (from Zwaan et al. 1997)

Red-shift Look Back Time H I Mass Limit Number of Detections −1 −2 (h Gyr) (h M⊙) 0.5–1.0 3.0–4.2 1.3108 3.2 105 1.0–1.5 4.2–4.9 2.8108 2.5 105 1.5–2.0 4.9–5.3 5.0108 9.9 104 2.0–2.5 5.3–5.6 9.0108 7.0 104 2.5–3.0 5.6–5.7 1.4109 3.6 104 3.0–3.5 5.7–5.8 2.2109 2.3 104 3.5–4.0 5.8–5.9 3.4109 1.0 104 4.0–4.5 5.9–6.0 5.1109 0.8 104

Table 2.1: Detectable H I Masses for an SKA Deep Pencil Beam Survey 34 CHAPTER 2. FORMATION AND EVOLUTION OF GALAXIES

The great advantage of a “pencil beam” survey as described here is, of course, that the selection of objects in the field is entirely based on H I content and not on the associated stellar component. The selection is therefore independent of the effects of extinction, color and optical surface brightness. The selection does, of course, depend on H I content and H I surface brightness. The combination of deep, H I selected samples and deep, optically selected samples will be extremely powerful for studying galaxy evolution over a large range of redshift. In addition to the H I content the survey will also measure the continuum emission of the galaxies in the field. The continuum emission is known to correlate almost perfectly with the far IR emission of spiral and irregular galaxies (Helou 1991, Condon 1992,Lisenfeld et al. 1996). The FIR emission appears to be a good indicator of massive star formation rate so that one can use the continuum emission to probe the star formation rates of the detected galaxies, independent of the effects of extinction. This information can be used to link the star formation rates to the H I contents of galaxies as a function of redshift and environment. In addition it will provide an independent estimate of the evolution of the comoving star formation rate density to be compared with the optically determined functions (Madau 1998). With such a pencil beam survey it will also be possible to verify the characteristics of the Tully-Fisher relation over distances out to z = 2 using the flat part of the rotation curves and establish whether it can be used as a reliable tool for independent distance determinations. Tully-Fisher work will even be possible out to higher redshifts, since all one needs to measure is a redshift and an H I profile width. An L∗ galaxy (assuming no evolution in the gas fraction, which is quite unlikely) can be detected out to redshifts of 3 in a 360 hour integration.

2.2.2 Large scale structure studies from a shallow, wide area survey In addition to a “pencil beam” survey one can perform a shallow survey covering a large area of sky to a depth of z 1. In 12 months of observing time one could cover 1000 square degrees and be able∼ to detect L galaxies out to a redshift of z 2 ∗ ∼ (or about 75% of the age of the universe). Assuming the Zwaan et al. (1997) H I mass function one expects to detect 107 galaxies in a volume of 107 Mpc3. This is orders of magnitude more than in∼ optical surveys, such as the∼ Sloan Digital Sky Survey (Gunn and Weinberg 1995) and the AAO 2dF Survey (Lahav 1995, Cannon 1995). The properties of the SKA wide ﬁeld survey are compared to those of large optical surveys in Fig. 2.4. The great potential of such an H I survey is the possibility of studying the large scale structure in the universe to greater depth than possible at present. The large coverage in redshift, coupled to the large number of detectable objects makes it possible to trace the evolution of large scale structure with redshift out to at least z =1.3. Or in other words provide the tools to determine structures and density ﬂuctuations on scales between 10h−1 and 4800h−1 Mpc. The evolution of large scale structure 2.2. LARGE SCALE STRUCTURE AND GALAXY EVOLUTION 35 Your text 108

7 10 SKA

106 Sloan

2dF 105 6dF Number

104

103

105 106 107 108 109 Volume (cubic Mpc)

Figure 2.4: Number of detected galaxies and volume sampled for a number of recent large galaxy redshift survey projects. Optical surveys are shown in blue and H I (HIPASS) surveys in red. The SKA would sample to greater depths than presently possible, and would sample galaxies in atomic hydrogen, thereby viewing galaxian masses independent of stellar content or star formation history.

with redshift contains information about the different cosmological parameters and is a very powerful tool for testing various structure formation models. Determining clustering properties at the earliest possible epochs will be crucial, as pointed out by Van de Weygaert and Van Albada (1997). The use of H I as a tracer of the galaxy population offers the additional advantage that for the galaxies with optical photometry one can use the Tully-Fisher method to derive distances independent of redshift and probe the peculiar velocity field to determine the mass density field in the universe. The number of data points from a survey as described above will greatly exceed the present catalogues of peculiar velocities (Dekel 1994, Sigad et al. 1998) and, moreover, have greater precision. A comparison of the mass density field with the actual distribution of galaxies provides a means of putting strong constraints on cosmological parameters.

2.2.3 The Lyα forest seen in the 21-cm H I line Numerical models suggest that the high redshift Lyα forest is part of a complicated structure of gaseous ﬁlaments and sheets formed in the gravitational ﬂuctuations in the underlying dark matter potential (Zheng et al. 1997, Hernquist et al. 1996, Cen et al. 1994). The exact structure of the gas giving rise to the Lyα forest and the precise connection to galaxies formed at early epochs from the original density 36 CHAPTER 2. FORMATION AND EVOLUTION OF GALAXIES

fluctuations still remains a matter of debate. Lanzetta and coworkers (Lanzetta et al. 1995 and references therein) compare Lyα forest redshifts with the redshifts of galaxies along the same line of sight and find redshift coincidences for galaxies with projected separations of typically 10 to 100 kpc. They argue that the Lyα forest arises from galaxy disks or halos. Shull and van Gorkom and coworkers (Shull et al. 1998, van Gorkom et al. 1996) use the VLA to measure the H I content of galaxies along the sight line to fairly low redshift Lyα forest lines (z 0.06) and find much larger projected separations of 400 kpc. These results can∼ no longer be explained ∼ as due to galaxy halos, but imply the kind of filamentary structure suggested by the simulations. Using the inner part of SKA with baselines up to 10 km and integrating for ∼ 100 hours one can reach limiting column densities of 1017 atoms cm−2. The strongest Lyα forest lines correspond to column densities just below 1017 atoms cm−2. On the other hand, one might expect to find somewhat higher column densities nearer to the intervening galaxies. Galaxy disks typically truncate at much higher column density levels ( 1019 atoms cm−2), so with observation such as mentioned here one would ∼ 17 19 −2 probe the H I outside galaxies at levels between 10 and 10 atoms cm . This is an extremely interesting regime and the structure of the H I in the few 100 kpc vicinity of galaxies will no doubt have clues for resolving the issue of whether the Lyα forest arises in galaxy halos or from filamentary gaseous structures.

2.2.4 High Redshift CO In order to understand galaxy formation and evolution, as well as the early history of the universe, it is essential to study the properties of galaxies at high redshift. Our knowledge of the processes by which gas becomes stars, and stars become galaxies is seriously incomplete. The primary tools used to study high redshift galaxies, optical and radio continuum observations, provide only an indirect measure of the gas content of early galaxies. CO observations, on the other hand, could provide us with a direct measure of the molecular gas content in high redshift galaxies. At high redshifts the CO 1 0 line will be sufficiently redshifted to be detected → by the SKA. The rest wavelength of the CO 1 0 line is 2.6 mm, which corresponds to a frequency of 115.38 GHz. Figure 2.5 shows→ the observed wavelengths for the lower CO transition lines as a function of redshift. If the SKA has a minimum operating wavelength of 1.2 cm, the CO 1 0 line will be shifted into the observable wavelength band at redshifts greater than→ or equal to 4. The CO 2 1 line will also be sufficiently shifted at redshifts greater than or equal to 8. → To date, there have only been a handful of CO detections at high redshift. The most significant of these are shown in Table 2.2. It should be noted that the detections of CO in IRAS F10214+4724 and the Cloverleaf QSO were aided by amplification due to gravitational lensing, and that the detection in PC 1643+4621A is somewhat controversial. 2.2. LARGE SCALE STRUCTURE AND GALAXY EVOLUTION 37

2.5

2 CO 1−0

1.5 SKA Minimum Wavelength Wavelength (cm) 1 CO 2−1

0.5 CO 3−2

0 2 3 4 5 6 7 8 9 10 Redshift

Figure 2.5: Wavelength of Observable CO transitions as a function of Redshift.

Here we consider two scenarios in predicting the detectability of CO at high redshifts: The ﬁrst requires a burst of initial star formation at an early epoch, and has been modeled theoretically by several authors. The second is a constant luminosity model in which a typical spiral galaxy is observed at large redshifts. Evans et al., (1996) reported negative CO detection from 11 high redshift powerful radio galaxies in the range 1

Object z Ref. Line L′ S ∆ν S · peak (h−2 K km s−1 pc2) (Jy km s−1) (mJy) IRAS F10214+4724 2.29 a 1 0 1.4 1011 2.4 13.3 b → × PC 1643+4621A 3.14 1 0 9.4 1011 10 15 c → × ≈ 53W002 2.39 1 0 1.2 1011 1.9 5.0 d → × BR 1202-0725 4.69 5 4 1.8 1010 2.7 9.3 e → × Cloverleaf QSO 2.56 3 2 6.1 1010 8.1 23 → × Table 2.2: CO Detections at High Redshift

aTsuboi & Nakai 1992 bFrayer et al., 1994 cYamada et al., 1995 dOhta et al., 1996 eBarvainis et al., 1994

Observed luminosities for IRAS F10214+4724, PC 1643+4621A, 53W002, BR 1202- 0725, and the Cloverleaf QSO are shown on the plot with asterisks. BR 1202-0725, and the Cloverleaf QSO are observed in high-J states. The assumption has been made for the plot in Figure 2.6 that the ratios with the J = 1 0 line is about 1 (see Solomon et al., (1992) and the discussion below). The CO non-detections→ observed by Evans et. al (1996) are also shown on the plot with open circles. (These indicate an upper limit to the CO luminosity.) Figure 2.7 shows the integrated ﬂux density as a function of redshift for a source with a CO luminosity as given by the galactic wind model, as well as for the standard galaxy (constant luminosity with z) with a CO luminosity of 1.5 109 h−2 K km s−1 pc2. The same observational data for the sources in Figure 2.6 are also× shown in Figure 2.7. The integrated ﬂux density, was derived from the expression,

L′ h2 ν2 q4 (1 + z) S ∆ν = rest 0 , (2.13) · 2.92 1014 Q2 × ′ −2 −1 2 where L is the luminosity in h K km s pc , νrest is the rest frequency in GHz, z is the redshift, q0 is the deceleration parameter, and Q is the cosmological term associated with luminosity distance:

Q = q z +(q 1) 1+2q z 1 . (2.14) 0 0 − 0 − q A ﬂat universe (q0 =0.5) was assumed in all calculations. The Hubble constant is given by H = h 100 km s−1 Mpc−1. L′ may be related to the luminosity L, which 0 × has units of solar luminosity, by the formula:

8πk ν3 L = b rest L′. (2.15) c3 ! A useful quantity in determining the ability of the SKA to detect CO at high redshift, however, is the predicted peak ﬂux density of CO at a given redshift. This 2.2. LARGE SCALE STRUCTURE AND GALAXY EVOLUTION 39

12 10

11 10 Galactic Wind Model

10 10

9 10 Constant Luminosity

Luminosity (h^(−2) K km s^(−1) pc^(2)) 8 10

7 10 0 1 2 3 4 5 6 7 8 9 10 Redshift

Figure 2.6: CO Luminosity as a Function of Redshift. Observations from Table 2.2 are shown as asterisks; open circles are upper limits observed by Evans et al., 1996

1 2 10 10

1 10 0 10

0 Galactic Wind Model 10 Galactic Wind Model −1 10

−1 10 Constant Luminosity −2 10 Constant Luminosity −2

Peak Flux Density (mJy) 10

Integrated Flux Density (Jy km s^(−1)) −3 10 −3 10

−4 −4 10 10 0 1 2 3 4 5 6 7 8 9 10 0 1 2 3 4 5 6 7 8 9 10 Redshift Redshift

Figure 2.7: Integrated (right) and peak (left) ﬂux densities of CO as a Function of Redshift for both constant luminosity and galactic wind evolution. 40 CHAPTER 2. FORMATION AND EVOLUTION OF GALAXIES

15 Constant Luminosity Signal to Noise Ratio

5 SNR = 5

0 5 10 15 20 25 Redshift

Figure 2.8: The signal-to-noise ratio for detection of CO in a 24 hour observations as a Function of Redshift. A bandwidth of 28 km s−1 is assumed.

can be approximated by dividing the integrated flux density S ∆ν by the average velocity width. Based on existing data, a reasonable value for· the average velocity width is 300 km s−1. Estimates of peak flux densities based on the galactic wind model, as well as a constant luminosity model, are shown in Figure 2.7. Existing data are again shown with asterisks. Between redshifts of about 4 and 8, the galactic wind model peak flux densities are between about 0.1 and 1 mJy. For a 24 hour observing run, with a bandwidth of 2 MHz (28 km s−1 at λ =1.4 cm, at a system temperature of 70 K, the SKA would have a flux density sensitivity of ∆S = 0.8216 µJy. (A bandwidth of 2 MHz would allow 10 points across a line spectrum integrated over a galaxy with a velocity width −1 of 300 km s ). Using the conservative estimate of Speak =0.25 mJy for the peak flux density, the signal-to-noise ratio for a CO detection would be Speak 300. ∆S ≈ This result indicates that, for the galactic wind model at least, the CO flux densities at redshifts observable by the SKA (z > 4) would easily be high enough for the SKA to detect. In this simulation, CO line is undetectable beyond z = 10, since z = 10 is the epoch of galaxy formation for this model. The signal-to-noise ratio is much lower for the constant luminosity model, as shown in Figure 2.8. Nonetheless, the plot indicates that an average spiral galaxy could be detected in CO (at 5σ) out to a redshift of about 20. In comparison, for a similar observing run at its system temperature of 160 K, the 2.3. DEEP CONTINUUM FIELDS 41

VLA would have a flux density sensitivity of ∆S = 141.7 µJy. This gives a signal- to-noise ratio that is too low for detection of both the constant luminosity CO line and the galactic wind model CO line. It is thus not surprising that the data points in Fig. 2.7 are above the model curves. It is possible that a large number of sources emit CO at flux densities matching the model, but since such low flux densities cannot currently be detected, only the more outstanding data values are seen. For very high redshifts we would expect the higher J lines to be excited as a result of the higher microwave background temperature, which increases with redshift as T = 2.7(1 + z). At z 4.5, this temperature becomes comparable to the cmb ≈ mean gas temperature of lower-redshift CO clouds. This results in a change in the relative populations of the CO states, increasing the strengths of the higher-frequency CO lines, and reducing the strengths of the lower-frequency CO lines. Solomon et al., predict that the CO 3 2 line is always comparable in strength to the CO 1 0 line for all very high redshift→ galaxies, because of the warmer microwave background.→ The analysis of Solomon et al., demonstrates that the relative strength of the higher excitation lines of CO is very dependent on gas density. Thus observations of a range of CO transitions is critical to sample the full range of gas conditions An analysis by Silk & Spaans (1997) for high-density gas, 2 106 cm−3 suggests that the × population of the J=1 state decreases by a factor of about 3 as z goes from 5 and 30. In conclusion, the SKA will to be able to observe redshifted CO lines beyond z = 10 from both large spirals and ellipticals. Starburst galaxies will be observable, although these galaxies will also be well observed using the large millimeter arrays (MMA/LSA). The SKA and the MMA/LSA will be complementary instruments to fully sample the range of physical conditions (density and temperature) of the CO that are expected to exist during the era of galaxy formation.

2.3 Deep Continuum Fields

2.3.1 Extragalactic Radio Sources Extragalactic radio sources cover a wide range of luminosity extending from 1019 Watts/Hz for normal spiral galaxies to more than 1028 Watts/Hz for the powerful FRII radio galaxies and radio loud quasars. Intermediate in luminosity are the less powerful FRI radio galaxies, the radio quiet quasars, and galaxies with active star formation. With few exceptions the observed radio emission is non-thermal synchrotron radiation. The nature of individual radio sources, e.g., normal galaxies, starbursts, AGNs, FRI or FRII radio galaxies, may be distinguished by their luminosity, the observed radio morphology, the radio spectral index, variability, their optical counterpart and spectral features, and observed characteristics in other wavelength bands, particularly x-ray and infrared. In the powerful radio galaxies and quasars, the source of energy is thought to lie in a massive central engine, whereas in many of the sources of intermediate luminosity, 42 CHAPTER 2. FORMATION AND EVOLUTION OF GALAXIES

the energy source appears to lie in regions of active star formation and supernovae activity which accelerate relativistic electrons into the interstellar medium. This star forming activity may be found in a population of faint blue galaxies or associated with the strong IR galaxies detected in the IRAS survey at 60 microns. Because of the close association of 60 micron and radio wavelength emission, both believed to be closely linked to star forming activity, deep radio observations are sensitive to star formation at early epochs, unaffected by the obscuration which plagues optical and infrared observations. Star forming rates may be estimated from the observed radio flux density, but only if the contribution to the observed radio flux density from a massive central engine can be determined from the radio or optical data. Typically, both the FRI and FRII radio galaxies, as well as radio loud and radio quiet quasars have linear dimensions of the order of a few hundred kiloparsecs with structure characterized by radio lobes which are well separated from the optical counterpart, but often joined to the optical counterpart by a thin jet. FRI and FRII radio galaxies are optically thin with radio spectral indices about -0.8, but may have a self absorbed flat spectrum core. The compact radio cores of quasars and AGN typically have angular dimensions much less than an arcsecond, and due to self absorption have flat radio spectra with observed spectral indices near zero. These very compact radio sources are often variable, and are identified with galaxies showing broad emission line spectra in their nuclei. The radio emission from galaxies with active star formation is typically confined to dimensions comparable to that of the galactic disk. Optical counterparts are often unusually bright at far infrared wavelengths as a result of the absorption of UV emission from young massive stars by dust and its subsequent thermal reradiation. Radio source surveys made over a wide range of wavelengths and flux density have catalogued about two million discrete radio sources above a limiting flux density of about 10 microJansky (µJy). Optical identification and spectroscopic redshifts show that most catalogued radio sources stronger than a few mJy are relatively distant powerful radio galaxies or quasars with radio luminosities greater than 1025 Watts/Hz, or nearby normal or nearly normal galaxies with much weaker radio emission. The space density of powerful radio galaxies and quasars quickly converges beyond redshifts of unity, so that nearly all of these powerful sources are included in the radio source counts above one mJy. At µJy levels, the radio source count again steepens, corresponding to a new population of radio sources. Nearly all µJy radio sources can be identified with a mixture of low luminosity AGN and faint star forming galaxies which are often found in pairs or small groups. These µJy radio sources mostly have redshifts between zero and one with corresponding radio luminosity between 1020 and 1025 W/Hz. Periods of active star formation may be driven by mergers which may also enhance their disk emission as well as fuel their central engine. The observed peak in the redshift distribution of µJy radio sources is comparable with that found for strong radio galaxies, for quasars, and the population of faint star forming galaxies suggesting that the evolutionary scenario is comparable for all of these populations, and that 2.3. DEEP CONTINUUM FIELDS 43 conclusions about star formation rates, the epoch of quasar and galaxy formation are valid and are not the result of obscuration by dust.

2.3.2 The SubmicroJansky Sky With the improvement in sensitivity given by the SKA, it will be possible detect radio emission as weak as 10 nanoJy in only a few hundred hours integration time, over a factor of a hundred fainter than the weakest sources seen so far with the VLA. At this level, not only will it be possible to detect radio emission from star forming activity out to redshifts of ten or more, if it exists, but for the first time, even normal spiral galaxies with P=1019 W/Hz will be detected out to cosmologically interesting distances (z=1). High resolution, multi-wavelength observations will be critical to distinguish among the different emission mechanisms for extragalactic radio sources. Nothing is known about the radio source count below one µJy. Extrapolation of the VLA 4 cm and 6 cm deep surveys suggests that there are about 100 sources/sq arcmin above 1 µJy at 20 cm. To model the radio sky, the known 1.4 GHz source counts (Windhorst et al. 1993; Hopkins et al. 1998) were used as the initial starting point. The known source counts were extrapolated down to a flux density of 1 nJy, subject to known limits on the source count slope (due to the CMB) and implied limits from the number of possible optical counterparts (Windhorst et al. 1993). The distribution in apparent size of radio sources at 1.4 GHz has been characterised as a function of flux density by Windhorst et al. (1990), and compared with the Phoenix Deep Survey sample (Hopkins, 1997) by way of verification. This distribution has been used to assign apparent sizes to a list of sources with given flux densities. The result of this is to produce a simulated distribution of sources with the same statistical properties (source counts and angular size distribution) as the real sky. The axial ratio of the simulated sources has been modeled simply by a uniform distribution between values of 0.2 and 1. With the angular size and the axial ratio for each source, a simulated image was constructed by adding elliptical gaussians at random locations and position angles. The peak value of the gaussian is defined by the flux density of the source. As a first step in refining this very simple model the source counts were divided between two populations, broadly described as “starbursts” and “AGNs.” This was accomplished by using the known fraction of these populations as a function of flux density (Wall & Jackson 1997; Hopkins et al., 1998). In addition, to mimic the double-lobed nature of many real AGNs, a pair of adjacent elliptical gaussians have been used, rather than the single elliptical gaussian used for starbursts. At brighter flux densities the angular size distribution will not necessarily be valid for the AGN population. A simulation of an 8-hour SKA observation at 20 cm of a region the size and shape of the Hubble Deep Field is shown in Fig. 2.9. With a flux density limit of 100 nJy, over 2700 sources are predicted (a source density of 5 109 sr−1). The different ≈ × populations are indicated by the colours, starbursts being blue and AGNs being red. If the source count continues unchanged down to 10 nanoJy, there will be about 44 CHAPTER 2. FORMATION AND EVOLUTION OF GALAXIES

Figure 2.9: Simulated λ20 cm continuum observation with the SKA of a region the size and shape of the Hubble Deep Field. The radio source population is derived from an extrapolation of the known population above 1 µJy. With a 5σ detection limit after 8 hours of about 100 nJy, the simulated SKA image contains 2700 sources. Starburst galaxies are shown in blue and Radio Galaxies and Active Galactic Nuclei in red. 2.4. PROBING DARK MATTER WITH GRAVITATIONAL LENSING 45

10,000 radio sources/sq arcmin above 10 nanoJy, or about three sources/sq arcsec. Assuming for the moment that these are point sources, with a resolution of 0.1 arcsec there will be about 30 beamwidths/source, just adequate to keep the effects of confusion negligible. More likely the nanoJy radio sources have dimensions comparable with those of the optical discs of galaxies, or a significant fraction of an arcsec at cosmologically significant redshifts, and at the full sensitivity of the SKA, individual sources will appear blended at the faintest flux density levels. Resolutions as least as good as 0.1 arcsec at 20 cm (500 km dimensions) will therefore be important not only to separate individual sources, but to image each source with sufficient detail to tell whether the emission comes from the entire disk characteristic of normal galaxies, is concentrated within a few hundred parsecs of the nucleus, characteristic of star formation, or is in a point source at the nucleus characteristic of a massive central engine. At longer wavelengths, the greater source flux density is roughly canceled by the lower sensitivity of the SKA, so that the limiting source density remains approximately constant with wavelength. But, the required dimensions of the array scales with wavelength to maintain a resolution of 0.1 arcsec, comparable with that of the Next Generation Space Telescope, and adequate to separate individual sources and to image radio emission from distant galaxies. This implies baselines of 2500 km at least for those elements working at 1 meter wavelength (300 MHz). At the short wavelength limit, the extraordinary sensitivity of the SKA will allow even normal galaxies at cosmologically interesting redshifts to be imaged with a resolution and image quality far superior to that of any other telescope, existing, or planned, operating at any wavelength band in space or on the ground. Such observations will be crucial to outline the early history of the formation of stars, galaxies, and quasars, without uncertainties due to possible obscuration by dust or other material.

2.4 Probing Dark Matter with Gravitational Lens- ing

The technique of weak gravitational lensing has proven to be an important tool to study mass distributions in the universe. The projected mass distribution of foreground gravitational structures distorts the images of the faint background galaxies. As a result, gravitational lensing provides a direct measurement of the projected mass density (e.g. Kaiser & Squires 1993). Until recently, massive structures in the universe were studied through dynamical analysis of their luminous components. These studies have shown that large amounts of dark matter exist in the universe. For clusters of galaxies a popular method uses the motions of the galaxies to estimate the mass using the virial theorem. One can also estimate the cluster mass proﬁle from X-ray observations when one assumes hydrostatic equilibrium and spherical symmetry (e.g. Allen & Fabian 1994). Both methods assume some dynamical state or geometry in order to obtain the 46 CHAPTER 2. FORMATION AND EVOLUTION OF GALAXIES

mass or a mass profile. The advantage of gravitational lensing is the fact that no such assumptions are needed. In the regime of weak gravitational lensing one can calculate the projected mass surface density up to some additive constant from the observed distortion pattern (Kaiser & Squires 1993; Kaiser et al. 1995; Schneider & Seitz 1995; Schneider 1995; Squires & Kaiser 1996). Since the first successful measurements of the weak gravitational distortions (Tyson, Valdes & Wenk 1990), many massive clusters of galaxies have been studied (e.g. Bon- net, Mellier, & Fort 1994; Fahlman et al. 1994; Squires et al. 1996b; Luppino & Kaiser 1997). In principle one can measure the gravitational distortion out to large radii from the cluster centre, beyond the radii where X-ray observations or cluster kinematics can be used to determine the mass distribution. So far, most weak lensing studies of clusters of galaxies have been undertaken using data from ground based optical telescopes. These are affected by atmospheric seeing, which causes the images of the faint background galaxies to be enlarged and more circular. Recently these efforts have been extended with the weak lensing analysis of a cluster of galaxies, CL 1358+62, using HST observations with a large field of view (Hoekstra et al. 1998). Up to now, other weak lensing studies of clusters of galaxies with HST have been limited to cluster cores (C. Seitz et al. 1996; Smail et al. 1997). Those observations consisted of single pointings, thus suffering from the limited field of view of the HST (about 2 arcmin). By using a mosaic of 12 pointings, a total field of view of approximately 8 by 8 arcmin could be obtained. This combination of space based observations and a large field of view provided an opportunity to study the cluster CL 1358+62 in great detail. An advantage of HST observations is the high number density of galaxies one can reach. Previous HST studies have achieved 100 galaxies arcmin−2 routinely (e.g. C. Seitz et al, 1996; Smail et al. 1997). With≃ a one hour exposure per pointing, a number density of 50 useful background galaxies arcmin−2 can be obtained. ∼ Another important advantage of HST over ground based observations is the size of the point spread function (PSF). Most of the faint objects are small. To recover the lensing signal one needs to correct for the effect of seeing (Bonnet & Mellier 1995, Kaiser, Squires & Broadhurst 1995, Luppino & Kaiser 1997, Fischer & Tyson 1997). For objects with sizes comparable to the PSF, these corrections become very large, amplifying the uncertainty in the ellipticity due to photon noise. As a result, the scatter in the derived ellipticities of the galaxies is larger than the expected scatter due to their intrinsic shapes. Consequently, for a given number density of background objects, the accuracy of weak lensing studies based on HST observations will be higher than the results from ground based data. Even with the high angular resolution of the HST, a major limitation to the detection of gravitational shear is introduced by the distortions of the PSF over the area of each detector array. This is illustrated in Fig. 2.10 where HST archival data for the globular cluster M4 have been analyzed for the width and orientation of the stellar profiles. After careful calibration and correction it is in fact possible to extract the weak 2.4. PROBING DARK MATTER WITH GRAVITATIONAL LENSING 47

Figure 2.10: The upper panel shows the “polarization” ﬁeld of stars taken from observations of the globular cluster M 4. The orientation of the sticks shows the direction of the major axis of the HST PSF whereas the length is proportional to the size of the anisotropy. (From Hoekstra et al. 1998) 48 CHAPTER 2. FORMATION AND EVOLUTION OF GALAXIES

Figure 2.11: Smoothed distortion field g from the blue and red background galaxies toward CL 1358+62. The length of the sticks denote the size of the distortion. The distortion field is calculated by smoothing the observed field from the individual galaxies with a Gaussian of width 0.4 arcmin(FWHM is indicated by the shaded circle). The characteristic pattern due to gravitational lensing is clearly visible. (From Hoekstra et al. 1998)

lensing signature as illustrated in Fig. 2.11 out to a distance of 1.5 Mpc from the ∼ cluster center. The observed distortion is consistent with a singular isothermal sphere model with a dispersion of 780 50 km/s. The total projected mass within a radius of 1 Mpc, corresponding to this± model is (4.4 0.6) 1014 M . The errors given here ± × ⊙ represent the random error due to the ellipticities of the background galaxies. The uncertainty in the redshift distribution introduces an additional, systematic error of 10% in the weak lensing mass. The weak lensing mass is slightly lower than dynamical∼ estimates and agrees well with X-ray mass estimates. The mass distribution is elongated similar to the light. The axis ratio of 0.30 0.15 and position angle ± of 21◦ 7◦ were measured directly from the observations and agree very well with the− previous± strong lensing determination (Franx et al. 1997). A two-dimensional reconstruction of the cluster mass surface density shows that the peak of the mass distribution coincides with the peak of the light distribution. A mass-to-light ratio of (90 13)h M /L is indicated, and this appears to be constant with radius. ± 50 ⊙ V ⊙ The impact of the SKA on the ﬁeld of weak lensing will be particularly profound. With the sensitivities noted in the Introduction, we expect detected source densities 2.5. ACTIVITY IN GALACTIC NUCLEI 49

of between 100 and 400 arcmin−2 in an 8 hour integration. Some 98% of these sources will correspond to the same normal galaxies visible to the HST, while the remaining 2% will be active galactic nuclei and radio galaxies. While these source densities are comparable, but superior to those of the HST, there are two important differences. Firstly, the SKA PSF is both compact (smaller than about 0.1 arcsec) and extremely well-defined (to about one part in a million) over the entire field of view. And secondly, the enormous instantaneous field of view (of about one square degree) is sufficient to probe scales of some 20 Mpc on a side (at z = 0.3) per pointing. This should enable clean measurement of cluster mass surface densities as well as routine detection of the weak lensing signature due to large-scale structure, for which the first tentative detections have recently been made (Schneider et al. 1998).

2.5 Activity in Galactic Nuclei

The properties of galaxies as a whole are profoundly influenced by the energy release processes which take place in their centres. The mass and spin of the central black hole and the rate of nuclear feeding are thought to play dominant roles in establishing the properties of galactic nuclei, and differences in these parameters may lead to the various forms of nucleus observed: active radio-loud, active radio-quiet, LINERs, normal galaxies etc. The process of accretion releases energy as gamma ray, X-ray and UV continua which, in active nuclei, can photo-ionize material orbiting within the central tenth of a parsec producing broad optical emission lines, and further out, up to a few kpc, narrow optical emission lines. In radio-loud nuclei, relativistic jets of material are ejected along the polar axis of the black hole and in some cases have sufficient energy to escape into intergalactic space as classical double-lobed sources. On their way out through the narrow line region the jets may play a role in shock-exciting the lines observed. In less active nuclei like LINERs, the photo-ionizing field is probably much more dilute. Orientation effects also play a major role in the appearance of galactic nuclei. For more than a decade, it has been postulated that geometrically thick and optically thick dusty molecular tori are located in the accretion plane of a nucleus which are capable of obscuring the central region from direct view. Only recently has it become clear through HST imaging observations and VLBI observations of H I absorption and maser emission that the postulated ring-like gas/dust structures are indeed found in the equatorial planes of galactic nuclei on scales of pc to kpc. The phenomenology of nuclei is very varied, and their nuclear properties are distributed over many orders of magnitude. Producing a unified picture of the factors governing the distribution of mass and mass motions in the inner few kpc in galactic nuclei and how this picture evolves with cosmological epoch are major tasks con- fronting us. Specific questions which may be asked are (e.g. Lawrence, 1999):

1. do all galaxies have massive dark objects in their centres; 50 CHAPTER 2. FORMATION AND EVOLUTION OF GALAXIES

2. are these massive dark objects super-massive black holes like those deduced to be present in the Galactic Centre and in NGC4258;

3. why are they so dark;

4. do all galaxies have an active nucleus like the 40% of optically selected samples or the 40% of ellipticals in radio selected samples;

5. does activity in the nucleus correlate with host galaxy properties e.g. radio loud quasar hosts are large luminous ellipticals like radio galaxies and 50% of radio quiet quasar hosts; radio power in the nucleus may correlate with the absolute blue magnitude of the host;

6. how is the activity deep in the nucleus fuelled. We have direct observations of radio structures on scales of less than 1 pc but there is no direct information on the inward transport of matter on scales close to the putative black hole;

7. is activity deep in the nucleus related to star formation activity further out, as may be the case in ultra-luminous infrared galaxies and high redshift galaxies;

8. what is the detailed physics of the energy outﬂows?

2.5.1 The SKA and Active Galactic Nuclei The Square Kilometre Array with its spectacular increase in sensitivity over conventional radio telescopes will play a major role in answering some of these questions, particularly the question of how widespread is nuclear activity in galaxies. Does every normal galaxy contain a black hole? The sensitivity of the SKA will also be crucial in greatly enlarging the number of objects in which the environment of the central engines can be studied through VLBI H I absorption measurements and in which the mass of the central engine can be estimated through water megamaser studies (see chapter 2.5.1). It will greatly increase the size of VLBI surveys for cosmological purposes, and allow studies of edge effects in compact radio emitting flows for the first time.

Black holes in normal galaxies The radio signature of a black hole in a galactic nucleus is generally taken as the presence of an unresolved core component with a ﬂat radio spectrum and variable ﬂux density, and a jet, or jets, emanating from the central core. In normal galaxies, detection of a jet(s) should resolve any confusion with starburst galaxies dominated by a central supernova remnant. Clearly, high angular resolution is an instrumental prerequisite for the radio detection of black holes in a large sample of normal galaxies, in addition to sensitivity (see Fig. 2.12 from Wrobel, Condon and Machalski for a display of current sensitivity limits). The ability to observe at frequencies up to 20 2.5. ACTIVITY IN GALACTIC NUCLEI 51

Figure 2.12: Dependence of peak VLA power on distance for 374 UGC galaxies

GHz is also a prerequisite in order to maximise the detection rate by observing near the peak in the radio spectrum. The sensitivity required is of order micro Janskys and the angular resolution of order milli arcseconds. The strawman SKA configuration with 500 km maximum baselines has the sensitivity and frequency coverage required, but not the angular resolution. Using the SKA as the major element in a global VLBI array or increasing the baseline lengths of the SKA itself are options that need to be examined in subsequent studies of the concept. The sensitivity of the SKA in VLBI configurations is summarised in the final section.

The environments of central engines: the structure and kinematics of gas in the nucleus A centrally condensed configuration for the SKA would be a magnificent instrument for the detection of H I absorption associated with active galactic nuclei or against background AGNs. The new UHF receiver complement on the WSRT can detect 0.2% lines against 1 Jy sources. With SKA, 0.2% lines could be detected against milliJy sources which opens up the prospect of studying the structure and kinematics of gas in radio nuclei with a wide range of luminosity. Once lines are detected, the higher angular resolution afforded by global VLBI (10 milli arcsec) is needed to make detailed studies of the structure and velocity field of the 52 CHAPTER 2. FORMATION AND EVOLUTION OF GALAXIES

Figure 2.13: H I absorption in 3C236

gas. Baselines to a space borne antenna are required to obtain suﬃcient resolution for the continuum structure. Work done so far by Conway and collaborators and Peck and Taylor have shown that H I gas in nuclei can be found in circumnuclear tori, amorphous clouds, and in distinct oﬀ-nucleus clouds. An example of the latter is shown in Fig 2.13 hich displays a cloud apparently blocking the SE-going jet 1 kpc from the nucleus in the giant radio galaxy, 3C236 (Conway and Schilizzi, in∼ preparation).

Physics of Outflows The most common parsec-scale radio structures in AGN have a so-called core-jet” morphology. The jets” are collimated and connect a flat spectrum and usually brighter feature (core) with extended steeper spectrum lobes” (e.g. Pearson 1996). Smaller groups of sources exhibit other types of morphologies, such as compact symmetric, compact doubles and “complex’ sources (Readhead 1995). At present, state-of-the-art VLBI images of these sources have noise levels of 50 microJansky per beam, where the beam size is of the order of one to several milli arcsec (Zensus et al. 1995). However, extensive studies of prominent targets (3C 84, Dhawan et al. 1998; 3C236, Conway 1996) make it clear that the depth of these VLBI images is several orders of magnitude short of reaching the level at which the lowest brightness features in AGN structures will be blended by the background radiation (see Figure 2.14). The SKA as a part of a global VLBI array will enable imaging with a sensitivity of micro Jy per milli arcsecond-size beam at centimeter and decimeter wavelengths. Such imaging sensitivity will allow us to “see” the edges of plasma flows along and across jets, and therefore to directly compare the predictions of various models with physical 2.5. ACTIVITY IN GALACTIC NUCLEI 53

Figure 2.14: The 1995 image of 3C 84 at λ = 2 cm. The contours are 10,10,14,20,. . . ,2560 mJy/beam. The peak is 2.97 Jy/beam, the oﬀ-source rms − residual is 0.33 mJy/beam, and the beam has a FWHM of 0.77 0.58 mas, at PA = 34◦. (from Dhawan et al. 1998). × 54 CHAPTER 2. FORMATION AND EVOLUTION OF GALAXIES

reality. This will be particularly interesting in regions where the jets pass through the optical emission line regions and in regions where jet-induced star formation is occuring. On larger scales, the SKA as a stand-alone instrument will provide data on the structure of plasma flows across jets on kiloparsec scales. A global ground-based VLBI array enhanced by the SKA will make it possible to study core-jet physics in new classes of extragalactic sources, previously unreachable by VLBI due to their faintness. A combination of the SKA with the largest ground-based radio telescopes (VLA, GBT, Effelsberg, WSRT) in the 5 – 8 GHz bands would allow the study of parsec-scale structures in cores of spiral galaxies similar to our Galaxy at cosmological redshifts (mini-AGN). Polarization VLBI observations have proved to be one of the most powerful tools of studying the geometry of magnetic fields and plasma flows on parsec scales (e.g. Wardle et al. 1994). At present, only several tens of highly polarized sources (with linearly polarized flux density of the order of several percent of the total flux density at cm-dm wavelengths) are within the reach of VLBI systems. To substantially enlarge the number of extragalactic radio sources measurable with VLBI polarimetry, one needs to increase the sensitivity of the VLBI systems by a factor of 10 to 100.

Cosmological Tests with Parsec-scale Radio Structures An increase in recent years in the amount of VLBI data on milliarcsecond structures in AGN has rekindled interest in the classical idea of using radio sources as cosmological “standard rods” (Hoyle 1959). Based on a comparison of the predicted and observed dependence of the apparent angular size of a standard object as a function of redshift, one can derive cosmological parameters (θ z test). Recent attempts to use VLBI images for θ z tests have been made by− Kellermann (1993), Gurvits − (1994,1998), Pearson et al. (1994), and Wilkinson et al. (1997). In several follow- up publications, authors have critically analyzed these attempts, pointing out the necessity to include in the models not only the deceleration parameter qo but also the cosmological constant Λ (see Krauss and Schramm 1993, Stelmach 1994, Kayser 1995, Jackson and Dodgson 1996). It has also been noted that statistical confidence in the results, in particular estimates of the value qo, is still in need of improvement (e.g. Dabrowski et al. 1995, Stepanas and Saha 1995). While the background principles of using parsec-scale sources as cosmological “standard rods” are simple and undisputed, a practical implementation of these tests faces serious difficulties, most of which are imposed by various selection and masking effects (eg. Wilkinson et al. 1998 and Vermeulen 1996). The most obvious selection effect is illustrated by Fig. 2.15, in which the luminosities of 330 AGN used for the θ z test (Gurvits Kellermann and Frey. 1999, in press) are plotted against redshift. Due− to the flux density limits of that ad-hoc sample of about 1 Jy, the luminosity of the low- and high-redshift AGN counterparts are mismatched by several orders of magnitude. Although one can try to use such a luminosity mismatched sample for cosmological tests, a much better result could be achieved by composing a luminosity- matched sample. To do so, one needs to image with VLBI, AGN’s at higher redshifts 2.5. ACTIVITY IN GALACTIC NUCLEI 55

Figure 2.15: Compact component luminosity as a function of redshift for a sample of 330 ﬂat spectrum extragalactic radio sources. The solid lines show luminosities of sources with ﬂux density of 1, 0.1 and 0.01 Jy, calculated under the assumption that the spectral index α = 0. (From Gurvits, Kellermann and Frey 1999.)

(z > 1) with total flux densities 10–1000 times lower than those imaged to date, i.e. at a level of 0.1 – 10 mJy. Furthermore, as shown by Dabrowki et al. (1995), a cosmologically conclusive result will require a sample of several thousand sources at this level of flux densities. If we add to this other masking effects (such as Malmquist bias, spectral properties, orientation bias, etc.), it translates into an extensive observing program of VLBI imaging of a sample of 104 AGN’s with total flux densities of milli- and sub-millijansky level.

2.5.2 Sensitivity of the SKA in VLBI Arrays The sensitivity of the SKA as an element of a VLBI array can be characterised in two ways - single baseline sensitivity and image noise. Assuming an integration time of s −5 120 , a recorded bandwidth of 512 MHz, and Tsys/Aeff =5 10 for the SKA and 10−2 for a 70 m telescope, the sensitivity for a single baseline× between SKA and a 70m telescope is 15µJy. The sensitivity to a 25 m telescope (e.g. a second generation ∼ space VLBI antenna) would be 40µJy. ∼ If the SKA is an element of an array with ten 70m class telescopes, the theoretical 56 CHAPTER 2. FORMATION AND EVOLUTION OF GALAXIES

image noise after an 8h observation is 250 nJy. For µJy sources, only the ten baselines to the SKA would contribute∼ and phase referencing would be required to phase up the array. Image simulation should be carried out to quantify the effects of such a configuration. For comparison, were the SKA configuration itself to be extended to global dimensions, the sensitivity of a single baseline between two of the 200m diameter stations would be 25µJy and the image noise would be 60 nJy. ∼ ∼ Note that the brightness temperature of a 1 µJy radio source which is 1 milli arcsec in size is 105 K. ∼

2.6 Circum-nuclear MegaMasers

Since the discovery of the first astronomical maser by Weaver et al. (1965), a great deal of effort has been expended on understanding their physics. However, even now, because of their complexity, in many maser scenarios we have only a vague understanding of the pumping mechanisms and the physical characteristics of the masers themselves. Despite this ignorance, masers have become powerful tools for probing the kinematics of gas in a variety of astrophysical environments. Maser regions are small and bright, and have a small velocity range, which makes them relatively easy to observe at extraordinarily high precision in location and velocity within their host environment. They have been used to study the astrophysics of both the birth and the death of stars, but recently a most spectacular result emerged from maser studies being the first conclusive evidence of a massive black hole in the center of a galaxy.

In this section, the extragalactic masers (OH, H2O, formaldehyde, and methanol) will be discussed in the context of the locations in which they are found and of the impact of the SKA on their science. In conclusion, we consider in some detail the area in which the SKA may make its greatest in this field: the determination of the mass distribution of massive black holes in the nuclei of active galaxies. Megamasers (so-called because their luminosity is about a million times greater than a standard Galactic maser) have been seen in both OH, first discovered in Arp 220 by Baan et al. (1982), and H2O, first discovered in NGC 4945 by dos Santos and Lepine (1979). Only recently formaldehyde emission has also been found in a number of nearby (active) galaxies, whereas H2CO emission has only been seen at two locations within the Galaxy (Baan et al. 1993). A review of megamaser characteristics can be found in Henkel et al (1993). While H2O megamasers appear to provide information of the parsec-scale surroundings of the galactic nucleus, the OH and H2CO megamaser emissions provide a view of the physics and dynamics of the inner few hundred. Curiously, extragalactic methanol masers have not yet been found although they are found widespread as Galactic masers (Phillips et al. 1998). 2.6. CIRCUM-NUCLEAR MEGAMASERS 57

Figure 2.16: This graphic shows the geometrical relationships of the jet emission, the disk of water molecules and the black hole at the center of the galaxy NGC 4258. The pseudo-colors show the relative intensity of radio emission from the jets. The black dot indicates the location of the black hole. The dots in the disk indicate the location of water maser “spots” observed with the VLBA. All components are to scale; the scale bar indicates 5,000 Astronomical Units. (from Herrnstein et al. 1997)

2.6.1 H2O megamasers

Since the first H2O megamasers were discovered (Dos Santos & Lepine 1979), they were suspected to be associated with accretion discs around black holes. Using VLBI techniques, Miyoshi et al. (1995) show that the H2O masers in NGC 4258 are confined to a thin molecular disk, only 0.5 pc in diameter, surrounding a central engine. This data provides the best evidence to date for the existence of massive Black Holes (MBH) in some active galactic nuclei (AGN) of the Seyfert 2 type. As a result, H2O megamaser studies are one of the most powerful tools available to us for probing the inner parsecs of active galaxies. For example, the rotation curve of the maser source in NGC 4258, which is Keplerian to high precision, has provided a mass estimate, accurate to a few percent, of the central engine, a well-defined geometric model as shown in Fig. 2.6.1, and the opportunity to measure the 3-dimensional velocity field of gas in the core of an AGN, using the proper motions of the masers. Recently, other authors (e.g. Herrnstein et al. 1997) have used these and related results to examine the turn-on of the radio jet at a distance of a fraction of a parsec from the MBH, as predicted by the standard Blandford & Königl model. Potentially, H2O megamasers could be used to examine the relationship between MBHs and their host galaxies of various types, and as a function of their evolutionary state. However, we are severely limited at present by the small number of H2O megamasers, and by the even smaller number that can be successfully observed with VLBI techniques. The most complete survey for H2O megamasers among Seyfert 2 and similar galaxies has so far been done by Braatz et al. (1996, 1997), made with a 1-σ sensitivity of typically 30 mJy in a 0.8 km/s channel spacing. The total of 16 known megamasers 58 CHAPTER 2. FORMATION AND EVOLUTION OF GALAXIES

discovered by Braatz et al and other authors have ﬂuxes in the range 60-16000 mJy, corresponding to an isotropic luminosity of 24 to 6100 L⊙. A typical VLBI array, such as the VLBA, has an rms imaging sensitivity in 0.8 km/s bandwidth of 7 mJy/beam in 8h (if there is a suitable nearby phase reference), or a baseline sensitivity of 700 mJy/beam in 2 min (i.e. without phase referencing). In practice, observations so far have been made without phase referencing, so that only the strongest few sources have been imaged. Such VLBI observations are essential if we are to use H2O megamasers successfully to understand the processes surrounding the MBH.

2.6.2 OH Megamasers OH megamasers probe somewhat larger scales within the nuclear region of their host galaxies. The OH megamaser galaxies are part of a much larger population of (ultra- ) luminous FIR galaxies and the pumping of the OH molecules is done by the FIR radiation field in these galaxies (Baan 1989). The strongest OH megamasers have line strengths of several hundreds of mJy. The characteristics of OH megamasers suggest that the OH line luminosity increases quadratically with the FIR-luminosity, which implies that, as one samples larger volumes of space, one encounters increasingly luminous (but increasingly rare) “gigamasers” (Baan 1989; Staveley-Smith et al. 1989). This quadratic relation may in part be explained by maser amplification of (weak) radio continuum emission. Clearly the OH-luminosity function must turn over at some point as a result of the evolutionary paths of luminous FIR galaxies and of galaxy mergers, but we have not yet discovered that point. The fluxes of the OH megamasers lie in the range 6-300 mJy, corresponding to an isotropic luminosity 4 of 3 to 1.4 x 10 L⊙ and with the source at the highest redshift of 0.265 also having the highest luminosity (Baan et al. 1992). Recent global VLBI studies of the prototype OH megamaser in Arp 220 (type Seyfert 2) have shown that the nuclear radio continuum is dominated by powerful (unresolved) SNR’s spread over a region of order 50 parsec (Smith et al. 1998). On the other hand, the OH line emission originates in both compact and extended emission regions centered on these SNR clusters (Lonsdale et al. 1998). The greatest contribution of the SKA for OH research will be its ultimate sensitivity, its RFI robustness, its speed in searching for gigamasers in large volumes of space, and its use as one element of a VLBI array. The array itself would be able to resolve a 50 pc disk at a redshift of 0.04, with a sensitivity of 1 µJy in 8h, which allows imaging of all nearby megamasers.

2.6.3 Formaldehyde Megamasers Formaldehyde emission has only been found in nearby prominent FIR galaxies (Baan et al. 1993). The known H2CO emission lines are still rather weak and may be explained by maser ampliﬁcation of weak continuum emission. The exact pumping mechanism is still not clear but their occurrence appears correlated with the FIR 2.6. CIRCUM-NUCLEAR MEGAMASERS 59 galaxy population; some FIR sources that are not “warm” enough to be OH megamasers but they do show H2CO emission. The strongest formaldehyde emitter known to-date is Arp 220 with a line ﬂux density of only 4 mJy, while the luminosity of the known sources ranges from 5-200 L⊙. The emission in Arp 220 is located at the two nuclei as for the OH emission but it also closely mimics the peculiar NIR emission structure found with HST-NICMOS (Baan & Haschick 1995; Scoville et al. 1998). It is anticipated that many galaxies will show formaldehyde emission across their most active starburst/NIR/FIR regions providing a new diagnostic for these activity regions. The large sensitivity and its imaging capability will make SKA a unique instrument for H2CO research.

2.6.4 The Impact of the SKA on Megamaser Studies At present we are limited by the small number of known megamasers particularly those strong enough to be suitable for VLBI. Only with VLBI measurements can we unleash the full power of megamasers as a tool for understanding the nuclei of active galaxies. The SKA will contribute to molecular and megamaser research in three ways:

as a detection instrument it will be able to detect megamasers of all ﬂavors, • that are too weak to be detected with existing instruments. Assuming that the multi-beam capability and the high sensitivity allow routine phase referencing, the sensitivity of the SKA will enable detection (and imaging) of

– all currently known megamasers in less than a second, – currently unknown megamasers in nearby galaxies, down to masers of strengths comparable with interstellar masers in our Galaxy, – moderate luminosity OH megamasers comparable to Arp 220 at the 20 µJy level with 10 km/s resolution at up to a redshift of 2, and

– moderate luminosity H2O megamasers comparable to NGC 4258 at the 5 µJy level with 1 kms−1resolution up to a redshift of 0.15.

as a stand-alone instrument the SKA will be able to resolve the nearby H2O • megamasers similar to NGC 4258 with a total extent of some 10 milliarcsec and OH megamasers similar to Arp 220 with total extent of about 0.2 arcsec.

as an element of a VLBI array with existing antennas as the other elements, • and assuming that the multiple beams and greater sensitivity enable routine phase- referencing, a baseline between a VLBA antenna and the would enable high-resolution imaging of

– all currently known megamasers to a dynamic range of at least 300, – many currently unknown megamasers in nearby galaxies, and 60 CHAPTER 2. FORMATION AND EVOLUTION OF GALAXIES

– H2O megamasers comparable in luminosity to NGC 4258 at a redshift of 0.06, at which the maser disk is just resolvable by the longest Earth-based baseline, and similarly – OH megamasers comparable in luminosity and structure with Arp 220 at redshift 0.2.

For megamaser studies the effect of the SKA is both to vastly increase the number of objects to be studied, thereby increasing our knowledge of the workings of active nuclei with or without MBH’s, and to extend the volume of space to be studied. In this manner megamasers may be used as diagnostic tools for studying the evolution of galaxies in the Universe. As a particular example of H2O megamasers, most currently known sources occur in Seyfert 2 and LINER galaxies. Given sufficient sensitivity, one may expect to find them in other types of galaxies as well and thus be able to compare the occurrence and mass of the MBH’s as a function of galaxy type. Specifically, one should be able to make significant inroads on answering questions such as:

What is the mechanism for fuelling a black hole, and what are the kinematics • of the accretion disk?

What is the relationship between the mass of the black hole and the type and • evolutionary history of the host galaxy?

How do black hole masses vary as a function of redshift? Are the massive black • holes a result of many mergers of small black holes, or do small gas-rich galaxies already contain such black holes?

Can we see accretion disks around black holes in merging galaxies? If so, can we • trace the kinematics of the circum-nuclear material as the black holes merge?

2.7 The Starburst Phenomenon

Radio observations of supernova remnants (SNR), and the relativistic electrons produced by them, greatly increase our diagnostic capabilities for investigations of the starburst phenomenon in galaxies. These observations can give a direct measure of the supernova rate, the massive star formation rate and the high end initial mass function in these galaxies. In addition, through absorption measurements, the SNR can be used to probe the ionised and neutral component of the starburst ISM on parsec scales. A further aspect of this research is that each starburst acts as a laboratory for the statistical study of well deﬁned samples of SNR. The work is currently strongly limited by brightness sensitivity and hence only a few nearby starbursts have been studied in depth. The sensitivity problem is illustrated with reference to recent MERLIN and VLBI observations. It is concluded that 2.7. THE STARBURST PHENOMENON 61

SKA would not only greatly advance research into the physics of starbursts and supernova remnants, but would also enable star formation processes to be investigated in the early universe.

2.7.1 The importance of Starbursts A major aspect of understanding the evolution of galaxies is to be able to parametrise their star formation history. Several decades ago it was inferred that the Milky Way passed though an era of high star formation in its early history (e.g. Eggen, Lynden- Bell & Sandage 1962), and recent studies of the integrated ultra-violet(UV), optical and infra-red(IR) emission from field galaxies (Madau et al. 1998) have suggested that the universal star formation rate may have peaked at redshift of z 1.5. Galaxies ∼ with high star formation rates are still relatively common in the nearby universe, and when the star formation rate cannot be sustained for the lifetime of the galaxy the phenomenon is loosely classified as a ‘starburst’. Studies of nearby starburst galaxies can give a unique insight into the processes and causes of high star formation rates which is not readily accessible in galaxies with high star formation rates at high redshift. However direct studies of optical and UV emission from the stars in these nearby starbursts are often hampered by extinction from the molecular clouds associated with the star formation, and many of their properties are inferred from studies of IR emission from thermal re-radiation by dust which has been heated by starlight. Radio observations of starbursts are not affected by extinction and give independent and, in many cases, more precise measurements of both the star formation rate and the structure of the starburst.

2.7.2 Current Radio Studies Extended continuum emission

Radio continuum emission gives unique information on the star forming regions in nearby galaxies. Two main processes are responsible for the radio emission - thermal via the free-free mechanism and non-thermal via synchrotron mechanism, the former originating from the HII regions associated with young stars and the latter is mainly due to relativistic electrons generated by supernova events. For most decimeter studies the non thermal emission dominates, although free-free absorption becomes increasingly important at longer wave lengths (Wills et al. 1997). The existence of a tight corelation between the non-thermal radio and far infra-red luminosity in numerous galaxy samples(e.g. Helou et al. 1985) demonstrated that non-thermal radio emission was also a strong indicator of the star formation rate in galaxies. Condon and others (e.g. Condon 1992) have quantiﬁed this process and are able to relate the supernova rate (and hence an estimate of the star formation rate) to the total radio luminosity via relatively simple equations. However it is important to consider that most of the extended emission is produced by relativistic electrons from remnants 62 CHAPTER 2. FORMATION AND EVOLUTION OF GALAXIES

which have faded away long ago and, compared with studies of the actual remnants (see below), gives a relatively crude measurement of supernova rates.

Compact radio components The discovery (Unger et al. 1984, Kronberg et al. 1985, Antonucci & Ulvestad 1988) of large numbers of compact radio components in nearby starburst galaxies (Fig. 2.17) has enabled a much higher degree of precision to be brought to studies of supernovae in starburst galaxies. Many of these compact components have non-thermal spectra, high brightness temperatures and show parsec scale shell structures and hence most of these objects are young supernova remnants (Muxlow et al. 1994). In M82, for example, most of these components are more luminous and compact than Cas A, and hence they almost certainly have ages of only a few hundred years. This assertion has been dramatically confirmed by recent EVN measurements of expansion velocities in the M82 remnants (Fig. 2.18) consistent with the youngest remnant having an age of 30 years (Pedlar et al. 1998). Once the age of a sample of these remnants is calibrated,∼ then a direct measure of the supernova rate can be made by simply making high resolution radio images of these objects. In regions with high supernova rates, the rate can, in principle, be checked by observing the number of new remnants which appear over a few decades or even years. In moving from supernova rate to the rate of star-formation both the initial mass function and the cutoff below which stars do not form type II supernovae needs to be understood. Nevertheless the measurement of the luminosity function and spatial distribution of these SNR in a starburst galaxy can give an accurate measurement of the star formation rate and trace out the structure of star forming regions. An additional benefit of these observations is in the area of statistical studies of supernova remnants. Although a large number of remnants in our own galaxy have been detected and parameterised, most of the samples have uncertain distances and were observed with differing linear resolutions and sensitivities, often with different instruments. Green (1984) has outlined the selection effects and other problems with such a sample. However a single radio observation of a starburst galaxy can give large sample of SNR all at essentially at the same distance (e.g. 3200 kpc with a scatter of only 0.5 kpc within the starburst region) and observed with the same angular (hence linear) resolution and sensitivity. Studies are already underway to investigate the relation between size and luminosity, spectral index evolution as well as real time measurements of flux density decay and expansion rate. There is much work to be done in this area, and many of the simple models of radio supernovae and SNR need to be reconsidered. However perhaps the most important area to urgently investigate is the relation between ’normal’ SNR and the radio luminous ‘hypernovae’ of which 41.95+575 in M82 is the best studied example (e.g. Wilkinson & deBruyn 1990). The recent discovery ( Smith et al. 1998) of a number of compact objects in Arp220 at a distance of 75Mpc, with a similar luminosity to 41.95+575, is particularly exciting, as, if we can relate these objects directly to the supernova/star formation rate, it will be possible to measure supernova rates at 2.7. THE STARBURST PHENOMENON 63

Figure 2.17: A λ6cm MERLIN/VLA image of nearby starburst galaxy M82. The discrete sources are mostly supernova remnants with ages less than 1000 years and compact HII regions. The non-thermal extended background is mainly due to relativistic electrons generated by older remnants. cosmological distances with the SKA.

Line studies

An increasing number of starburst galaxies have been the subject of neutral hydrogen emission studies. However the brightness temperature limitations of current instruments prevents these measurements having higher resolution than 5 arcsec which corresponds to size scales > kpc in all but the nearest starburst∼ galaxies. How- ever, high angular resolution neutral hydrogen measurements can be achieved via absorption measurements, and the absorption spectrum against individual supernova remnants can be used to probe the interstellar medium of the starburst host on parsec size scales. In a recent study of M82, Wills et al. (1998) measured H I absorption against 33 remnants and investigated non-circular gas motions associated with the starburst. Hydroxyl masers as well as molecular absorption have been observed in nearby starbursts, although much of this work is currently limited by sensitivity. A small number of starbursts, however, contain ‘megamasers, and can be observed to cosmological distances even with current instruments. 64 CHAPTER 2. FORMATION AND EVOLUTION OF GALAXIES

Figure 2.18: An example of a shell supernova remnant seen in M82 by the EVN. The remnant was observed at two epochs and has expanded consistent with an age of 30 years (Pedlar et al. 1998). ∼

2.7.3 The Potential of SKA for Starburst Studies Nearby starbursts Although important advances have recently been made in radio continuum studies of nearby starbursts, almost all the work is severely limited by sensitivity. In principle we have sufficient angular resolution ( via MERLIN, EVN, VLBA etc) to resolve most remnants in starbursts out to 100Mpc (1mas 0.5pc at 100Mpc). The major limitation is, of course, the dramatic decrease in brightne∼ ss sensitivity which is a natural consequence of obtaining high angular resolution with comparable collecting areas. This is evident in the EVN observation of M82, which at 15mas resolution, only detected 5 out of the 50 objects detected by MERLIN/VLA despite having ∼ comparable flux sensitivity. Although there is considerable dispersion, it appears that the total flux density of the remnants in M82 decreases approximately inversely as the diameter (Muxlow et al. 1994)

−1 S5GHz = 30Dpc mJy and hence the average brightness of an SNR decreases as D−3 which even in M82, results in the more extended (> 4pc) remnants escaping detection by MERLIN. It is indeed fortunate for present studies that we have two relatively strong starbursts 2.7. THE STARBURST PHENOMENON 65

in M82 and NGC253 at 3Mpc. These two are the most luminous objects within 15Mpc and hence studies∼ of other starbursts are severely limited either by being at greater distances or having lower supernova rates, and hence MERLIN observations of a number of nearby starbursts have only detected relatively small numbers of remnants. SKA with sub µJy sensitivity would not only enable much deeper studies to be made of older, more extended remnants in nearby strong starbursts such as M82 and NGC253, but would also result in the detection of statistically usable numbers of SNR in many nearby lower luminosity starbursts. The ability of SKA to image starburst galaxies at low radio frequencies can also be used to constrain the distribution of ionised gas in the object via free-free absorption studies. A recent subarcsecond study of M82 at 73cm (Wills et al. 1997) has revealed extensive regions of ionised gas which are completely inaccessible in the optical because of extinction. Finally SKA will enable subarcsecond neutral hydrogen emission studies of nearby starbursts to be made which will complement the current absorption studies and determine the gas dynamics and structures of neutral material associated with the starburst. In addition, by comparing emission and absorption spectra against individual SNR it will be possible to measure the spin temperature of neutral gas in the starburst.

Distant Starbursts Even the weakest remnants currently seen in M82 would be detectable by SKA at 100Mpc and hence supernova remnants in hundreds of starburst (and ‘normal’) ∼ galaxies could be studied statistically. This would enable the effect of different environments on supernova remnant evolution to be investigated, as well as constraining the physics of the star formation process. The stronger remnants, particularly ‘hypernovae’ such as 41.9+58 in M82 and possibly the compact objects in ARP220, would be detectable by SKA at cosmological distances ( 1000Mpc), and assuming ∼ that these brighter objects can be related to the normal supernova remnants via a luminosity function, it will be possible to measure supernova and star formation rates in the early universe. One important parameter for these studies is, of course, the angular resolution necessary to separate individual SNR in a starburst. Assuming a typical separation of 10pc then an angular resolution of 10mas is required at 100Mpc. This would, of course, require baselines of several∼ thousand kilometers, and hence SKA would need to be linked either to the existing VLBI networks or a number of purpose built SKA elements need to be situated at these spacings. Even if SKA is just used as an additional element in existing VLBI networks, then the resulting sensitivity of 1µJy/beam would detect statistically useable numbers of SNR in all starburst galaxies∼ to 100Mpc. ∼ In addition, as we have discussed in section 2.7.1, the diffuse radio emission from a starburst can be related to the star formation rate (Condon 1992). Such emission would be easily detectable by SKA. For example an object such as M82 would have a 66 CHAPTER 2. FORMATION AND EVOLUTION OF GALAXIES total 20cm flux density of 100 µJy at 1000Mpc. Indeed several of the objects already seen in 20cm MERLIN/VLA∼ observations of the Hubble deep field are extended and are very likely starburst galaxies (Muxlow Private Communication). Hence it is clear that SKA would enable radio studies of starburst galaxies to be easily carried out at cosmological distances. Unlike radio loud objects, which only constitute a small percentage of all galaxies and also have a wide dispersion in intrinsic properties, starbursts are present in a large fraction of galaxies and appear to form a relatively uniform group. Such studies would be complementary to studies at other wavebands, but would not require complex modeling to account for uncertainties in extinction etc. Hence it seems likely that, with the sensitivity of SKA, starbursts can be used to constrain the parameters of the universe at cosmological distances with high precision.

2.8 Interstellar Processes

The interstellar medium is the matrix within which the processes of galaxy evolution occur. It exists in many states: atomic and ionized hydrogen, relativistic plasma, molecular gas, and dust, each containing velocity and density structures over a vast range of scales. At the low spatial scales, a highly disturbed state is maintained by point-like energy input from stars at all phases of the stellar life cycle. On the other end of the scale, energy input can take the form of global, large-scale phenomena, such as viscous dissipation or magnetic stress from Galactic rotation, and the motion of spiral arm density waves. Despite the apparent flux of energy on all scales, pockets of relative quiescence exist, where cold gas can self-gravitate and the process of star formation begins. The life cycle of stars, and the state and evolution of the ISM environment, are intimately intertwined. To date we are restricted to studying these processes in detail only in our own and a very few nearby galaxies. While several large ambitious projects are now underway to take advantage of the unique perspective we have in our own Galaxy, we have essentially no detailed information on the evolution of the ISM of galaxies from the epoch of formation, through the multiple stages of star formation and recycling to the present epoch. A multi-wavelength approach is required to fully understand the complex processes and phenomena that govern the evolution of the interstellar medium. To extend these studies back in cosmic time we must be able to image the ISM of external galaxies to high redshift in as many states of the ISM as possible. The SKA offers the possibility of studying the interstellar medium of a significant number of external galaxies with a detail that has heretofore been possible only in our own and nearby galaxies. The SKA will allow imaging of four major components; the atomic hydrogen gas, the relativistic plasma, the ionized medium and the molecular medium. 2.8. INTERSTELLAR PROCESSES 67

2.8.1 HII Regions: High Resolution Imaging of Thermal Emis- sion The physics of heating and cooling in a photoionized plasma of characteristic astrophysical abundance results in an equilibrium kinetic temperature of a few 104 K. The photoionized interstellar medium, thus has brightness temperatures of this value or lower, depending on optical depth. The sensitivity of the SKA will open for the first time the possibility of imaging of this low surface brightness, thermal radio emission at milli-arcsecond resolution. This capability will have a revolutionary impact on the field of radio astronomy, one that will spill over into many areas of astrophysical inquiry. This advance is illustrated in Fig. 2.19, which shows a plot of angular size versus brightness temperature. The three dashed diagonal lines show the angular radius of a source required to produce flux densities at λ6 cm of 100 mJy, 1 mJy and 1 µJy as a function of the brightness temperature of the source. The dark solid lines characterize the imaging capabilities of the most powerful existing radio telescope arrays. The horizontal portion shows their maximum resolution, and the diagonal portion shows the minimum detectable flux density. For a radio source with a given angular radius and brightness temperature to be resolved by a particular radio telescope array, it must lie above the solid line. The current suite of most sensitive radio telescope facilities occupies the upper right portion of this diagram, and is able to image only sources with either very high brightness temperature or large angular size. For thermal radio sources (below 104 K) the maximum attainable resolution is 0.1′′ (with the VLA). To make inroads toward the small angular diameter, thermal brightness∼ temperature region, a giant step in sensitivity is needed. The SKA will uniquely occupy this region of parameter space, providing angular resolution of 10 milli- arcseconds at λ6cm down to continuum brightness temperatures well below 100K, and, as an element of a global VLBI array, a few milli-arcsecond resolution at temperatures of 103 - 104 K. Ionized hydrogen in external galaxies is a direct tracer of massive star formation. At a wavelength of a few cm, the SKA would be able to detect the continuum emission from the Strömgren sphere surrounding an O5 star at a distance of 100 Mpc in 12 hours. A B0 star would be detected out to 10 Mpc. An HII region surrounding a luminous early-type star has typical dimension of tens of pc. More compact and ultracompact HII regions are known in our own Galaxy with dimensions 1 to 0.1 pc and brightness temperature 103 - 104 K. Since in the local Universe, the surface brightness of resolved objects is independent of distance, the SKA will be able to image bright HII regions at linear resolution below 0.1 pc at up to 20 Mpc distance, and 0.5 pc at 100 Mpc. Such observations would provide, for example, complete counts of luminous young stars in galaxies well beyond the Virgo cluster, providing direct measurement of the high end Initial Mass Function in a very large number of galaxies and in a range of cluster environments. These measurements would be unaffected by extinction. Comparison with Hα images from sensitive optical telescopes would yield measurements of the extinction and allow derivation of temperatures and densities. The deconvolution of non-thermal and thermal emission in galaxy disks from high 68 CHAPTER 2. FORMATION AND EVOLUTION OF GALAXIES

Figure 2.19: The area of brightness temperature – angular size space that will be opened up by the SKA. Photoionized hydrogen gas at a temperature of 104 K and below will be imaged at resolutions of a few milli-arcseconds. Compact HII regions in external galaxies could be images with 0.5 pc linear resolution at 100 Mpc distance. Also shown on the plot, for reference, are the temperatures and angular radii for stars in the supergiant branch at a distance of 100 pc. 2.8. INTERSTELLAR PROCESSES 69

resolution imaging at decimetre and centimetre wavelengths will provide the means to investigate the origin of the far infrared-radio luminosity correlation in galaxies. By comparing such images to high resolution images of dust emission obtained from large submillimetre arrays, we will be able make detailed spatial studies and determine the dominant dust heating mechanisms. Observations of HII regions at low frequencies (330 MHz) have demonstrated that measurements of optically thick HII regions can constrain source electron temperatures, emission measures, and ﬁlling factors (Kassim et al. 1989; Subrahmanyan & Goss 1996). At even lower frequencies these regions appear as cooler regions against a much hotter Galactic background, allowing kinematic distance ambiguities within our Galaxy to be resolved and the superposition of thermal and nonthermal sources to be separated along complex lines of sight through the Galaxy. Kinematic distance ambiguities resulting from radio recombination line measurements can be resolved using the detection, or non-detection, of HII regions in absorption below 100 MHz (Kassim et al. 1990). This is because foreground HII regions would be much more prominent absorption features on low frequency SKA maps than distant ones.

2.8.2 Centimetre Wavelength Molecular Probes of the ISM Diffuse Molecular Lines The SKA will also be a powerful probe of the molecular interstellar medium, which will be an important complement to CO line studies. The cm wavelength transitions of molecular lines naturally sample the low-temperature environment that characterizes much of the volume occupied by the molecular ISM. The wavelength range longer than about 1 cm favours molecular hyperfine and fine structure, λ-doubling, and rotational transitions of large molecules. The physical conditions in this gas are conducive to narrow lines that tend not to overlap, and are very useful probes of turbulent motions of only a few km s−1. Designed with these observations in mind, the resolution and sensitivity of the SKA will open up the field, enabling large scale sampling of important tracers as well as new extragalactic comparisons with Galactic observations. The molecules H2CO, OH, NH3, and the H2O maser, have strong cm wave lines that have been studied extensively. H2CO (λ6 cm) is seen in absorption against the 2.8K microwave background, thus allowing unrestricted mapping. OH (λ18 cm) is seen in both emission and in absorption in more diffuse gas. With resolutions of a few arcsec, the SKA will be able to trace the structure and dynamics of the diffuse molecular gas (OH) as well as the denser gas (H2CO) in most nearby galaxies, and in galaxies at modest redshift. The OH molecule is highly important in astrochem- istry. It is key to primordial chemistry since it participates in the very first chemical reactions that form molecular hydrogen. NH3 is one of the most important interstellar molecules, and its large suite of lines near λ 1.3 cm are powerful probes of the conditions in dense (> 104 cm−3) molecular gas (see Avery 1991), the sites of star formation. Since these lines are strong, they could be observed in galaxies out 70 CHAPTER 2. FORMATION AND EVOLUTION OF GALAXIES

to several 100 Mpc at resolution of 10 milli-arcseconds, enabling the study of the density of protostellar sites at sub-parsec∼ resolution in a large number of galaxies. There are several tens of other molecular species known in the cm wavelength range, each with many lines. Most have been detected by sensitive single dish observations, but their applicability to a broad range of interstellar physics has not been fully explored because the lines are too weak, and single dish resolutions are insufficient to allow full comparisons with other tracers of the interstellar medium. A prominent example is CH, an important constituent in chemical networks, which has weak lines near λ10 cm. The combined sensitivity and resolution of the SKA will provide a major impetus to this field and spark the application of weak line observations to many new problems in ISM astrophysics. Brightness temperature sensitivity is important for detecting and mapping weak lines, and so the array configuration of the SKA will play an important role. For these types of observations the compact central core of the array (containing a large fraction of the total collecting area) would be used. It is likely that the array could be designed to include about 30% of the collecting area in baseline spacings less than 2000 m. The brightness temperature sensitivity for this array would be about 10 mK in an observing time of 24 hours, and the resolution would be about 10” at λ10 cm, allowing imaging of large molecular clouds in nearby galaxies. A number of molecules emit spectral lines of astrophysical importance in the 20 - 50 GHz range, and could be imaged with the SKA at λ1.3 cm from galaxies at z 1 and greater. These include the lower transitions of CS, SO and HC N. CS and ∼ 3 HC3N have been detected in nearby galaxies at levels of several tens of mK. At λ1.3 cm, the compact core of the array would have 10 mK sensitivity and resolution of 1′′. Detection of emission from these molecules at high redshift would complement CO studies, and provide important information on the physical conditions of molecular gas and on chemical evolution at early epochs.

Interstellar Masers

OH, H2O, and methanol masers are frequently found in the cool molecular gas surrounding newly formed stars, and have been used extensively to probe the physical conditions and kinematics of this gas. Numerous observations have shown them to consist of clusters of small maser components often associated with compact 12 HII regions. The individual maser components have typical sizes of 10 m (corresponding to 0.01 arcsec at 1 kpc), and are arranged in clusters typically 3.1014 – 1015 m ( a few arcsec) in diameter, although there are significant differences between the∼ three species. Most of these masers appear to be situated in the warm gas accreting on to a pre-main sequence massive star. The properties of OH and H2O masers are summarised by Elitzur (1992). A recent development is that about 30% of methanol masers (at 6.7 and 12.2 GHz) appear to be located in edge-on circumstellar disks (Norris et al., 1993, 1998; Stecklum et al, 1998) around high-mass stars. The existence of these disks will require revision of those theories which as- sert that such disks should be destroyed by the strong stellar winds from high-mass 2.8. INTERSTELLAR PROCESSES 71 stars. Interstellar masers tend to be so strong that present-day radio-telescopes are quite adequate for studying those in our Galaxy. However, the enormous sensitivity of the SKA opens up the possibility of using these masers to probe extragalactic star formation. At present it is difficult to study star formation in the nuclei of active and starburst galaxies because the high dust extinction (which can be hundreds of magnitudes at optical wavelengths) prevent even mid-infrared observations from pen- etrating the dense shroud of dust. Whilst the far-infrared observations of ISO can penetrate this, ISO does not have sufficient angular resolution to measure the distribution of star formation on the parsec scale in the nuclei of these galaxies. Such resolution is important if we are to understand, for example, the potential role of star formation in feeding the massive black hole (MBH) at the centre.

At the distance of NGC253, individual methanol and H2O maser spots corresponding to our own normal galactic masers might have observed flux densities of 3 and 250 mJy respectively, while at Cen A they might have fluxes of 0.4 and 30 mJy. In both cases, individual maser spots would be detectable at the five-sigma level in a few minutes, and could be identified in the parsec-resolution images obtained from an 8 hour synthesis. SKA would therefore tell us about the structure and kinematics of star formation regions in nearby active and starburst galaxies at a level of detail approaching that which we have in our own galaxy. This would help enormously in solving questions such as the relationship between star formation and the AGN.

2.8.3 Supernova Remnants

Supernovae and their remnants play a central part in the dynamics and evolution of the Galaxy. Supernovae inject massive amounts of energy into the interstellar medium (ISM), powering a large fraction of the turbulent motions seen there. The expansion of the supernova remnant (SNR) both illuminates pre-existing structures in the ISM and carves out new structures, transferring kinetic energy from the original supernova to the ISM. Supernovae act as recycling centers, taking material that would be otherwise trapped in stellar form and returning it to the ISM to form new stars. This recycled material is processed by the supernova explosion into iron or heavier elements, those which profoundly eﬀect the energetics of the ISM and the next generation of stars which form out of it. SNRs are central to other areas of astronomy also. SNRs are thought to be the sites of cosmic ray production, but the exact mechanism is ill-understood, and the connection between the two as-yet unproven. Supernovae associated with the death of massive stars are expected to leave behind stellar, as well as diﬀuse, remnants. These stellar remnants usually take the form of neutron stars, but some fraction are also expected to form black-holes, possibly the most enigmatic objects in the Universe. The interaction between SNRs and dense molecular clouds drives chemical reactions that cannot be duplicated on earth. As such, these interaction regions serve as invaluable laboratories, sampling chemistry which can only take place in extreme environments and situations. 72 CHAPTER 2. FORMATION AND EVOLUTION OF GALAXIES

Yet, despite the importance of supernovae and SNRs in our understanding of our Galaxy, our knowledge about the SNRs and their physics is far from complete. Exactly where are cosmic rays accelerated, and how? How do the supernova ejecta couple to the ISM? Where are the neutron stars (and black holes?) in SNRs? Where are the young SNRs? Where are the oldest? These are just some of the questions that new radio telescopes can help us answer. Radio wavelength observations of SNRs provide an important complement to studies at other wavelengths. Optical observations of SNRs provide important information on thermal particles and abundances in the SNR through recombination line emission. Infrared emission traces the shocked dust within the SNR, as well as dust trapped within the hot interior of the SNR. X-ray emission arrives from the region of hottest gas in the SNR, and provides invaluable information on the temperature structure of the SNR through high-ionization state line emission. To this impressive arsenal is added radio wavelength observations of synchrotron emission, tracing the magnetic ﬁeld and relativistic particle distribution within the SNR, thus providing important information about both the SNR shock and the ISM. The SKA promises to revolutionize the way we look at SNRs in the radio regime and, as a result, our understanding of these important objects. The possibilities of- fered by the SKA are many, and their inter-relation complex. Instead of concentrating on the new observational capabilities oﬀered by the SKA, we will illustrate the way that the SKA will be able to illuminate some of the important questions in SNR astronomy.

Where are all the young SNRs? Studies of supernova rates in external galaxies suggest that we should expect roughly 2 per century in our own. Despite this, we are only aware of a handful of Galactic SNRs younger than about 1000 years. The lack of young Galactic SNRs can be partially explained in two ways. First, the lack of an obvious optical supernova over the last few centuries, if not simply a statistical ﬂuke, suggests that they must have been heavily obscured, and thus probably very distant. Second, the emission from a SNR is a direct result of its interaction with the ISM; the lack of obvious emission from young SNRs could thus suggest they occurred in low density regions of the ISM and thus are faint. Combined, these explanations suggest that future searches for young SNRs require high sensitivity (to detect the faint objects) and high resolution (to resolve the distant, and thus angularly small objects). These requirements are exactly where the SKA excels.

Old Supernova Remnants The explosion of supernovae is a dynamic and exciting event and study of young supernova remnants can often tell us valuable information about the progenitor star and its circumstellar medium. Understanding of the ﬁnal merging of the ejected material with its surroundings and the eventual dissolution of the energy throughout the in- 2.8. INTERSTELLAR PROCESSES 73

terstellar medium requires study of the old SNRs. These old SNRs are generally very large and faint so they are difficult to study with current instruments; observations suggest that objects older than about 50,000 years are not currently radio detectable. The SKA with its high sensitivity and its capability to image large fields of view will be an ideal instrument for studying older objects. Most old SNRs have 80-90% of their total emission in a smooth component, which is usually missed with current aperture synthesis telescopes, and the remainder is in thin, often unresolved filaments. Both components need to be imaged to fully understand how the shocks decay as the remnant’s expansion slows down and the emitting material diffuses into the galactic background. Many large remnants, such as the nearby Cygnus Loop, cover several degrees on the sky and show a variety of features in different regions. In some places clumps of material are being overrun and there are characteristic changes in the position of the optical line radiation as the shock progresses into the clump. In other places the shock appears to be well in front of the expanding material and most of the emission appears to be from dense cool regions which are compressed under pressure equilibrium with the warmer interclump gas. Does the synchrotron radiation have different spectral signatures in these different regions which can be used to help determine their physical characteristics? To date, resolution of individual filaments has been limited by sensitivity so that the thinest radio filaments detected are limited to over 0.01 pc whereas the HST can approach a resolution of 0.0005 pc for the closest remnants. It is important to match the optical resolution at radio wavelengths in order to see how the relativistic particles track the thermal ones and how the energy in different components changes with the size and structure of different components. For example, does the spectrum change over a shock as it progresses into a clump? The SKA will have the sensitivity to allow use of its full resolution for answering questions like these.

Particle Acceleration in SNRs The synchrotron emission detected by radio telescopes is powered by the SNR blast- wave. The connection between the two is very unclear however. Where are particles accelerated in the shock? How high of energies can they be accelerated to? What role does the magnetic ﬁeld play? How important is turbulence? The SKA promises to help us understand these questions by allowing us unprecedented details of the spectral index variations across SNRs. The radio spectral index provides information on the particle acceleration mechanism and the underlying seed particle population. The variations across a remnant will allow us to trace how these change with ISM density, shock velocity, magnetic ﬁeld strength, and past history of the region. One of the most exciting ways the SKA will allow us to study these questions is by simply extending the useful wavelength range we can study. By being able to observe at low frequencies the uncertainties in our determinations of the spectral index will be reduced by as much as a factor of two. The high resolution of the SKA will allow us to measure the spectral index of small regions, making it possible to measure these 74 CHAPTER 2. FORMATION AND EVOLUTION OF GALAXIES

quantities on a scale similar e.g. to the Hubble space telescope. The high sensitivity and resolution will combine to allow us to see how these values change with time, a technique that has yielded amazing insights into the few objects, such as Cas A, the oungest Galactic SNR known, we are currently able to study in this way. Recent centimeter wavelength studies of Cas A (Anderson & Rudnik 1996) have suggested that the spectrum of the emitting regions may be determined not only by current acceleration processes but also by the history of particle acceleration in the environment through which the particles have moved. Observations with the 74 and 330 MHz VLA systems have recently confirmed this surprising conclusion. Perhaps even more exciting has been the unique absorption measurements provided by the 74 MHz observations which reveal evidence for unshocked ejecta within Cas A, as predicted by theory (Kassim et al. 1995). This measurement raises the prospect that many young supernova remnants may harbor a cool thermal core, which the low frequency measurements can uniquely detect. The detection of thermal absorption from within the first two SNRs observed with the new VLA system suggests that these effects may be common at low frequencies. Useful information can be gained simply from the integrated low frequency spectrum of SNRs. Predictions from Fermi acceleration theory imply concave integrated spectra such as has been claimed in the case of the Tycho and Kepler’s SNRs (Reynolds & Ellison 1992). But large error bars on the lowest frequency measurements hamper these conclusions and restrict their extension to many more sources. More accurate, higher resolution measurements can extend such studies to many more objects and confirm whether the line-of-sight thermal absorption is indeed related to envelopes of normal HII regions as is currently speculated. The low frequency abilities of the SKA will provide unique tests to theory. Fermi acceleration theory, for example, predicts a concave integrated spectrum at low frequencies. Current observations are not sensitive enough to reliably test these theories, but the SKA will be.

Dynamics of SNRs Perhaps the most basic information needed for understanding a SNR is the dynamics of the remnant. Exactly how fast is the remnant expanding? How does this vary across the remnant? How does it correlate with ISM density? How does it correlate with spectral index? What dynamical stage is the SNR itself in? What is the distance to the SNR? We will be able to address these questions with the SKA. The high resolution and sensitivity of the SKA will make it possible to measure proper motions for SNRs out to the other side of the Galaxy. As an example, a young SNR expanding at 104 km/s at a distance of 20 kpc would grow in diameter by 0.21” over one year, a growth which would be easily measurable using the SKA at its full resolution. A SNR expanding at 200 km/s at a distance of 4 kpc would increase its diameter by the same amount. These examples suggest that the SKA will provide unprecedented information on the expansion of SNRs. The proper motion is, of course, dependent on both the expansion velocity of the 2.8. INTERSTELLAR PROCESSES 75

remnant, and its distance. As a result, proper motion alone cannot uniquely constrain either of these quantities. Combining the proper motion with other information including radial velocities from optical spectroscopy will, however, shed light on these quantities. The sensitivity of the SKA will allow unprecedented measurements of high-velocity, post-shock, HI. The velocity of this gas will provide a lower limit to the present shock velocity, and thus – in combination with the proper motion information – provide a lower limit on the distance to the remnant. Conversely, if other information is available on the distance to the remnant e.g. associated pulsar dispersion measure, HI absorption, interaction with molecular material, the shock velocity will be able to be determined. The high spatial resolution of the SKA will allow the expansion rates to be mapped across the remnant. This will allow the study of e.g. SNR blow-outs into lower density ISM, the expansion of fragments vs the blast-wave as a whole, and the interaction of the shock with high density, molecular material. The interaction with molecular material can then be used to help constrain models of shock chemistry and dynamics.

SNR/ISM Interaction

The near-instantaneous injection of energy into the ISM by a supernova has profound effects on the surrounding medium. The material immediately around the explosion is swept into a fast-moving shell of dense material. Small dense clouds are destroyed, and shocks are driven into larger clouds, initiating chemical reactions impossible to find elsewhere. Left behind are low density bubbles of million-degree gas. The shells of material swept-up by the expanding blast-wave eventually cool and compress further to create shells tens of parsecs in diameter. These bubbles and shells can merge with similar bubbles and shells created by stellar winds and other nearby supernovae to form superbubbles hundreds of parsecs in diameter. The interaction of the SNR with the ISM can be seen in a number of ways. High- velocity H I , from recombined material behind the shock, has been used to identify shocks interacting with high-density material. Detection of this material gives an insight to the shock velocity. This material is very faint, however, requiring high surface-brightness sensitivity to detect; with its large collecting area, the SKA should be able to detect high-velocity, post-shock, HI in a large number of SNRs. The 1720 MHz OH maser line has recently been shown to be a powerful indicator of interactions between SNRs and dense molecular material. Follow-up observations of sites of OH maser emission have begun to reveal new sites of shock-induced interstellar chemistry, promising to revolutionize our understanding of this subject. Maser emission has also been used to measure the magnetic fields within the interaction zone, casting light upon this ill-understood topic. Present observations can detect only strong maser emission, but the SKA – with its increased sensitivity – promises to detect much weaker emission, increasing our ability to sample the interaction between SNRs and their surroundings. 76 CHAPTER 2. FORMATION AND EVOLUTION OF GALAXIES

SNRs and Pulsars The stellar remnants created in core-collapse supernova, in addition to being fascinating in their own right, can provide invaluable information about the associated diffuse supernova. Ages of SNRs, for example, are notoriously hard to come by unless the object is the result of an historical explosion. If an associated pulsar is found, its spin-down age can be used to estimate the age of the SNR. Additionally, although somewhat unreliable, the pulsars dispersion measure distance can be used to estimate the distance to the remnant. As a result, discovering pulsars associated with SNRs is of considerable interest. As noted elsewhere in these proceedings, the SKA will excel at finding new pulsars, some fraction of which will be associated with SNRs. Even without the detection of pulses from the pulsar, however, it is possible to infer the existence of a neutron star within a remnant. The synchrotron nebulae associated with Composite and Filled-Center SNRs are expected to harbor central rotating neutron stars, even though they may not be beamed towards us, and thus the detection of a synchrotron nebula within a SNR is enough to deduce the presence of a neutron star (NS). These pulsar powered nebulae can reveal their presence as hard X- ray sources, the spectra of which can be modeled to include a hydrogen column density from which a distance to the NS, and thus SNR, can be estimated. The sensitivity and resolution of the SKA will combine to allow researchers to detect more and fainter synchrotron nebulae, in currently identified as well as newly discovered SNRs. The interaction of the pulsar/NS with the SNR can also be illuminating. The Crab pulsar, for example, is associated with nearby, slightly elongated, features known as “wisps.” These wisps are thought to be the termination shocks of the pulsars free- flowing relativistic wind; the distance of this shock from the pulsar depends on both the energy of the wind and the physical conditions within the remnant, thus providing a probe of both. A similar wisp is seen adjacent to a hard X-ray point source in the Filled-Center SNR 3C58; this suggests that these wisps may be common within SNRs with pulsars. The resolution and sensitivity of the SKA will make it possible to detect similar wisps in more distant SNRs, and thus teach us more about the relationship between pulsars and SNRs. Finally, the SKA may be able to finally put to rest one of the most enduring problems in our understanding of pulsar-powered SNRs: the lack of a limb-brightened shell associated with the Crab Nebula and similar remnants. Despite being the first nebula to be conclusively associated with a historic supernova explosion, the Crab Nebula is among a handful of remnants which do not have the signature limb-brightened shell of a SNR. While expectations abound that the visible nebula is surrounded by an, as yet, undetected shell, speculation will run rampant until either a shell is detected, or sufficiently strong limits are placed on any emission from a shell. The high surface-brightness sensitivity of the SKA will be able to shed considerable light on this question and will, no doubt, be among the most eagerly-awaited first results from the SKA. The intrinsic surface-brightness sensitivity will be aided by the low-frequency abilities of the SKA, pushing sensitivity limits on a non-thermal shell even further. Supernova remnants, extended nonthermal emitting sources which are the princi- 2.8. INTERSTELLAR PROCESSES 77 pal source of energy input into the ISM, are natural targets for study with the SKA. Moreover, their often large angular size is well matched to the SKA large fields of view. High resolution, multi-frequency images will serve to anchor spectral index studies of SNRs whose spatially resolved continuum spectra uniquely constrain the energy distributions of relativistic electrons. The key is to relate measured source spectral variations to dynamical structure, since models of particle acceleration in SNRs, either by shocks or by second-order Fermi (stochastic) acceleration in interior turbulence, predict structure in the particle distributions. Measured variations must be related to acceleration processes or the injection spectrum of the seed particles. Variations in older SNRs can also be related to compression of cosmic ray gas and interstellar magnetic fields. Previous studies have had far too poor angular resolution at the lowest frequencies to explore such issues in detail, if at all. Sensitive low frequency observations should lead to the discovery of older, low surface brightness SNRs which are known to be missing from catalogs due to severe selection effects (Green 1991). Discovery of such older SNRs, at the last stage of evolution before blending into the ISM, are potentially of great importance in discovering new pulsar-SNR associations and drawing links to unidentified γ-ray sources. Presently only the youngest pulsars can be associated with SNRs since remnants older than about 1000 yr have surface brightnesses too low for detection by current instruments. The sensitivity and angular resolution of the SKA would be sufficient to extend these SNR studies to nearby external galaxies, thus greatly extending the available data base. Recent statistical studies (e.g., birthrates, distribution, energetics, etc.) of complete, co-distant samples of SNRs in nearby galaxies are proving extremely useful for exploring problems in stellar evolution, ISM structure, and for increasing samples sizes of poorly understood SNR-subclasses (Wills et al. 1997; Duric et al. 1995; Jones et al. 1998). Sensitive, high resolution low frequency observations are required to compliment existing higher frequency data and to anchor the derived spectra and search for absorption effects. VLA and Westerbork observations at 330 MHz and MERLIN observations at 151 MHz have been utilized successfully for these purposes, but lower frequency observations would be even more useful.

2.8.4 The Origin of Cosmic Rays The origin of cosmic rays has been a challenge ever since their discovery. The current paradigm holds that high energy phenomena, related to supernovae and/or active galactic nuclei (AGNs), are involved. However, no direct connection between the particles that we observe locally and any identified cosmic sources has been made, leaving their origin uncertain. A key barrier is observational. Because they are charged and deflected by Galac- tic, interplanetary, and geophysical fields, it is impossible to deduce the origin and complete spectrum of the cosmic ray particles from direct measurements, with the exception of the very highest energy particles. These appear to be extragalactic, but 78 CHAPTER 2. FORMATION AND EVOLUTION OF GALAXIES

are so few in number that good source statistics have not been obtained. Fortunately, cosmic ray particles generate radiation at both the highest (γ-rays) and lowest (radio) frequencies from their interaction with interstellar matter and magnetic fields. Hence interpretation of these radiation measurements may hold the key to unlocking the origin problem. However, while high energy capabilities have advanced quickly and have produced important results such as the ASCA X-ray evidence for cosmic ray production in the shell of SN1062, observational capabilities at the lowest frequencies remain primitive. A new instrument with fundamentally improved characteristics in angular resolution and sensitivity can make a major impact which, when interpreted within the context of modern high energy observations, may hold the answer to this puzzle. The low- and high-energy observations are related because the distributed γ-ray and the low frequency radio emission are both generated by cosmic rays (Longair 1990; Webber 1990). The γ-rays and radio waves are coupled through the distribution of interstellar hydrogen and magnetic fields, respectively, and their nonthermal character reflects the energy signature of the poorly understood cosmic ray “source” spectrum. The higher energy γ-rays (E > 100 MeV) result mainly from the pion decay that results from the collision of high energy (E > 300 MeV) cosmic ray nuclei with hydrogen. However lower energy γ-rays (E < 70 MeV) result mainly from relativistic bremsstrahlung of cosmic ray electrons of energies below 200 MeV in interstellar matter, and these particles also generate synchrotron radio radiation below 100 MHz. Hence a comparison between the low energy γ-ray and low frequency radio spectra could, in principle, allow us to uniquely separate the matter distribution from the magnetic field distribution, and to deduce the distribution and primary energy spectrum of the cosmic ray particles. The problem can be approached at the radio end in two ways. One can measure the distributed emission directly; however single-dish measurements have always had too poor angular resolution to properly deconvolve the true distributed emission from the myriad of discrete sources which pile up along the most interesting lines of sight, making comparisons with potential cosmic ray source distributions difficult. A more direct approach is to use an interferometer to resolve optically thick HII regions against which the distributed synchrotron emissivity could be accurately determined (Kassim 1990). The power of this approach is in the availability of relatively well determined path lengths to the HII regions, allowing us to derive the true three-dimensional space distribution of the cosmic-ray generated radiation field. From highly sensitive low-frequency radio observations and comparison with γ-ray observations, a number of important studies would follow immediately: the lifetimes of electrons in the interstellar and intergalactic gas and the competition between escape and energy loss; evidence of electron acceleration processes; the ratio of high energy electrons to protons in the ISM and intergalactic media (IGM) and in their sources; and, of course, the identification of absorption and emission regions with the positions of known objects, e.g., giant molecular clouds, SNRs, nearby normal galaxies and AGNs. As a general tool, the method could prove invaluable for study- 2.8. INTERSTELLAR PROCESSES 79

ing both Galactic and extragalactic energetic source populations, thereby enabling the extraction of the matter and magnetic ﬁeld distributions. Here the high and low energy radiation measurements provide a delineation of the electron and nuclear components, revealing unique information concerning the acceleration and transport of the diﬀerent species in many types of environments.

2.8.5 Interstellar Plasma Turbulence All Galactic and extragalactic radio sources are observed after their radiation has propagated through the Galactic plasma. Variations in the plasma density produce refractive index fluctuations, scaling as ν−2, which in turn scatter the radiation. The magnitude of radio-wave scattering from the interstellar plasma is strongly direction dependent, but the effects can remain significant even at frequencies as high as 10 GHz. The density (refractive index) microstructure responsible for interstellar scattering occurs on scales of order 1 AU. The density fluctuations, in turn, are thought to arise from velocity and/or magnetic field fluctuations. In addition to their corrupting effects, interstellar propagation effects are a powerful sub-parsec probe of the interstellar plasma, can provide a tracer of energy input into the ISM, and may be linked to cosmic ray propagation. Low frequency observations of compact sources provide a powerful diagnostic of propagation effects from the interstellar medium. The scatter-broadened angular diameter of a compact nonthermal source scales as λ2, while the resolution of a telescope and the minimum apparent size constrained by synchrotron self absorption both scale as λ. Thus, interstellar scattering observations are optimized with high- resolution, low-frequency observations. Current studies of interstellar scattering focus on regions of intense scattering (e.g., Cygnus and the Galactic center), where the scattering effects can be detected at frequencies near 1 GHz. SKA observations would be able to probe the density and field microstructure in less intense scattering regions, such as those near the Sun. Further areas of study include the search for supernova-generated turbulence, the z-distribution of scattering material, and the search for volume-limited scattering. Current shock acceleration theories (Ellison et al. 1984), relevant to the origin of cosmic rays, also suggest that upstream of a SNR should be an ideal site for the generation of the density fluctuations responsible for interstellar scattering. High frequency searches for the signatures of such upstream turbulence have a mixed record. The λ2 dependence of interstellar scattering would allow much more stringent tests to be applied. Ionized gas is found several kiloparsecs above the Galactic plane (the “Reynolds layer”). The scale height of the density fluctuations responsible for interstellar scattering is about 1 kpc, inferred from observations of high-latitude pulsars (particularly those in globular clusters) and low-frequency interplanetary scintillation measurements. The agents presumed to be responsible for generating density fluctuations— SNRs and HII regions—have a much smaller scale height ( 0.1 kpc). The SKA ∼ 80 CHAPTER 2. FORMATION AND EVOLUTION OF GALAXIES

could probe to higher latitudes than existing instruments and would provide additional information about the vertical distribution of the scattering material and clues about any other agents responsible for generating the density ﬂuctuations. Particu- larly valuable would be deviations from the smooth distribution of scattering material predicted by current models (Taylor & Cordes 1993).

The density fluctuations responsible for interstellar scattering have a spatial spectrum. The largest scale on which these density fluctuations occur is about 1 pc near the Sun and may be of the order of 0.01 pc in regions of intense scattering; this scale is presumably related to the injection of energy into the ISM that produces the density fluctuations. At low frequencies the angular extent of the scattering region may be less than the nominal scatter-broadened angular diameter of a background source. If so, the shape of the scatter-broadened image may be affected by the fact that the scattering is occurring only in a limited volume. The orientation and distortion of the image would then provide information about the volume in which the scattering was occurring.

2.8.6 Recombination Lines

Radio recombination lines of H, He, C, and heavier elements oﬀer the possibility of tracing temperature, kinematics, and ionization structure as well as abundances of heavy elements in ionized gas. Available from meter to mm wavelengths, strong radio lines are found in HII regions, but narrow, weak lines are also found in the very diﬀuse, ionized gas that pervades the Warm Ionized Medium. The sensitivity/resolution regime of the SKA will allow the kinematic imaging of ionized gas in a new range of astrophysically interesting circumstances, such as, for example, kinematic studies of the impact of SNR shocks on the surrounding ionized gas or metallicity maps of nearby galaxies. The stronger lines associated with bright HII regions could be imaged within starburst galaxies at higher redshifts.

As interstellar carbon recombines into very high Rydberg states (up to n = 768), absorption lines below 150 MHz are generated. The carbon atoms in these high states are very sensitive to the interstellar environment and permit excellent measurements of density, temperature, and ionization levels to be carried out (Payne et al. 1994). A number of Galactic regions that produce these lines have been found, including a large region that stretches 40◦ along the Galactic plane in the central region of the Galaxy (Erikson et al. 1995). The SKA would provide the sensitivity to identify many more regions for such diagnostic studies of the ISM and would allow these studies to be extended to external galaxies. 2.9. MAGNETIC FIELDS 81

2.9 Magnetic Fields

2.9.1 Rotation Measure Synthesis The radio emission of galactic and extragalactic radio sources often shows a significant amount of linear polarization. This linearly polarized signal can be utilized as an extremely effective probe of the intervening magneto-ionic medium between the source and observer since the intrinsic plane of polarization undergoes Faraday rotation as this emission propagates. The amount of that rotation is designated the Rotation Measure (RM) which is defined by:

2 θ = θ0 + RMλ (2.16)

2 where θ is the observed and θ0 the intrinsic polarization angle, RM is in rad/m and λ is the wavelength in meters. The diﬀerential rotation dθ across a bandwidth, dν, centered at frequency ν is given by:

dθ =2 RMλ2dν/ν (2.17) Suﬃciently narrow frequency channels must be used to measure θ, particularly at low frequencies, otherwise bandwidth depolarization will result. The RM depends on the properties of the medium as:

RM =0.81 Bkne(l)dl (2.18) Z where Bk is the longitudinal component of the magnetic field measured in µGauss, −3 ne is the electron density in cm and l is measured in pc. The λ2 dependence of the polarization plane noted above has often been used to produce estimates of the magnetic field structure and electron density distribution in the vicinity of radio sources and within the intervening medium. Traditionally only a small number of observations of the polarization angle at discrete wavelengths have been obtained and a least squares fit has been used to calculate the RM. The utility of the RM probe can be vastly enhanced by using very large instantaneous bandwidths together with high spectral resolution. In this case a coherent addition of the polarized signals, P~ =Q+jU, in the various bands can be formed as a function of the RM from:

2 I~(RM)= W (λ)P~ (λ)e−2iRMdλ dλ (2.19) Z where we have deﬁned the frequency window function, W(λ), which is determined by the locations, widths and shapes of the various frequency bands. This is eﬀectively a Fourier Transform with respect to the RM and produces a cube with two spatial axes and a third of RM. The point spread function in RM is given by:

2 T~(RM)= W (λ)e−2iRMdλ dλ (2.20) Z 82 CHAPTER 2. FORMATION AND EVOLUTION OF GALAXIES

PSR J0214+4232

8 bands, 319-380 MHz

0 -150 -100 -50 0

Figure 2.20: Plot of the polarized intensity as a function of RM for the pulsar J0214+4232.

With a suitable choice of the function, W(λ), a clean RM beam is achieved, providing both high RM resolution and a low sidelobe level. A practical example of this technique is given in Fig. 2.20 in which the RM in the direction of the pulsar J0214+4232 has been derived from only 8 frequency channels of 5 MHz width spread over the interval 319–380 MHz. Due to the very limited frequency sampling and uniform weighting in this example the RM sidelobe level is relatively high. Even so, the RM resolution determined by the inverse of the total observing bandwidth corresponds to about 12 rad/m2. Given a substantial signal-to-noise ratio, the RM centroid can of course be determined to much higher precision than the RM beamwidth. The range of RM’s which is accessible to a given observation depends on the observing wavelength, the channel bandwidth and the total observing bandwidth in addition to the instrumental sensitivity. Inserting the speciﬁcations for the SKA, we see that with a spectral coverage of 50% (ν/∆ν = 2) and a spectral resolution, ν/dν = 104, a one degree precision in the polarization angle would allow measurement of RM < 875 with an accuracy of about 0.02 rad/m2 at a wavelength of 1 m and RM| <| 2.4 105 with 4.8 rad/m2 precision at a wavelength of 6 cm. | | × An example of the application of this technique to a supernova remnant (SNR) is given by Gaensler et al (1998). With the sensitivity of the SKA accurate RMs will be obtainable for most SNRs. In conjunction with other information, this can be used as a (weak) indicator of distance. This new tool should be particularly powerful in the study of weakly polarized extended sources, such as giant radio galaxies where the RMs are known to be small. 2.9. MAGNETIC FIELDS 83

Small changes in the RM across the surface of extended sources should be easily traceable, especially if these RMs show some spatial coherence. Very small variations in RM can be expected to occur close to the core of AGN when polarized structure moves relative to a foreground Faraday screen. With a sensitivity of 0.02 rad/m2, extremely sensitive measurements of the ionized gas in front of polarized radio sources can be made. The observed RM is an integral along the line of sight. If the medium emitting the polarized waves is mixed with the medium rotating the wave vectors, the observed RM does not vary with λ2 anymore (Sokoloﬀ et al. 1998), and the observed RM cannot be used directly to compute magnetic ﬁeld strengths. On the other hand, any variation of RM with wavelength contains valuable information about the rotating medium. A large number of frequency channels, as planned for the SKA, is required to check the wavelength dependence of RM.

2.9.2 Polarization Studies of the Interstellar Medium in the Galaxy and in Nearby External Galaxies The bulk of the radio emission from the Galaxy at frequencies below 5 GHz is generated by the synchrotron process through the interaction of relativistic electrons with magnetic fields. Studies of the polarization state of that radiation gives us the opportunity to measure many properties of those fields, and the enhanced resolution and sensitivity that will become available with the SKA will open new opportunities. As synchrotron emission propagates through the intervening magneto-ionic medium its inherent linear polarization suffers Faraday rotation effects depending on both the field strength and the electron density, and measurements of Faraday rotation effects can give information on parameters which is unobtainable by other means. Lastly, Zeeman splitting of spectral lines from maser sources offers the possibility of direct field strength measurements under a variety of astrophysical conditions. Our knowledge of the magnetic field configuration of the Galaxy is based on measurements of the Faraday rotation of signals from extragalactic sources as they propagate through the Galaxy (e.g. Simard-Normandin et al. 1981, Broten et al. 1988, Han et al. 1997). The deficiency of this method of probing the field has been the relatively small number of sources suitable for these measurements (at most a few hundred). The SKA will increase the number of sources available for this kind of work by a large factor, leading to a dramatic improvement in our detailed knowledge of the field configuration. It will also be possible to extend this field mapping technique to external galaxies. However, with the present telescopes only a handful of polarized background sources can be observed even in the most extended galaxy M31 (Han et al. 1998). With the SKA, the much larger number of suitable sources which could be seen through a face-on galaxy, such as M101, would allow a mapping of the field and, with the aid of other data, an estimate of the field strength that would complement information available from mapping of the synchrotron emission from the M101 itself. By this method even magnetic fields without detectable synchrotron emission together 84 CHAPTER 2. FORMATION AND EVOLUTION OF GALAXIES

with thin ionized gas, e.g. in galactic halos, can be discovered. In a similar way Faraday rotation probing of SNRs (e.g. Kim et al. 1988) and other objects, such as clusters of galaxies (e.g. Kim et al. 1990), would be improved by a very large factor. RM determinations using pulsar signals add to our knowledge of the Galactic field configuration, with the added advantage that we have some knowledge of pulsar distance, albeit with considerable errors. The SKA will contribute in this field too, largely through the discovery of many additional pulsars. Numerous nearby galaxies have now been studied in some detail with angular resolutions from arcminutes to tens of arcseconds, yielding a resolution of some 100’s of parsecs at the distance of galaxies in the Local Group. Accurate flux determinations at various frequencies can be used to deduce the magnetic field strength through the equipartition argument. The mapping of linear polarization (at several frequencies to correct for Faraday rotation) has shown surprisingly well-aligned magnetic fields (Beck et al. 1996). In fact, magnetic fields commensurate with the spiral arm structure have been traced. However there are indications of deviations from the large-scale structures both in the spiral arms and in the centres of galaxies. In several galaxies regular fields are concentrated between the optical spiral arms. The sensitivity and high angular resolution of the SKA will play a key role in understanding these structures in view of the competing theories of the origin of the magnetic field. There has been a resurgence of interest in the polarization of the Galactic synchrotron emission, usually referred to as the “background” emission, but actually in the foreground. After a lull since the 1970s due to a lack of improvement in instruments, the subject has been revived through the advent of sensitive systems capable of precise polarimetry over wide areas of the sky. Large regions of the sky have now been surveyed with resolution of the order of 10 arcminutes, in the Southern sky by Duncan et al. (1997) and in the Northern sky by Uyaniker et al. (1998 and 1999). Synthesis telescopes have also become capable of wide-field polarimetry, with good understanding of instrumental effects and calibration. This has enabled polarization mapping with arcminute resolution at 327 MHz with the Westerbork Telescope (Wieringa et al. 1993) and at 1420 MHz with the DRAO Synthesis Telescope (Gray et al. 1998 and 1999). Many effects are seen in the quoted papers which are predominantly effects in angle, since they have no significant counterpart in polarized intensity. Fig. 2.21 shows an example of such a region from the Canadian Galactic Plane Survey being made with the DRAO Synthesis Telescope. Such features are interpreted as the effects of a “Faraday screen”, a region of ionized gas threaded by magnetic field lying between the source of the emission and the observer. Of course, Faraday rotation also occurs within the emitting volume, complicating interpretation. Faraday rotation depends on the product of the electron density and the field strength. In a typical Galactic field of 3 µG, it is easy to detect the Faraday rotation produced by an ionized region in the foreground with an emission measure 1 cm−6pc 2.9. MAGNETIC FIELDS 85

Figure 2.21: An image at λ21 cm of effects of the diffuse Faraday Screen in the Milky Way. This image from the Canadian Galactic Plane Survey shows the spatial variations in polarisation position angle imposed by changes in the structure of the Galactic magnetic field and diffuse electrons distribution on scales of 10’s of pc. The SKA allow this technique to be used on other galaxies to investigate the structure and generation mechanisms of galactic magnetic fields. 86 CHAPTER 2. FORMATION AND EVOLUTION OF GALAXIES

at 1.5 GHz. This is much higher sensitivity than most other methods of detecting ionized gas, which may be either an asset or an insuperable complication. In certain circumstances sufficient information is available to allow the product of field and electron density to be separated. For example, Gray et al. (1999) have been able to measure the magnetic field strength in the extended ionised envelope of the HII region W4. The enhanced angular resolution of the SKA will permit such measurements to much greater distances, and with a vastly increased number of HII regions. The high density and the turbulent structure of the ionized gas in HII regions cause very high and very rapidly variable values of rotation measure through these objects, which effectively eliminates the polarization of synchrotron radiation propagating through them (by bandwidth or by beam depolarization). This effect can be exploited to limit the distance to polarization structures (Gray et al. 1998). this technique is analogous to the use of HII regions to absorb the background synchrotron emission in order to measure the synchrotron emissivity in the foreground. With two orders of magnitude improvement in sensitivity and with the small beam of the SKA such studies at high frequencies (e.g. 5 GHz) can be extended to much greater distances (right across the Galaxy). Depolarization will be much less so that the polarized emission from distant regions can be studied. Dickey (1997) has shown that HI absorption can be used to place limits on the distance to the polarized Galactic emission. Spectral-line mapping in Stokes Q and U removes all emission features because the HI emission is not polarized. Any absorption detected must arise from absorption of polarized continuum emission. With the sensitivity of the SKA, this technique can be extended to obtain absorption distances to other Galactic synchrotron emitters too faint for other methods. Examples include the non-thermal emission from WR-stars (Williams et al. 1997), from novae (Reynolds & Chevalier 1984), and from low-surface-brightness SNRs. The SKA will play an important part in the study of SNRs. In addition to the detailed morphological information which will come from total-power images of excellent sensitivity and resolution, there is much to be learned from polarimetry of these objects. It is generally believed that cosmic-ray acceleration occurs at the shock fronts of SNRs, but many aspects of acceleration theory can only be improved through much more detailed knowledge of magnetic field structure at the shock front. Detailed examination of field structures in young SNRs will allow study of turbulent field amplification at the shock front. Detailed examination of the interaction of the shock fronts of older remnants with the clumpy interstellar medium (ISM) will show the role of the magnetic field in moderating these interactions, which are a significant source of energy input to the ISM. The magnetic field must play an important role in the coupling of SN ejecta with the ISM, currently a poorly understood process. Such coupling can be seen in some SNRs (e.g. Cas A (Anderson & Rudnick 1995) and the Vela SNR (Aschenbach et al. 1995)) and the sensitivity and resolution of the SKA will permit such studies to be extended to other objects. 2.9. MAGNETIC FIELDS 87

Finally, the very high sensitivity of the SKA will make it the premier instrument for the measurement of magnetic fields through the Zeeman effect. Measurements of the Zeeman effect require the detection of a very small difference between the signals received in the two senses of circular polarization and it is extremely difficult to account for all instrumental effects. Measurements using the H I line at 1420 MHz have been controversial, and it appears that measurements to date which have detected field strengths of the order of 10 µG are inconclusive. However, the method has proved useful in circumstances where the field is enhanced, for example in the shells of SNRs. Frail et al. (1994) and Claussen et al. (1997) have used Zeeman splitting of the 1720 MHz line of OH, excited into maser emission by the interaction of the SNR shock with molecular material, to measure fields of order 0.2 mG. The SKA has the potential to make Zeeman measurements at much lower field strengths. An angular dimension of 0.1′′; corresponds to only 0.3pc at the distance to M31, and magnetic fields are expected to be highly uniform on this scale in many regions of M31. The average field strength in M31 is only 5 MicroGauss, but on small scales we may expect 10 MicroGauss, enough to see the Zeeman effect. Galaxies with higher field strengths are e.g. NGC6946 and M51 with 10 MicroGauss on average, locally up to 20 MicroGauss. Here the spatial resolution of the SKA is a few pc, but uniform fields can be expected even on this allowing Zeeman measurements of the strength. The successful pursuit of these goals will require careful attention to the polarization performance of the instrument.

Chapter 3

Formation and Evolution of Stars

3.1 Continuum Radio Emission from Stars

Since the construction of large arrays like the WSRT and VLA in the 1970’s, the ﬁeld of stellar radio astronomy has advanced tremendously. Radio emission has been detected from all stages of stellar evolution, from birth to death, and in all these stages has shown us astrophysical phenomena and stellar activity not detectable by any other means. However, further major advances in stellar radio astronomy are limited by sensitivity. Only a few hundred stellar radio sources are now known. The SKA will increase this number by over four orders of magnitude to at least 106 stars (Seaquist 1996). This phenomenal increase will bring about a new era in the ﬁeld of stellar radio astronomy. New classes of stars and previously unknown phenomena related to stellar radio emission are certain to be discovered. The discipline will literally be reborn. The loci of 441 radio-detected stellar systems on a Hertzsprung-Russell Diagram (HRD, brightness MV vs. color B V ) is shown in Fig. 3.1. We clearly see the main sequence, a clump in the G-K subgiant/giant− area, mostly due to RS CVn binary systems and PMS, and the concentration in the red giant/supergiant area, mostly associated with the cool components of symbiotic stars (to which the radio emission has been attributed). Despite the apparent plenitude of radio detections across the HR diagram, many features result from strong selection bias and detection limitations. Obviously, the observed radio luminosity increases drastically with increasing absolute brightness (decreasing MV). But since the space density at the same time steeply decreases, optically luminous stars are on average distant objects, with only very radio-luminous examples being detected by radio telescopes. The radio luminosity distribution is graphically illustrated in two panels of Fig. 3.2. The 2-D surfaces represent an approximation to the envelope of the peak radio luminosity detectable as a function of MV and B V (binning, peak search, and median smoothing was applied for this representation).− The right panel represents only single stars (or components in binaries that are not tidally interacting and which are the likely sources of the observed emission; PMS are excluded).

89 90 CHAPTER 3. FORMATION AND EVOLUTION OF STARS

-10

-5

0 V

M 5

size: log(LR) color: type

12 single stars 13 RS CVn & Algol binaries 10 14 BY Dra & W UMa binaries 15 other binaries 16 PMS, T Tau, Herbig Ae/Be stars 17 symbiotic & shell stars 15 18 Bp/Ap stars 19 WR stars

OBAFGKM 20 -0.5 0 0.5 1.0 1.5 2.0 B-V

Figure 3.1: The Radio HR Diagram showing the loci of 441 radio detected stellar systems. The stellar classes are color-coded, and the symbol sizes indicate the radio luminosity. Only detections within 1–10 GHz have been considered. The radio star catalog by Wendker (1995; 1998) has been used, complemented by (mostly Hipparcos) distances and B V colors, with some color corrections due to emission lines (WR − stars) or to extinction. 3.1. CONTINUUM RADIO EMISSION FROM STARS 91

Figure 3.2: The distribution of radio luminosity with position in the Hertzsprung- Russell Diagram. The envelope of peak radio luminosities is shown as a 2-D surface. The left panel includes all stars (see Fig. 3.1). The left includes only single stars (or components in binary systems that are not tidally interacting).

With very rare exception, normal stellar photospheric emission cannot be detected with current sensitivities. Detectable radio emission is virtually always associated with active stellar phenomena, such as energetic outbursts or mass loss. Radio emission is generated in diﬀerent atmospheric environments characterizing diﬀerent areas in the HRD. In hot stars, radio emission is believed to stem from electrons accelerated in shocks that form in strong stellar winds. In the cool-giant area, many sources are binaries, some of which produce radio emission in shocks of colliding winds; also, chromospheric emission is detected from red giant stars. These range of phenomena can be broadly categorized based on the characteristic power law of the radio continuum spectrum, ranging from:

normal stellar photospheres (ν+2) • bremsstrahlung emission from circumstellar environments and stellar ejecta in • the form of winds (WR and OB stars), shells (novae, planetary nebulae), and jets (symbiotic stars, class 0 PMS stars) (ν+1)

synchrotron and gyrosynchrotron emission from ﬂare stars, active binary sys- • tems, pre-main sequence (PMS) stars, and x-ray binaries (ν−1)

pulsars (ν−2) • For all these stellar types we have only a very poor census of properties. Our current understanding of these phenomena is based on investigations of a few bright 92 CHAPTER 3. FORMATION AND EVOLUTION OF STARS

Figure 3.3: The ﬂux density expected for a typical Mira variable as a function of frequency and distance compared to the continuum sensitivity of the SKA after 8 hours integration (dashed line). Also shown is the sensitivities of the VLA and LSA/MMA for the same integration time.

members of the class. For pulsars, non-thermal stars, and thermal circumstellar emission (type 1, 2 and 3 listed above), the SKA will uniquely afford the sensitivity to provide measurements of radio properties for whole stellar populations, allowing meaningful studies of the relationship between radio phenomena and other stellar properties and evolutionary states. Many classes will be detected over the entire volume of the Galaxy and in the Magellanic Clouds. Pulsars will be detected throughout the Local Group of galaxies. In the first class of phenomena, stellar photospheres, emission rises rapidly toward shorter wavelengths, and planned millimeter arrays will detect a large number. However, at centimeter wavelengths, the SKA will study a similar number of these stars and the high angular resolution at the shortest wavelengths will allow direct imaging of the stellar surfaces. Figure 3.3 shows the flux densities expected from the photosphere of a Mira variable star at various distances. Also shown is the continuum sensitivity of the SKA as a function of frequency. Presently, isolated Miras can be detected at radio wavelengths out to a distance of only a few hundred pc. With the SKA this will be increased to more than 10 kpc, making possible the detection of emission from Mira stars throughout most of the Galaxy. The expected flux densities for several classes of non-thermal stars versus distance are shown in Fig. 3.4, along with the continuum sensitivity of the SKA at 5 GHz after 8 hours integration. With the SKA the quiescent radio emission from the Sun would be detectable to a distance of 100 pc. This volume contains 104 G-dwarfs. It ∼ ∼ would be possible for the first time to explore the Solar-stellar connection in the radio, and place the radio phenomena of the Sun within the context of the properties of stars such as rotation, mass, magnetic field and chemical abundance. Even the nearest stars with solar radio luminosity are undetectable at present. 3.2. IMAGING THE SURFACES OF STARS 93

Figure 3.4: Flux density at 5 GHz versus distance for a number of non-thermally emitting stars compared with the continuum sensitivity of the upgraded VLA and the SKA for an integration time of 8 hours. The luminosities used are typical values. (after Seaquist 1996).

Quiescent radio emission from ﬂare stars is detectable out to several kpc, and typical RS CVn systems, PMS stars, peculiar magnetic stars, and non-thermal components of WR winds would be detectable over the entire Galaxy and the Magellanic Clouds. X-ray binary systems and brighter WR stars would be detectable in M31. For most of these classes, only a handful of objects are now known. With the SKA, in addition to studying large and complete samples of these populations, we will be able to trace their spatial distributions in star formation regions and clusters and within the Galaxy as a whole, thereby relating radio properties to characteristics such as age and environment.

3.2 Imaging the Surfaces of Stars

3.2.1 Red Giants and Supergiant Stars Imaging the surfaces/atmospheres of other stars is one of the major frontiers in stellar astronomy. Because red giant and supergiant stars present the largest angular diameters in the sky, they are the first stars apart from the Sun to have been imaged, albeit so far with angular resolution just sufficient to resolve their disks. These observations confirm that red giant and supergiant stars have highly extended atmospheres, which from millimeter and infrared observations are known to give rise to a massive outflow. The SKA will be able to image the structure and directly measure the temperature of these stellar atmospheres over a relatively large range of heights, leading to more ac- 94 CHAPTER 3. FORMATION AND EVOLUTION OF STARS

Figure 3.5: Left panel shows a false color near-ultraviolet image of Mira made with the HST with an angular resolution of 45 mas (from Karovska et al. 1997). The optical photosphere of Mira is clearly asymmetric,∼ perhaps as a consequence of nonradial pulsations. Right panel shows a contour plot of Betelgeuse’s optical photosphere made with aperture-masking interferometry (from Tuthill et al. 1997). The two apparent bright spots may be produced by large convection cells distributed inhomogeneously over the stellar surface. curate empirical models of their physical properties as well as a better understanding of the mechanisms that drive the stellar atmosphere outwards. Optical images of both red giant and supergiant stars reveal the common occurrence of asymmetric photospheric structures on these stars. Examples are shown in Fig. 3.5 for the red giant Mira and the red supergiant Betelgeuse. Their asymmetric photospheric structures are attributed to bright spots produced by large convection cells distributed inhomogenously over the stellar surface, or intrinsic stellar distortions possibly produced by nonradial pulsations. Red giant and supergiant stars are ubiquitous ultraviolet emitters. An ultraviolet image of the red supergiant Betelgeuse made by the Hubble Space Telescope (HST) reveals a hot chromosphere that extends to many stellar radii and is apparently asymmetric (Gilliland & Dupree 1996a,b), as shown in Fig. 3.6. Acoustic waves, Alf´ven waves, and radial pulsations have all been proposed for heating thereby greatly extending the stellar atmosphere as well as driving its mass outﬂow, but the available observations do not permit us to choose between the various possibilities. Both red giant and supergiant stars lose mass at a prodigious rate of 10−7– −6 −1 ∼ 10 M⊙ yr . On red giant stars, radiation pressure on dust grains condensed from dense gas in the lower stellar atmosphere is thought to drive the mass outﬂow (Kwok 3.2. IMAGING THE SURFACES OF STARS 95

Figure 3.6: Left panel shows a false color ultraviolet continuum image of Betelgeuse’s chromosphere made with the HST with an angular resolution of 38 mas (from ∼ Gilliland & Dupree et al. 1996b). Right panel shows a ﬁtted model image comprising an oﬀset hot spot superposed on a circular disk. The black circle has an angular diameter of 55 mas, somewhat larger than the typically inferred optical photospheric diameter of 45 mas. The chromosphere seen here in the ultraviolet continuum ∼ extends to a height of nearly 3R∗, but that seen in the line of Mg II extends to at least twice this height (Gilliland & Dupree 1996a). 96 CHAPTER 3. FORMATION AND EVOLUTION OF STARS

Figure 3.7: The ﬂux density (left panel) and angular diameter (right panel) expected for a Mira-like (late-M) red giant star as a function of radio frequency and distance, assuming S = 5.1 ν2 and an absolute diameter of 8 AU as found by Reid & µJy × GHz Menten (1997). The 3σ continuum sensitivity of the SKA is plotted for an integration time of 8 hrs in one polarization (left panel), as is the angular resolution of the telescope (right panel).

1975). The situation is less clear for red supergiant stars as dust is not expected to form in a hot chromosphere, but the mechanisms involved presumably also operate in the lower regions of the extended stellar atmosphere. This mass loss has important consequences for it significantly affects the evolution of the star, forms a circumstellar shell that later helps shape the structure of protoplanetary and planetary nebulae (on post red giant stars) and supernova remnants (on post red supergiant stars), and is one of the most important sources of enrichment for the interstellar medium. Thus, study of the properties and dynamics of the stellar atmosphere leading to a better understanding of the mechanisms responsible for their large extents and mass outflows is one of the important topics in red giant and supergiant star research. The large radio photospheres of red giant and supergiant stars give rise to detectable thermal radio emission. From a sample of 6 stars, Reid & Menten (1997) showed that the radio flux density spectrum of late-M giant stars such as Mira follows a S ν2 dependence at centimeter wavelengths (8–22 GHz), suggesting that their radio∝ photosphere has a sharp opacity edge. Fig. 3.7 (left panel) shows the expected flux density from Mira-like giant stars as a function of radio frequency at various distances, and for comparison the sensitivity of the SKA for an 8-hr integration over a range of radio frequencies. As is apparent, the SKA will be able to detect the radio photospheres of these stars out to many kiloparsecs at high frequencies, and right across the entire galaxy at 20 GHz. Although the radio photospheres of red giant stars have yet to be imaged, Reid & Menten (1997) have partially resolved the radio photosphere of the red giant W Hydra. They find an average diameter of 0′′.080 0′′.015 (corresponding to 8 AU), and a ± ∼ brightness temperature of 1500 570 K. This corresponds to a disk size approximately ± 3.2. IMAGING THE SURFACES OF STARS 97

Figure 3.8: The flux density (left panel) and angular diameter (right panel) expected for a Betelgeuse-like (early-M) red supergiant star as a function of radio frequency 1.3 and distance, assuming SµJy = 240 νGHz as found by Newell & Hjellming (1982) and the radio photospheric diameters measured× by Lim et al. (1998). The 3σ continuum sensitivity of the SKA is plotted for an integration time of 8 hrs in one polarization (left panel), as is the angular resolution of the telescope (right panel). twice that measured in the optical, and a temperature two-thirds that of the optical photosphere. This relative dimension and temperature, when applied to the other red giant stars in the sample of Reid & Menten (1997), also suitably explain their observed flux density. With a brightness sensitivity of < 10 K at frequencies above 300 MHz (for an integration time of 8 hrs in one polarization)∼ , far below the brightness temperature of the stellar radio disk, the SKA will be able to image with high fidelity the radio photospheres of all red giant stars with angular sizes large enough to be resolved. Assuming a radio photosphere of diameter 8 AU, Fig. 3.7 (right panel) shows the angular diameter a Mira-like giant star will subtend at various distances compared to the angular resolution of the SKA as a function of frequency. At the highest frequency of 20 GHz, it will be possible to resolve the radio photosphere of Mira-like red giant stars to distances of up to 1.5 kpc. For these stars, comparison ∼ of the radio photospheric structure to the optical photospheric structure will provide important constraints or help elucidate the mechanism(s) responsible for driving the stellar atmosphere outwards. Early-M supergiant stars have a flatter radio flux density spectrum of S ν1.0−1.3 ∝ as observed for Betelgeuse (Newell & Hjellming 1982) and Antares (Hjellming & Newell 1983). Fig. 3.8 (left panel) shows the expected flux density from Betelgeuse- like supergiant stars as a function of radio frequency at various distances, and for comparison the sensitivity of the SKA for an 8-hr integration over a range of radio frequencies. As is apparent, the SKA will be able to detect the radio photospheres of all red supergiant stars in our Galaxy, and at 20 GHz even those in the Large and Small Magellanic Clouds. Using the VLA, Lim et al. (1998) have imaged the radio photosphere of Betel- 98 CHAPTER 3. FORMATION AND EVOLUTION OF STARS

geuse at 7 mm. As shown in Fig. 3.9 (left panel), the star appears to be asymmetric. Measurements taken simultanously at cm-wavelengths partially resolve the stellar atmosphere, and directly measure its radial temperature profile. As shown in Fig. 3.9 (right panel), these measurements reveal that the stellar atmosphere has a temperature close to the photospheric value at a height of R∗, and from there decreases in temperature with height. The height range probed∼ in radio is identical to that probed in the ultraviolet by the HST, but the latter reveals a hot chromosphere at temperatures > 5000 K. Although these two components must therefore coexist in the lower stellar∼ atmosphere, the much lower radio opacity per particle of the cooler radio atmosphere implies that it must be much more abundant than the hotter chromospheric component and forms the dominant component in the stellar atmosphere. Lim et al. (1998) suggested that the elevation of photospheric gas by large convection cells distributed inhomogeneously over the stellar surface naturally explains the observed temperatures and asymmetric structure of the stellar atmosphere. Shock waves produced by the elevated gas, particularly strong over the convection cells, may be responsible for heating a small fraction of the atmosphere to chromospheric temperatures. These measurements resolve the previous puzzling observations of episodic dust formation at heights of 3R∗ (Bester et al. 1996), difficult to explain in the presence of a pervasive hot chromosphere.∼ Instead, the measured dominant gas temperature of 2000 K at this height is just sufficiently low for dust grains to condense. Radiation pressure∼ on these dust grains could further elevate the stellar atmosphere, and ultimately drive its mass outflow.

The SKA will have an especially significant impact on studies of red supergiant star atmospheres. Fig. 3.8 (right panel) shows the angular diameter subtended by Betelguese’s radio photosphere at different frequencies as a function of distance, and for comparison the angular resolution of the SKA. For those nearer than 2 kpc, the SKA will be able to directly measure the temperature and structure of the stellar atmosphere over a large range of heights. Comparisons of the atmospheric structure at different heights will be vital for understanding both the physical properties and dynamics of the stellar atmosphere, and (combined with optical photospheric images) will provide powerful constraints on the mechanisms responsible for driving the atmosphere outwards.

Finally, an exciting prospect not yet studied numerically is the possibility of detecting and imaging circularly polarized radio emission from red giant and supergiant stars from regions of relatively strong and ordered magnetic fields. The detection of polarization in circumstellar SiO maser emission demonstrates that at least some red giant stars possess significant magnetic fields (10–100 G at the observed heights), which furthermore appears to be relatively well ordered on a global scale (Kemball & Diamond 1997). It is not known whether red supergiant stars possess significant magnetic fields. Such magnetic fields, if strong and pervasive on red giant and supergiant stars, may play an important role in shaping the structure of the stellar surface and influencing both the rate and direction of the mass loss. 3.2. IMAGING THE SURFACES OF STARS 99

Figure 3.9: A false color image of Betelgeuse made at a wavelength of 7 mm with the VLA (left panel). The radio photosphere at λ7 mm has a size approximately twice the optical photosphere (45 mas), and is asymmetric. Simultaneous observations at longer wavelengths partially resolve the radio photosphere, which increases in size with increasing wavelength, and are used to deduce the radial temperature proﬁle of the atmosphere (right panel). From Lim et al. (1998). With surface brightness sensitivity of a few 10’s of K at angular resolution below 10 mas, the SKA will directly image the surfaces of red giant and supergiant stars beyond a kpc distance. 100 CHAPTER 3. FORMATION AND EVOLUTION OF STARS

Name Max No. Baselines Max Baseline length Element Diameter LBT 1 20 m 8 m ISI 1 35 m 1.7 m GI2T 1 60 m 1.5 m Magellan 1 60 m 6.5 m PTI 1 110 m 0.4 m I2T 1 140 m 0.3 m SUSI 1 640 m 0.2 m IOTA 1-3 45 m 0.5 m Keck 1-15 180 m 10/1.5 m COAST 3-6 100 m 0.4 m NPOI 3-15 250 m 0.4 m VLTI > 6 200 m 8/1.8 m

Table 3.1: Some Current and Planned Optical-IR Interferometers

3.2.2 Complementarity to Planned Optical-IR Interferome- ters

By the time the Square Kilometer Array is constructed, a number of powerful optical and IR interferometers (OIRI) will also be in operation. Some of the major eﬀorts are summarized in Table 3.1. Names of the interferometers are given in column 1. Column 2 shows the maximum number of baselines existing or planned for that instrument. The maximum baselines (in meters) is shown in column 3 while column 4 shows the diameters of the elements in the array. Some instruments, like COAST, NPOI and SUSI are already operational. Consideration of angular resolution suggests that the SKA and OIRI will be quite compatible. Fig. 3.10 shows a plot of SKA operating wavelength plotted against those of the planned VLTI. The angular resolution of the VLTI which scales inversely with wavelength is shown across the top of the graph. The diagonal line indicates for which radio wavelengths the SKA produces the same angular resolution. We see that the match is excellent for the IR wavelength range 5 - 50 microns and the radio range of 1 to 20 cm. Since the SKA is not planned to operate below 1 cm, the match with the optical resolution capabilities of the OIRI is not as good, although the latter can always be reduced to match the radio capabilities. Because of the common angular resolution and the common sensitivity to thermal sources the most obvious science overlap is in the study of stellar surfaces. At its highest operating frequency the SKA can approach an angular resolution of 5 milliarcseconds and can therefore resolve the 400 or so stars whose angular diameters are greater than about 10 milliarcseconds. The largest stars have angular diameters approaching 0.1 seconds of arc. Thus, in the best cases SKA will produce stellar images 3.3. STAR FORMATION 101

0.005 0.05 0.1

Figure 3.10: Angular Resolution of the SKA and the VLTI Optical-IR Inteferometer. The diagonal line indicates for which radio wavelengths the SKA produces the same angular resolution. There is an excellent match for IR wavelength range 5 - 50 microns and the radio range of 1 to 20 cm. with as many as a couple of hundred resolution elements. Such imaging can record surface features like hot spots and starspots as has been seen in the extensively studied supergiant Betelgeuse. The OIRI will match or exceed the (1 cm) SKA resolution from the optical up to about a wavelength of 5 microns. The availability of common- resolution imaging at optical, IR and radio wavelengths should provide the ability to image the atmospheres of these stars in 3-dimensions because the optical depth is a strong function of observing frequency. This type of imaging should allow the inves- tigation of the vertical structure of the photospheres and via long term monitoring of surface and stellar wind features make it possible to correlate surface activity with changes in the stellar wind (e.g. Lim, 1998). Such studies of the atmosphere-wind interface will not only lead to a better understanding of the wind phenomenon but may also shed light on the dynamo that drives surface activity (e.g. G¨uedel, 1997). Additionally, such data place important constraints on convection models leading to a better understanding of energy transport in these stars.

3.3 Star Formation

The formation of stars, particularly those of low mass like the Sun, is fast becoming one of the cornerstones of modern astrophysics. This is due in part to the advance of observational techniques, which have brought us to the point where studies of the formation of individual stars and their planetary systems are now feasible. One of 102 CHAPTER 3. FORMATION AND EVOLUTION OF STARS

Figure 3.11: VLA images of HL Tau at λ7 mm a showing structures on scales of 10’s to 100’s of AU. At low” resolution (0.2 arc sec) the extended emission from the disk dominates (left), while at higher resolution (0.05 arc sec) the brighter jet dominates (right). From Wilner, Ho and Rodriguez (1999). the links in the chain of cosmic events leading from the birth of the Universe to the emergence of intelligent life, the problem of low mass star formation and planet formation is poised to make enormous advances in the early 21st Century. This eﬀort also commands wide interest in society as a whole. Since stars form in dense molecular clouds, proposed millimetre and submillimetre wave arrays will take a leading role in star formation research in the next century. However, it is clear that the SKA working at high angular resolution at centimetre wavelengths will make critical and unique contributions to star formation. There are many molecular line transitions available at low frequencies, some of which are uniquely capable of probing certain aspects of the physics of molecular clouds and their collapse to form stars. At radio continuum wavelengths emission from protostars is comprised of a dusty moleclar disk and collimated ionized jets (see Fig 3.11). Widespread study of the physical regimes and processes underlying these phenomena at AU scales will require the SKA.

3.3.1 Protostellar Cores Dust around protostars Observing the formation and evolution of circumstellar disks is crucial for understanding the star formation and planet-building processes. These disks, 100 AU ∼ 3.3. STAR FORMATION 103

in radius, tens to a few AU thick, and with masses 0.1 solar masses, are usually studied in the dust thermal continuum at (sub)millimetre∼ wavelengths in the case of young deeply embedded objects, although for optically visible pre-Main sequence stars, such disks can be seen with advanced optical telescopes as the HST or the new generation of ground-based optical telescopes. Little is yet known about the physical conditions and processes within protostellar disks. Planned (sub)mm interferometers will observe dusty disks at an angular resolution of 10 milli-arcseconds, corresponding to a size scale 5 AU at a distance of ∼ ∼ 500 pc. This angular/linear resolution will permit the study of the mass and thermal structure of disks, disk vertical stratification, the spatial distribution of dust properties, the distribution, kinematics, and chemistry of molecular components, and the development of circumstellar structures in binary systems. Given these capabilities in the millimetre and submillimetre range, the SKA will nevertheless have a unique advantage when observing protostellar disks at centimetre wavelengths: dust emission is optically thin at centimetre wavelengths, whereas it may be optically thick in the (sub)millimetre in very dense gas condensations. The consequence of optically thick dust emission is that it complicates the determination of dust masses and dust optical properties through measurement of the continuum spectral index, and attenuates the spectral line emission from molecules within the condensation. It is thus possible that only at centimetre wavelengths with the SKA we will be able to probe the bulk of the material in the inner tens of AU where surface mass densities may be greater than 103 g cm−2. Recent observations of the embedded protostellar object L1641N (IRAS 05338- 0624) by Chen et al. (1995) provide evidence that dust continuum emission in low mass protostars can be optically thick even at millimetre wavelengths. These researchers found that the continuum spectrum of L1641N between 5 GHz and 200 GHz can be understood in terms of two components: optically-thin free-free emission from an ionized protostellar wind (spectrum ν−0.1) dominating at low frequencies, and dust ∼ thermal emission (spectrum ν+2.1) dominating at high frequencies. Assuming the standard dust emissivity relationship∼ (β 2), the spectral index of approximately ∼ +2.1 found between λ7mm and λ1.3cm wavelength suggests that the dust emission is optically thick in that wavelength range (in the optically-thick regime we expect the spectrum to go as ν+2, whereas for optically-thin emission we expect ν(+2+β)). The alternative interpretation, that the emission is optically thin, would imply β 0.1, meaning that the dust grains were really large fluffy “snowballs” and implying∼ an unreasonably massive disk. Note that the determination of the dust emissivity parameter β is of great importance for understanding dust properties and determining dust masses, and can only be determined from observations in which the emission is optically thin. In the case of L1641N, observations longward of λ1cm are required. ∼ The wavelength at which dust emission becomes optically thick depends critically on the dust properties. In the simplest model, the optical depth is given by

λ −β τ = κo σ ( ) · · λo 104 CHAPTER 3. FORMATION AND EVOLUTION OF STARS

Figure 3.12: The wavelength for optical depth unity for a spherical cloud of total mass 1 M⊙ (dust mass of 0.01 M⊙), as a function of linear diameter. Typical ISM dust 2 −1 properties have been assumed (κo = 0.1 cm gm , λo = 250 µm and β = 2). High resolution imaging in the cm regime is required to probe the cloud on scales below 10-20 AU. With linear resolution of 1 AU at distances of a few hundred pc, the SKA will uniquely image AU-scale structures∼ within such dense protostellar cores.

2 −1 where κo is the opacity in units of cm g at the reference wavelength λo, β is the frequency dependence of the dust emissivity, and σ is the dust mass column density in g cm−2. Typical values in the diffuse ISM might be κ 0.1 cm2 g−1 at λ 250 o ∼ o ∼ µm, and β 2. yhe wavelength of unit optical depth through a spherical dust cloud with yhese∼ dust properties is shown as a function of cloud diameter in Figure 3.12. it has been suggested that due to the evolution of dust properties in cold dense cores and in dense protoplanetary disks, κo might be > 10 times larger than the diffuse ISM values and β may be as small as 1. ∼ As an example, following the discussion of Men’shchikov & Henning (1997), for 2 −1 κo = 7 cm g at λo = 1 mm, and β 1, a protoplanetary disk with a total gas plus dust mass of 0.1 M (dust mass of 0.001∼ M ) and disk diameter of 100 AU has ⊙ ⊙ ∼ σ = 1 g cm2. With these assumptions, the dust opacity τ is > 1 shortward of λ7mm. Dust emission from very young Class 0 protostellar objects (< 105 years old?) such as NGC1333/IRAS-4 accounts for 20% to 100% of the emission at λ1.3cm (Mundy et al. 1993). Other somewhat older objects, Class I and T Tauri stars also have very steep centimetre wave spectral indices ν(2.3−3.2) which cannot be explained by ionized gas emission, and thus must be dominated by the long wavelength tail of thermal dust emission. Most such observations of centimetre wave dust emission have been made with the current VLA. The SKA will have no difficulty in detecting dust emission from even the most evolved pre-Main sequence Class III objects, in which disk evolution has gone on the longest and in which the planet formation process may be in full 3.3. STAR FORMATION 105

swing.

Dynamics and Chemistry

SKA observations of the NH3 molecule will allow dynamic imaging of the dense molecular gas associated with star formation cores in our Galaxy on sub-AU scales. Many regions of active star formation exist within a few hundred parsecs of the Sun. At these distances milli-arcsecond angular scales correspond to dimensions 0.1 AU. As demonstrated above, it could well be that the optical depths due to dust∼ in very dense regions may be high at (sub)millimetre wavelengths, meaning that molecular line emission from very high column density regions may be severely attenuated. If so, observations of molecular species such as NH3, H2CO, CH3OH, and carbon chain molecules at centimetre wavelengths (where the dust is optically thin) will be required in order to probe the gas chemistry and dynamics in the densest molecular condensations and on (sub)AU scales. What the characteristics of these regions might be are suggested by recent new models of dust emission towards L1551/IRS5, which suggests the existence of a region at the core of size 100 AU having a gas density ∼ 109 cm−3. Men’shchikov & Henning 1997). There are no observations of molecular ∼lines towards this source which probe densities > 107 cm−3, so we presently have no solid understanding of the structure and dynamics in the highest density regions. Opacities in the millimetre and submillimetre are predicted to be large, τ 1 at λ1mm. ∼

3.3.2 Protostellar Jets Strong mass loss occurs in star formation, in which processes occuring very near the protostellar object (accretion, rotation, magnetic fields) drive large scale bipolar outflows of mass. Driven probably by strong ionized winds which are produced and collimated very near the star/accretion disk (tens of AU scale). Recent studies of heavily obscured YSO’s at centimeter wavelengths have revealed very weak thermal radio continuum jets on AU scales in a large fraction of objects. Approximately 80% of what are currently thought to be the youngest objects (extreme class“I” sources or class “O” sources) have been detected in the centimeter range (Anglada 1996). Only a fraction of the centimeter sources have been resolved. In those cases the emission is in the form of thermal radio jets from collimated, partially ionized flows with dimensions of 10-100 AU and dynamical time scales of order 1 year. We need to understand how they are driven and collimated, how fast they are, how massive, and how mass outflows change as the star evolves. Of particular interest is the question of how the mass loss varies on the shorter timescales years to tens of years and its connection with short term variations in the protostellar object and/or episodic disk/accretion events. Some jets appear to be time variable up to the extent that the jets may be pulsed with monopolar phases. What role do magnetic fields play in driving and collimating the flows? Which objects precess, and why? 106 CHAPTER 3. FORMATION AND EVOLUTION OF STARS

Figure 3.13: Approximate range of lengths of collimated outﬂows from Young Stellar Objects as observed at diﬀerent wavelengths (after Anglada 1996).

To answer such questions we need high angular resolution in order to resolve the jets and to be able to do variability and proper motion studies on them. The combined high resolution and high surface brightness sensitivity provided by the SKA will make it one of the most powerful tools for studying the nature of these deeply embedded objects, and thus understanding the final stages of the star formation process. The thermal jets are aligned with the larger scale molecular and optical outflows. Figure 3.13 shows a schematic illustration of the range of scales of the outflows that are seen around YSO’s and the physical scales that will be probed by the SKA. Ionized jets such as those associated with L1551 (see Fig. 3.14), will be probed on sub-AU scales! The ionized jets die out as protostellar objects approach the Main sequence, so we need the sensitivity of the SKA in order to trace the long term evolution of the mass outflow phase. Multifrequency observations are required in order to determine spectral indices and thus probe jet fine structures (such as opening angles and temperatures). The number of resolved protostellar jets, which emit strongest at centimetre wavelengths, is only about 20, and of those only about 10 are reasonably well studied, but there are at least 200 molecular outflow sources known, not to mention many more older pre-Main sequence objects without detectable molecular outflows but which have weak ionized winds. A typical jet may have a flux of < 1 mJy, and the VLA has the sensitivity to detect winds in only about 10% of nearby young stellar objects. The SKA will be able to detect and resolve these winds in essentially all nearby low mass protostellar and pre-Main sequence objects. Proper motions have been detected in a few objects, most notably in HH80-81 (1000 km/s) (see Fig 3.15), the Serpens jet source (300 km/s), and HH1-2 (400 km/s), with evidence that some jets are unipolar (HH111) or bipolar but asymmetric 3.3. STAR FORMATION 107

Figure 3.14: A λ3.6 cm VLA image of double? jets in L1551. The yellow rectangles indicate the position and size of two protoplanetary disks found a λ7 mm (Rodriguez et al. 1998). With the sensitivity and resolution limits of the VLA the nature and origin of this ionized gas cannot be uniquely deﬁned. The SKA will allow such structures to be probed at sub-AU scales.

S4 N4

1994.3-1990.2

1995.5-1990.2

1997.0-1990.2 E

N 1"

Figure 3.15: Multi-epoch images of the ionized jet of HH80-81 at λ3.6cm. The three images are diﬀerence maps from an image taken earlier at epoch 1990.2. High resolution images at low ﬂux levels allow detection of faint components of the jet and measurements of proper motion. From Marti, Rodriguez, and Reipurth (1998). 108 CHAPTER 3. FORMATION AND EVOLUTION OF STARS

(Serpens, Re 50). These must be intrinsic effects, since variable dust extinction cannot affect these radio continuum observations. Current instruments which have some capability to resolve proper motions in ionized protostellar jets (VLA, VLBI) can observe only a handful of the fastest, nearest, and brightest jets. Confirming proper motions of 0.1′′ year−1 with the VLA requires observations spaced over a 10 year period, while∼ the brightness sensitivity of VLBI observations is rather poor. ∼ The SKA will excel in observing in essentially “real time” the ejection, development, and interactions of the jets which act as a dynamical and energetic interface between young stars and their circumstellar environments over many decades of size scale. SKA will also be able to detect and map for the first time radio recombination lines from hydrogen in these jets. The understanding of how these jets are accelerated and collimated, most probably by magnetohydrodynamical mechanisms, requires of a knowledge of the kinematics of the ionized gas. Only the SKA will have the sensitivity to provide this information.

3.3.3 Uncovering the Evolutionary Sequence One of the key stages in stellar evolution is the period just prior to the formation of a protostar when a cloud core achieves the critical state that transforms it into a collapsing object. The conditions which lead to this transformation, and the processes by which it occurs, determine how solar systems form and how galaxies evolve. In the commonly presented scenario, a proto-stellar nebula forms out of one of the fragments of a collapsing cloud core. Observations of dense regions have in come cases shown evidence for a “layered chemistry”, or chemistry that varies greatly through a section of a dense cloud. This could be an excitation eﬀect due to increasing density towards the centers of knots, or an “age” eﬀect in which a time dependent chemistry both creates and destroys various species. Or it could be time-dependent or density-dependent depletion on to grains. The structure of these regions also indicates fragmentation, which could lead to the coagulation of fragments to form proto-stellar systems. The layered separation of two important species is illustrated in Fig. 3.16, where thorough observations of NH3 and CCS have been carried out for the dense cloud L1498. In this case Kuiper et al. (1996) argue that this pre-protostellar cloud shows evidence for growth by accretion. Ammonia, which takes about 106 year to form (Herbst et al. 1989), is concentrated near the center. CCS, on the other hand, forms early in a chemically evolving cloud, and is destroyed within a few 105 yr (Millar × and Herbst 1990). Although only the ends of L1498 could be mapped with high angular resolution, the pattern is strongly suggestive of layering. The dense ridge in Taurus MC1 illustrates the value of the SKA in studying star forming clouds (Fig. 3.17). Over a region about 20 arcmin in size, an evolutionary sequence appears to be laid out. The ridge in TMC1, however, appears to be quite unusual (at least in terms of its extraordinarily rich chemistry), and may not be indicative of typical cloud-collapse circumstances. At the NW end there is a con- 3.3. STAR FORMATION 109

Figure 3.16: L1498 is an example of a quiescent core, possibly in a pre-protostellar phase. Ammonia, a molecule which forms slowly, is concentrated near the center, indicating the presence of relatively old gas. CCS, a molecule with a short (few 105 yr) lifetime, is concentrated in an outer layer. This suggests that the cloud is growing× by accretion.

Figure 3.17: Taurus Molecular Cloud 1 shows a linear progression from young CCS- rich gas in the SE to older gas with prominent NH3 in the NW. (Adapted from original Hirahara et al. 1992.) 110 CHAPTER 3. FORMATION AND EVOLUTION OF STARS

centration of NH3, far-infrard, and outflow sources (Chandler et al. 1996). At the SE end, CCS predominates. There are five condensations, labeled A through E, along the cloud, and star formation is evident only near core A. If this is an age effect, it seems likely that, in time, star formation will occur, probably sequentially, in at least some of the other condensations. However, with the best present instrumentation, it took several years and several hundred hours of telescope time to observe Core D alone. Although CCS illustrates the layering effect particularly well in L1498 and TMC1, other species have also been found to vary along the TMC1 ridge(e.g. C4H and HC7N (Olano etal. 1988)). A quantitative explanation must be based on the statistics of many clouds and cloud-clumps that constitute the low mass end of the structure of the interstellar molecular medium. Sensitive (T ∗ 0.1 K) spectral line maps of molecular clouds, A ∼ with high spectral resolution ( 0.05 km s−1) and good spatial resolution, are required to provide data on the density,∼ mass, and temperature of the fragments, the space density of the fragments and their relative velocities. Comparison of the abundances of key molecular species in many objects could elucidate which of the possible processes is producing the layered effects. A large set of cases is needed to provide information on the ages of fragments, and to yield evolutionary sequences as the abundances respond to the changing conditions during the pre-star-formation assembly process. Observations of this type are best done at centimeter wavelengths. These weak lines may be one of the best hopes of tracing star formation. At these low temperatures ammonia only emits significantly in a few transitions near 24 GHz. CCS and similar carbon chain molecules, because of their large moments of inertia, radiate predominantly at centimeter wavelengths. In dark cloud cores, CCS and NH3 abundances are anti-correlated, with NH3 abundant in cores with signs of star formation, and CCS is abundant in cores without star formation (Suzuki et al. 1992). Thus it currently appears that these are two key molecules which probe the beginning and end of the star formation evolutionary sequence. Existing telescopes and arrays are not well suited to measuring on the relevant size scales. The largest single apertures do not have enough resolution, and the arrays cannot achieve the required aerial coverage and brightness sensitivity. As a phased array the SKA would take about 10 minutes to integrate to an r.m.s noise level of 20 mK with a spectral resolution of 0.05 km/s in the 22 GHz band and an angular resolution of about 10′′. This would, for example, enable a map of the 3 4 arcmin × area of L1498 to be carried out in the NH3 and CCS lines near 24 and 22 GHz, respectively, in less than an hour.

3.3.4 Magnetic Fields in Frotostellar Objects Zeeman Splitting in Molecular Gas Polarisation observations of molecular spectral lines with the SKA will yield measurements of magnetic ﬁeld strengths in dense molecular regions via the Zeeman eﬀect. 3.3. STAR FORMATION 111

At centimetre wavelengths the Zeeman effect has been observed in the interstellar medium for both the H I line and for the OH radical. The H I line does not, however, serve as a good magnetic field probe of primarily molecular regions. Neither does OH, since its abundance is much higher in the less dense envelopes around molecular gas than in the interior regions, and OH maser observations sample only a particular type of environment inside a cloud whereas a general probe is needed for mapping fields in dense molecular clouds. In the millimetre and submillimetre range, possible Zeeman probes include the SO, CN, CCS, and CCH molecules. However, there are strong advantages to be gained by making Zeeman observations in the centimetre wave lines of appropriate molecules: the splitting of the line into Zeeman components is approximately independent of the line frequency whereas the Doppler width of lines is proportional to frequency. Therefore, the ratio of Zeeman splitting to Doppler width (and hence the ability to detect the Zeeman effect) is greater for the lower frequency lines. Potential probes of the Zeeman effect in low frequency lines include SO (13 GHz) and CCS (11 GHz, 22 GHz). The transition which has the greatest sensitivity to this effect is the CCS line near 11 GHz. The measurements require high sensitivity and high spectral resolution to make very precise determinations of line profiles, and high angular resolution to resolve the magnetic structure in protostellar cores. Present efforts to observe this effect with single aperture telescopes involve many tens of hours of integration time for one position. As an aside, attempts to detect the Zeeman effect in hydrogen and carbon radio recombination lines have not yet been successful, and the SKA will be the most powerful instrument for renewed attempts.

Non-Thermal Emission Processes Classical T Tauri stars (CTTSs) have weak radio emission detected in about 10% of objects, 0.3 mJy at 5 GHz. All these seem to be associated with jets/collimated ∼ outflows. Herbig Ae/Be (HAEBE) stars, which are intermediate mass CTTSs, are detected for about 20% of all nearby objects. All of the above are thought to be dominated by thermal emission from ionized winds. In contrast, some weak lined Tauri stars (WTTSs) and related Class III sources are detected in the radio continuum much more often (up to 50% of objects) at 1 mJy at 5 GHz. Such objects ∼ have only remnant disks and weak ionized winds at best, and emit with nonthermal characteristics: they are variable on timescales of hours to days, have a moderate degree of circular polarization (several percent), and high brightness temperatures 107 K. The emission is thought to be due to the gyrosynchrotron mechanism (electrons∼ moving around large scale, dipolar-like stellar “magnetospheres” up to 30 solar radii in diameter). The radio spectra are quite flat, with indices 0 during quiescence and 1 during outbursts at 1–5 GHz. At higher frequencies,∼ 5–15 GHz, the opposite behaviour∼ is seen, suggestive of a turnover in the spectrum around 5–10 GHz during flares. Little is known about the frequency/time dependence of the circular polarization of the nonthermal emission from these objects. The current lack of detections of nonthermal emission from CTTS and younger Class I and Class 0 objects (in contrast to Class III/WTTS objects) brings up the 112 CHAPTER 3. FORMATION AND EVOLUTION OF STARS

question of what are the magnetic field strengths and structures around these very young objects and how do they evolve? Many theoretical models of protostellar outflows require strong, 1 kGauss, fields near the stellar surface. It has also been ∼ suggested that strong fields in CTTSs may couple the star to its accretion disk and thus provide a way for the star to regulate its angular momentum. This regulation is needed in order to keep CTTSs rotating well below breakup speed even though they are accreting high angular momentum material from their disks, and to explain why CTTSs (generally younger) rotate only half as fast as WTTSs (generally older). The SKA sensitivity would be essential for the concerted searches for nonthermal emission from CTTSs. The above “theoretical” requirements of CTTS magnetic fields can probably be satisfied with field strengths roughly similar to those of the nonthermally emitting WTTSs. There is, however, no radio evidence of large scale magnetic structures in very young objects. These objects have ionized winds/jets which, in a few cases at least, may have appreciable optical depths in their thermal free-free emission. It is thus possible that nonthermal emission from very young objects may be partially or completely absorbed in the ionized gas. What fraction of the nonthermal emission that would be masked will depend upon the size of the magnetosphere with respect to the scale length of the ionized jet gas. The SKA will have the sensitivity to go two orders of magnitude deeper than current centimetre wave telescopes in the search for spectral signatures of partially attenuated nonthermal emission. A key method to find evidence for the existance of magnetospheric-type structures around jet sources may be to look for short period variability and/or circular polarization at high angular resolution and sensitivity. A demonstration of this possibility is the polarized radio emission from around T Tauri: evidence for magnetic fields in young jet-driving sources can be detected when they extend outside of the inner obscuring free-free “blanket”, as seen from the MERLIN λ6cm map of this system. The near infrared companion to T Tauri is observed to have distinct lobes of right and left circularly polarized emission around the infrared source, separated by 20 AU, suggestive of ∼ magnetic structures extending on tens of AU scales. The most likely explanation of this phenomena is that the magnetic fields (a few Gauss) are part of a collimated flow from the star. The sensitivity of SKA will be needed to study similar phenomena in other young stellar objects.

3.4 Cool Star Astronomy

Cool stars play a major role in astrophysics; they deﬁne the largest stellar class, comprising our Sun as an average example. Most cool stars maintain magnetically conﬁned atmospheres which give rise to particle acceleration and plasma heating. Investigations of the Sun in spatial, spectral, and temporal detail have provided us with a considerable knowledge on energy release, structuring, and evolution of stellar atmospheres. However, the Sun represents a particular state of stellar evolution, for a particular stellar mass. Understanding the full range of phenomena related to 3.4. COOL STAR ASTRONOMY 113

stellar activity, mass loss, and evolution requires the study of solar-like phenomena in large samples of stars. Up to the present day, no star at an activity level of the Sun’s, believed to be typical for the vast majority of cool stars, has been detected at radio wavelengths. The SKA will drastically change this situation and is certain to yield a major breakthrough in our understanding of stellar atmospheres, by detecting several thousand normal stars in the solar vicinity. A broad range of diagnostic tools is available. “Solar-stellar connection” studies will be fruitful means to better understand both the Sun and stars. Cool stars (typically defined as stars with photospheric temperatures < 7000 K, or spectral types later than A) serve as the natural bridge from solar to stellar∼ physics. On the one hand, the Sun is a moderate, average example of a middle-aged main- sequence star, in no ways special nor extreme, believed to well represent the basic phenomena of magnetic activity that are thought to occur in many other star classes with outer convection zones. On the other hand, mid-mass main-sequence stars play, among stellar populations, a major role in our understanding of the dynamo, the long-term evolution, the atmospheric energy release from magnetic fields, and the spin-down history of stars due to transport of angular momentum away from the star by a magnetized stellar wind. The stellar astronomer’s vantage point with a prototypical source close-by sets a particularly exciting task to this part of astronomy. While extremely detailed physical models can be tested on the Sun in spectroscopic, spatial, and temporal detail, the diversity of stellar sources plays also a global role in the evolution of the galaxy, the recycling of matter and enrichment of the interstellar medium, the formation of planetary systems, and perhaps in the acceleration of soft cosmic rays. The Sun presents only a small range of atmospheric behavior; stars with different masses, ages, rotation periods, and chemical composition need to be understood as well to attack many fundamental problems of stellar evolution and galactic structure. It is the task of the ‘solar-stellar connection’ to build the bridge between the great knowledge gathered from solar observations and modeling, and the diversity of stars. The flow of information is in both directions: Many of the physical principles are known from the Sun, but much about global systematics and stellar (and therefore: solar) evolution can be learned only from stellar observations. This section summarizes how far we have gotten in this endeavour in radio astronomy, and how important a role the SKA will play in solar-stellar connection studies. Single A or early F stars are usually not detected as radio sources (the same holds for X-ray emission; see gap in Fig. 3.1, bottom right, around B V 0.2). They do not possess strong winds, and their outer convection zone becomes− too≈ shallow to maintain a solar-like magnetic dynamo. Notable exceptions are chemically peculiar Bp/Ap stars that do support strong magnetic fields and are radio sources, although the physical mechanisms and geometric structures are probably very unlike the Sun’s (Linsky, Drake, & Bastian 1992). We are in the following mainly concerned with cool stars with outer convection zones and their pre-main-sequence relatives. There is evidence that each cool MS 114 CHAPTER 3. FORMATION AND EVOLUTION OF STARS

star possesses an outer magnetic atmosphere (e.g., Schmitt 1997) in which radio phenomena described in this chapter are generated. Toward the optically faintest stars (M dwarfs to the lower right in the HRD) the maximum radio luminosity declines. This trend is not due to a selection effect but is real and parallels observations in the X-rays: The total (non-flaring) X-ray luminosity of a coronal star is bounded by 10−3 times its bolometric luminosity. This upper bound is ascribed to a ‘saturation’∼ effect either intrinsic to the magnetic dynamo, or to the X-ray emitting plasma trapped in the corona. In contrast to the X-ray emission, the radio emission apparently signifies a saturation of the number of relativistic electrons in the corona. The statistics therefore clearly indicate that the most active coronal stars are found in spectral class G, i.e., among analogs to the Sun!

3.4.1 The Radio Sun

The non-flaring Sun has been studied at metric/decimetric/centimetric wavelengths in appreciable detail. Fig. 3.18 shows the Sun observed near its activity maximum with the Very Large Array (VLA) at 20 cm wavelength (Dulk & Gary 1983). Figure 2b illustrates a close-up view of 6 cm radio emission (contours) in the vicinity of a sunspot group (Gary & Hurford 1994). The emission contributions vary strongly during the magnetic activity cycle; at maximum, most of the measured flux stems from active regions. The dominant emission mechanism depends on the observing frequency and on the source location with respect to active region magnetic field structures. The type of dominant emission can be derived in particular from the spectral behavior on the optically thin side of the spectrum, with gyroresonance emission showing a steeper decrease toward higher frequencies; polarization properties are also used as discriminators (e.g., Gary & Hurford 1994; Vourlidas & Bastian 1996). Roughly speaking, the longer-wavelength emission (e.g., at 20 cm, Fig. 2a) is dominated by optically thick coronal free-free emission that correlates well with soft-X-ray features (identical emission mechanism; Dulk & Gary 1983; Gopalswamy, White, & Kundu 1991; Vourlidas & Bastian 1996). At wavelengths < 6 cm, active region emission consists of bright compact sources and a diffuse halo.∼ The former, strongly sunspot- related component with a brightness temperature of a few times 106 K is attributed to optically thick gyroresonance emission of coronal plasma in rather strong ( 1 kG) magnetic fields (e.g., Lang et al. 1987; Vourlidas, Bastian, & Aschwanden∼ 1997). The halo component is often due to optically thin thermal free-free emission from the corona, but also from the cooler transition region and the chromosphere. Both types of emission can be effectively used to derive principal source parameters, namely the electron temperature from either of the optically thick emissions, the electron density ne from the bremsstrahlung components, and the magnetic field strengths from the gyroresonance emission (Gary & Hurford 1994). Somewhat surprisingly, strong magnetic fields up to 2 kG can be present even at coronal levels. Such values are close to photospheric levels; they indicate that the magnetic field 3.4. COOL STAR ASTRONOMY 115

Figure 3.18: The 20 cm radio Sun observed with the VLA during the magnetic activity maximum on September 26, 1981. The patches of enhanced emission are predominantly due to optically thick free-free emission from coronal plasma. (Copyright by the Smithsonian Institution). Right: Radio emission (contours) from above a sunspot group (grayscale plot), observed at 6 cm (4.8 GHz) with the Owens Valley Radio Observatory. Much of the emission is due to optically thick gyroresonance emission in concentrated ( 1 kG) magnetic ﬁelds. (Gary & Hurford 1994; courtesy of D. ∼ Gary.) 116 CHAPTER 3. FORMATION AND EVOLUTION OF STARS

Figure 3.19: Extract from a dynamic spectrum of a solar type U radio burst group (9 July 1980). U bursts are generated by the beam plasma instability. The emission at the plasma frequency traces closed coronal loops. (From the RAG data archive, Institute of Astronomy, ETH Z¨urich.)

divergence from the photosphere to coronal layers is rather small (White, Kundu, & Gopalswamy 1991; Lang et al. 1987; Lang et al. 1993; Vourlidas et al. 1997). This aspect may be much more important still for active stars in which the photospheric filling factor of strong magnetic fields is large, and their divergence is even more confined on geometric grounds. During flares, the Sun can be a source of copious radio gyrosynchrotron emission for up to tens of minutes. The emission is produced by accelerated, mildly relativistic electrons spiraling around the magnetic field lines. This radiation often reaches its peak flux (transition from optically thick to optically thin) in the 1–10 GHz region. Its spectrum contains rich diagnostic information: The electron energies can be derived from the optically thick portion, magnetic field strengths from the turnover frequency, and the electron energy distribution from the optically thin part. Many of these diagnostic tools have been applied to easily observable, strong stellar flares, albeit with considerably smaller temporal and spectral (and no spatial) resolution. For a comprehensive review of solar radio flares and their diagnostic power, we refer to Bastian, Benz, & Gary (1998). Solar radio astronomy has greatly profited from the diagnostic power of coherent radio emission during flares (Bastian et al. 1998). Traditionally, dynamic spectra (i.e., flux as a function of frequency and time) at a high temporal sampling rate (>1 spectrum per second) with high frequency resolution (f/∆f > 100) over a broad frequency band (several 100 MHz) have helped disentangle a veritable∼ zoo of coherent 3.4. COOL STAR ASTRONOMY 117

radio emissions. The best studied examples are thought to be due to emission near the plasma frequency or its second harmonic, induced by an electron beam instability (type III). Since the emission frequency is proportional to √ne, the time history of the emission on the dynamic spectrum traces the path of the electron beam in density, in cases producing ‘pseudo-images’ of coronal loops as the electrons travel along closed magnetic field lines (Fig. 3.19). Beams traveling along open magnetic field lines eventually end up in interplanetary space. The timing of these bursts provides information on the electron beam acceleration and the possible fragmentation of the accelerator. In cases where simultaneous upward- and downward-moving beams can be identified, information about the density in the acceleration region (represented by the demarcation frequency) are obtained. Shorter, narrow-band emissions (spike bursts) have been interpreted as signatures of energy release fragmentation and may provide insight into the shortest time scales of the energy release (Benz 1985). Meter- wavelength type IV and type II bursts can be associated with coronal mass ejections (CME’s) and related shocks and therefore provide important clues on transient coronal mass loss and magnetic field geometry (Kundu et al. 1989).

3.4.2 Observing Solar Analogs at Radio Wavelengths

−1 −1 The non-flaring radio luminosity of the Sun at several GHz is logLR 10.7 [erg s Hz ] (principally due to gyroresonance and bremsstrahlung emissions)≈ and therefore well below the base level in the logarithmic 3-D diagrams in Fig. 1. No single radio star with LR equaling that of the Sun has ever been detected. The lowest-luminosity radio star found so far, the F5 subgiant Procyon at a distance of 3.5 pc, is at logL 11.7, R ≈ an order of magnitude higher than the Sun (Drake, Simon, & Brown 1993). Presently, the VLA detects a source like the Sun marginally out to 1.5 pc after 10 hours of integration time. The SKA, in contrast, will detect solar twins at 50 pc distance! Almost all radio detections among cool MS or subgiant stars are interpreted as nonthermal gyrosynchrotron emission, based on their shallow spectra (Fig. 3.20) and on estimates of their brightness temperatures (often reaching 108 1010 K; e.g., White, − Kundu, & Jackson 1989). This sets active stars apart from solar behavior. But since the Sun is also a gyrosynchrotron source during flares, it has been speculated that the steady nonthermal emission from magnetically active stars is the envelope of a large number of unresolved flare contributions. This view is supported by a correlation between quiescent radio and X-ray luminosities of active stars that appears to be similar for time-integrated solar flare gyrosynchrotron and X-ray emissions (Güdel 1994 for a review). In rare cases, a second spectral component is detected that is compatible with optically thick gyroresonance emission from hot plasma (Fig. 3.20, right panel). It can be used to derive coronal magnetic field strengths (Güdel & Benz 1989; White, Lim, & Kundu 1994). However, most dMe stars do not show this component at a detectable level, providing stringent upper limits to the filling factor of strong, low- coronal magnetic fields that contain hot plasma (White et al. 1994). 118 CHAPTER 3. FORMATION AND EVOLUTION OF STARS

Figure 3.20: Radio spectra of the active dMe star UV Cet during quiescence. The left figure shows three optically thin gyrosynchrotron spectra. The right figure additionally reveals optically thick components above 10 GHz interpreted as gyroresonance emission. (Güdel & Zucker 1999.)

Flares on active stars are commonly interpreted in terms of solar flare physics (Bastian 1990 for a review). Their emission is often strongly circularly polarized and rapidly varying in time. Together with the high brightness temperature (up to 1016 K) this clearly suggests coherent emission mechanisms. Dynamic spectra would therefore provide great diagnostic insights into the timing and temporal fine structures of the accelerator, electron beams (type III or U bursts), mass loss by coronal mass ejections (type II and IV bursts), magnetic field topology, or electron densities. However, non- solar radio astronomers have traditionally been utilizing single-frequency receivers. Only a few dynamic spectra of stellar flares exist so far; their information content is clearly limited given their total bandwidth of typically less than 100 MHz at a frequency of 1400 MHz (Fig. 5; see also Bastian & Bookbinder 1987; Jackson, Kundu, & White 1987; Bastian et al. 1990). A broadband device such as proposed for the SKA could produce dynamic spectra routinely.

3.4.3 Where are the many other Radio Suns? Why do we have numerous nonthermal radio detections of nearby M dwarfs (with smaller surface areas than the Sun; White, Jackson, & Kundu 1989) while solar analogs have escaped detection until recently despite their clear trend toward a much higher upper envelope of radio luminosity (Fig. 3.2)? Presently, only about a dozen GV stars have radio detections (see list in G¨udel, Guinan, & Skinner 1998). Ob- servations show that the spin-down of initially rapidly rotating G stars to projected equatorial surface velocities of vsini 10 km s−1 occurs within a few tens of Myr, ≈

3.4. COOL STAR ASTRONOMY 119

{

{ MHz

requency

{

j j

Time UT on Nov

Figure 3.21: Dynamic spectrum of a very strong ( 700 mJy) coherent radio burst from the dMe star AD Leo, observed simultaneously∼ at Arecibo (top; time resolution = 0.2 s), Effelsberg (middle; 0.125 s), and Jodrell Bank (bottom; 2 s). Note the narrow frequency range. The Effelsberg range is contained within the Jodrell Bank band. (From Güdel et al. 1989). 120 CHAPTER 3. FORMATION AND EVOLUTION OF STARS but takes 10 times longer for M dwarfs (Soderblom et al. 1993). At the Hyades’ age (600 Myr),∼ G stars rotate slowly, while M dwarfs easily maintain vsini> 10 km s−1. The value of vsini is a good measure for magnetic activity (presumably since rotation directly influences the operation of the magnetic dynamo). Further, the space density of field M dwarfs is 8 times higher than for G dwarfs (Allen 1973). Both properties strongly favor M dwarfs for radio detections. We thus identify two major areas in which the SKA will lead to a breakthrough: 1) The most “active” single MS stars are found among solar analogs. These would be the objects of choice for many investigations were they not so rare or weak. The SKA will detect such stars easily out to several hundred pc. 2) By far the largest population of stars in the solar neighborhood is made of spun-down stars that presumably behave grossly like the Sun. Except for the weak detection of the subgiant Procyon, none of these stars has been found at radio wavelengths so far.

3.4.4 Prospects for the SKA Physics of Solar-Like Stars The SKA will detect solar twins out to 50 pc as illustrated in Fig. 3.22. Extrap- olating from the Gliese & Jahreiss (1991)∼ Catalog of Nearby Stars or Allen’s (1973) volume densities of stars in the solar neighborhood, we infer that 30,000 stars (99% of them being cool MS stars) are potentially within reach of a telescope like the SKA if the solar activity level is taken as a lower limit for all MS stars. Clearly, the SKA will open a completely new window to ‘normal’ stars. Numerous physical problems so far studied only on the Sun can be investigated on a wide range of stars; the diagnostic power of bremsstrahlung or gyroresonance components will allow us to infer low-coronal magnetic field strengths, magnetic filling factors, or coronal electron densities and to study these parameters as a function of stellar age, rotation period, mass, or other magnetic activity indicators. The potential for infer- ring structuring and composition of the (non-flaring) corona and transition region will be considerable. Radio observations will help determine coronal temperatures from the low-coronal optically thick emission and will therefore complement soft X-ray observations. The magnetic filling factor, if identified from radio observations, should be compared with spot coverage maps from simultaneous optical Doppler imaging techniques. Rotational modulation can be studied at high sensitivity. Apart from providing a unique collection of rotation periods (difficult to measure otherwise), it gives invaluable information on the size and the distribution of active regions. Knowledge of the frequency-dependent source sizes is important for a proper interpretation of the microwave spectra, including variations of the polarization (Bastian et al. 1998). Since the presence of active regions presumably varies with magnetic activity cycles as on the Sun, radio observations will also uncover such cycles. Polarization measurements could perhaps even detect a predominant magnetic polarity! Nonthermal sources will be detectable out to hundreds of parsecs. This provides 3.4. COOL STAR ASTRONOMY 121

the exciting opportunity to study several open stellar clusters at radio wavelengths (Lim & White 1995). Such studies are important to probe systematics in the evolution of coronae. This extends back in time to many star forming regions of which the most important ones will be easily surveyed by the SKA. If the SKA will be connected with VLBI antennas or will itself contain baselines of Earth-diameter size, its sensitivity will make routine, spatially resolved observations of radio stars at the milliarcsec level possible. So far, only few main-sequence stars reach an appropriate level for VLBI studies, but they begin to yield exciting results on coronal structure, such as global magnetic ﬁelds, in active stars not expected before (Benz, Conway, & G¨udel 1998).

3.4.5 Flares and Microflares One major aspect of the SKA is its ability to provide high time resolution thanks to its high sensitivity. Coronal energy releases in the Sun and on stars occur on time scales between less than 0.1 s and several tens of minutes. Although the VLA and other modern radio telescopes observe large, energetic flares from many stellar systems, the main physical questions require high time resolution of flare light curves, broad bandwidth with spectral resolution, and observations of small and short events (‘microflares’). The limitations to the latter are illustrated in Fig. 3.22. The two upper panels show the light curve from a VLA X-band observation (100 MHz bandwidth) of one of the most productive, most variable, and radio-brightest nearby flare stars, the M dwarf binary UV Cet. About ten flares with characteristics similar to major solar “gradual” flares are well visible (Güdel et al. 1996). From simultaneous X-ray observations, their radiated energy has been estimated to be2 1030 2 1031 ergs, i.e., corresponding to large solar flares. × − × The thermal noise sets severe limitations. The remaining panels show a simulation of a synthetic light curve (bottom panel) that conforms with the assumption of small- scale events building up part of the “quiescent” emission of active stars. About 15,000 flares with a power-law distribution of their energy and duration have been stochastically placed in time (on a constant background level). The middle panels in Fig. 3.22 show the input light curve degraded to the sensitivity of the VLA at X-band (1 min resolution) and to the SKA (in the 5 GHz band with maximum bandwidth). We expect that SKA easily observes flares on nearby stars 100 times smaller in their total energy (and shorter by a factor of ten) at time resolutions of a few seconds. We note that a typical solar radio flare may reach about 0.1 mJy at a distance of α Cen (1.3 pc). Such fluxes are accurately measurable at sub-second time-resolution with the SKA! The great sensitivity to flares over a very wide range in energy will address the fundamental problem of coronal heating. The energy distribution of a large set of flares can be used to study whether their rate of occurrence suffices to power the corona (the “microflare/nanoflare” hypothesis; see Parker 1988). Observations of the Neupert effect (characteristic lag of soft X-ray emission with respect to the flare gyrosynchrotron peak, owing to a causal relation between electron/heat front impact 122 CHAPTER 3. FORMATION AND EVOLUTION OF STARS

Figure 3.22: Expected sensitivity range of the SKA for detecting cool stars. A number of stars are plotted on the distance-luminosity plane, with the blue dots representing selected bright objects (mostly RS CVn binaries; Drake, Simon, & Linsky 1993) and the red dots F/G-type MS stars so far detected. Many M dwarf flare stars fall within the schematic quadrangle at left (within 10 pc; White et al. 1989). Various other radio detections (Wendker 1995) are plotted as small crosses. The average radio luminosity at 8 GHz of the non-flaring Sun is indicated by the bar at bottom, and the luminosity of a typical solar radio flare is marked by the dotted line. The diagonal lines represent the 3σ detection limits for three different continuum integration times (1 min, 1 hr, 10 hrs), both for the VLA (8.4 GHz, 100 MHz bandwidth; upper set) and the projected SKA (maximum bandwidth; lower set) – Right: Light curve simulation for the SKA. Top panel: VLA observation of the flare star UV Cet at X band, 5 min resolution; 2nd panel: same observation, 1 min resolution; 3rd and 4th panel: simulated light curves for the VLA X band and the SKA, respectively, using the input model light curve shown in 5th panel. The numbers in the plot indicate the time binning. 3.4. COOL STAR ASTRONOMY 123 on the chromosphere and plasma heating and expansion) in a large number of such events may yield a qualitative or even quantitative picture of mechanisms that lift cool material into the strongly radiating coronae of magnetically active stars.

3.4.6 Summary of Scientific Objectives The scientific objectives in stellar astronomy for an instrument with the capabilities of the SKA are manifold; most of the stellar phenomena described above and known on the Sun will for the first time be within reach of a radio telescope. Many scientific objectives so far only addressed in the solar case will be extended to a large variety of stars. Some of the major steps forward are expected in the following areas:

Routine detection of non-ﬂaring, thermal radio emission from most cool • stars in the solar neighborhood, out to approximately 50 pc distance. The SKA will probe the physics and structuring of the lower corona and transition region. This telescope will open a new research area in the ‘radio solar-stellar connection’.

Bremsstrahlung and gyroresonance components will yield information on the • electron temperature and density and on the magnetic ﬁeld strengths in active regions for a large sample of stars. In the case of more active stars, optical Doppler imaging can be used to relate the observed coronal emission to surface magnetic features. Rotational modulation provides information on the size and distribution of active regions, while long-term monitoring could uncover magnetic activity cycles.

By studying the production of nonthermal electrons via their gyrosynchrotron • emission in active stars and its fading toward less active stars, the physics of energy release can be studied as a function of stellar evolution. This is relevant to the problem of magnetic ﬁeld production, the stellar dynamos, and the interaction of high-energy particles with the stellar environment. In terms of solar history, we can study the ‘Sun in Time’ and infer the behavior of the Sun in past times. Extension to pre-main-sequence objects will probe the possible early history of the Sun and provide important information on star formation.

Luminous nonthermal coronal sources such as young, active main-sequence stars • will be detectable out to several hundred pc easily. Many open clusters will be studied to infer important systematics in the evolution of coronal activity.

High-sensitivity observations will detect numerous small ﬂaring events. Their • statistics could yield important clues for the long-standing problem of coronal heating.

Broad-band observations will provide diagnostics for coronal instabilities and • energy release. This ﬁeld has strongly fostered solar radio astronomy for 50 124 CHAPTER 3. FORMATION AND EVOLUTION OF STARS

years but has only been touched by stellar astronomers, even after considerable eﬀorts.

The SKA can contribute greatly in sensitivity to VLBI studies of normal stars. • Such observations will be ground breaking for the study of coronal structure in many types of stars.

Stellar radio astronomy is in urgent need of a new instrument as powerful as the proposed SKA. About 20 years ago, the VLA became operational; it has set new standards in stellar radio astronomy by discovering many stellar classes routinely for the ﬁrst time, and thus yielding access to new physics at a rapid pace along with leading satellite observatories at other wavelengths. By now, however, most new investigations of stellar chromospheres, transition regions, and coronae rely on data in the ultraviolet, the EUV range, and X-rays from devices that detect moderately active stars out to many parsecs routinely, and active stars out to hundreds of parsecs. But radio astronomy has its own diagnostic power that cannot be replaced by information from other wavelengths. A new suite of space-borne observatories will soon be, or is already, available across the electromagnetic spectrum (e.g., AXAF, XMM, HST, GRO, SIRTF, FIRST), with next-generation follow-up missions already under discussion (NGST, Constellation-X, XEUS, etc). Clearly, radio astronomy needs to keep pace with these exciting developments.

3.5 Imaging of Circumstellar Phenomena

The new regimes of milli-arcsecond imaging parameter space depicted in Figure 2.19 open up a wealth of phenomena in stellar astrophysics that have as yet only been revealed as point sources of radio emission. Imaging with linear resolution of 1 AU will be possible out to distances above a kpc. As noted in the previous section,∼ this dimension is characteristic of the radii of red giant stars. This dimension is also typical of the binary separation of semi-detached interacting binary systems containing giant stars, such as symbiotic stars, the winds of Wolf-Rayet, OB and Be stars, planetary nebulae, and the very early stages of novae outbursts. At current resolving power, imaging investigations of these types of objects are constrained to examination of material that has propagated to large distances from the point of origin before becoming resolved. The SKA will ﬁnally allow imaging of thermal radio emission on the scale of the phenomena that give rise to these sources. The circumstellar disks of Be stars can be imaged out to several hundred pc; red giant photospheres are resolvable out to more than a kpc; the binary separation of ‘D’ type symbiotic stars, the ﬁrst week of nova ejecta evolution, and the winds of WR and OB stars can be resolved to several kpc. 3.6. STELLAR ASTROMETRY 125

3.6 Stellar Astrometry

Determination of distances is a fundamental problem in studies of objects within our Galaxy. Spectroscopic or kinematic distances have a high degree of uncertainty, and that uncertainty generally propagates to even higher levels of uncertainty in source physics. The large increase in the number of stellar detections with the SKA combined with milli-arcsecond resolution at centimeter wavelengths will enable astrometry on a large number of radio stars at sub-milli-arcsecond accuracy. This level of astrometric accuracy would allow accurate parallax measurements out to distances of a large fraction of a kpc. It would also allow the direct measurement of binary motions for nearby (within a few 100 pc) non-thermal radio binaries such as RS CVn and Algol systems and weak thermal emitters such as Be star binaries. One of the chief problems with using stellar positions to tie the optical and radio coordinate frames to very high accuracy is the binary nature and high degree of variability of the few bright non-thermal radio stars for which accurate radio astrometry is currently available. Thermal emission from stellar photospheres and the winds of single stars is quasi-steady state and centered on the position of the optical image of the star. These objects are ideally suited for matching optical and radio reference frames, a problem that will become more critical when coordinated optical and radio studies on sub-arcsecond scales become commonplace with the new generation of optical and radio facilities. Combined with the observations from missions like Hipparcos, radio astrometric positions for a large number of thermal emitting, single stars will provide an extensive grid of accurate measurements to tie the optical and radio reference frames over the entire sky with milli-arcsecond precision. Dimensions of order 1 AU have crossing times of order one year for velocities of 10’s of km s−1. Hence milli-arcsecond resolution images will allow direct measurements of the angular expansion of slowly expanding objects such as planetary nebulae and wind interaction shells with the ISM. Combined with spectroscopic data, such measurements can provide accurate distances for a signiﬁcant population of objects out to kpc distance scales.

3.7 Supernovae

Detailed radio observations of extragalactic supernovae are critical to obtaining valuable information about the nature and evolutionary phase of the progenitor star in the period of a few hundred to several tens-of-thousands of years before explosion. Ad- ditionally, radio observations of old supernovae (>20 years) provide important clues to the evolution of supernovae into supernova remnants, a gap of almost 300 years (SN 1680, Cas A, to SN 1923A) in our current knowledge. Finally, new empirical ∼ relations indicate that it may be possible to use some types of radio supernovae as distance yardsticks, to give an independent measure of the distance scale of the Universe. However, the study of radio supernovae is limited by the sensitivity and resolution of current radio telescope arrays. Therefore, it is necessary to have more sensitive 126 CHAPTER 3. FORMATION AND EVOLUTION OF STARS

arrays, such as the SKA, to advance radio supernova studies and our understanding of supernovae, their progenitors, and the connection to supernova remnants. Supernovae (SNe) play a vital role in galactic evolution, through explosive nucle- osynthesis and chemical enrichment, energy input into the interstellar medium, their stellar remnants (i.e., neutron stars, pulsars, and black holes), and the production of cosmic rays. SNe are also being utilized as powerful cosmological probes, both through their intrinsic luminosities and expansion rates. A primary goal of supernova research is an understanding of progenitor stars and explosion mechanisms for the different SNe types. Unfortunately, little is left of the progenitor star after explosion, and only the progenitors of three (SNe 1987A, 1978K, and 1993J) out of more than 1350 extragalactic SNe have been directly identified in pre-explosion images. Without direct information about the progenitors, thorough examination of the environments of SNe can provide useful constraints on the ages and masses of the progenitor stars. SNe come in three basic types (e.g., Filippenko 1997): Ia, Ib/c, and II. Both SNe Ia and SNe Ib/c lack hydrogen lines in their optical spectra, whereas SNe II all show hydrogen in their optical spectra with varying strengths and profiles (Schlegel 1996). SNe Ib and SNe Ic subclasses do not show the deep Si II absorption trough near 6150 A˚ that characterizes SNe Ia, and SNe Ib show moderately strong He I lines, while SNe Ic do not. These spectral differences are theoretically explained by differences in progenitors. SNe Ia are currently thought to arise from the total disruption of white dwarf stars, which accrete matter from a binary companion. In contrast, SNe II, SNe Ib, and SNe Ic are likely the explosions of massive stars. SNe II presumably result from the core collapse of massive hydrogen-rich supergiant stars with masses 8 < M(M⊙) < 40. On the other hand, SNe Ib/c are believed to arise from a massive∼ progenitor∼ which has lost all of its hydrogen envelope prior to explosion (e.g., Porter & Filippenko 1987). One candidate progenitor for SNe Ib/c is exploding Wolf-Rayet stars (which evolve from stars with M > 40 M⊙; e.g., Conti et al. 1983: Humphreys, Nichols, & Massey 1985). An alternative∼ candidate is exploding, relatively less-massive helium stars in interacting binary systems (Uomoto 1986; Podsiadlowski et al. 1992). Possible variant of normal SNe II are the “Type IIn,” (Schlegel 1990) and the “Type IIb” (Filippenko 1988), which both show unusual optical characteristics. SNe IIn show the normal broad Balmer line profiles, but with the narrow peak sitting atop the broad base. The narrow component presumably arises from interaction with a dense (n > 107 cm−3) circumstellar medium (CSM) surrounding the SN. SNe IIb look optically∼ like normal SNe II at early times, but evolve to more closely resemble SNe Ib at late times.

3.7.1 Radio Supernovae Radio emitting supernovae (RSNe) have been extensively searched for since at least 1970 (e.g., de Bruyn 1973; Allen et al. 1976), but, due to low resolution, background confusion, and sensitivity limitations, only with the Very Large Array (VLA) was 3.7. SUPERNOVAE 127

the first example found which could be studied in detail at multiple radio frequencies (SN 1979C; Weiler et al. 1981; see also Weiler et al. 1986, 1991, 1992a; Montes et al. 1998b). So far, about 22 RSNe have been detected since then with the VLA, with about 17 objects which have been extensively studied, including the SN II 1980K (Weiler et al. 1986, 1992b; Montes et al. 1998a), the SNe Ib/c 1983N (Weiler et al. 1986) and SN 1990B (Van Dyk et al. 1993b), the SNe IIn 1986J (Weiler, Panagia, & Sramek 1990) and SN 1988Z (Van Dyk et al. 1993a), and the SN IIb 1993J (Van Dyk et al. 1994). The SNe IIn are unusual not only in the optical, but also in the radio, in being exceptionally powerful radio sources ( 1028 erg s−1 Hz−1, or several thousand times the luminosity of Cas A, at 6 cm). ∼ Fig. 3.23 shows several examples of detailed RSNe light curves. Fig. 3.24 shows an example of a particularly well-measured RSN, the SN IIb 1993J, which is only 3.6 Mpc distant, in M81. Analysis of the radio emission provides vital insight into the interaction of the SN shock with preexisting circumstellar matter lost by the progenitor star or progenitor system, and, therefore, into the nature of presupernova evolution. All RSNe appear to have the common properties of 1) nonthermal synchrotron emission with high brightness temperature; 2) a decrease in absorption with time, resulting in a smooth, rapid turn-on first at shorter wavelengths and later at longer wavelengths; 3) a power-law decline of the flux density with time at each wavelength after maximum flux density (optical depth 1) is reached at that wavelength; and 4) a final, asymptotic approach of spectral index≈ α to an optically thin, non-thermal, constant negative value (Weiler et al. 1986, 1990). The observed RSNe can in general be represented by the Chevalier (1982a, b) “mini-shell” model, which involves the acceleration of relativistic electrons and enhanced magnetic field necessary for synchrotron emission, arising from the SN shock interacting with a relatively high-density CSM, which has been ionized and heated by the initial SN UV/X-ray flash. The CSM is presumed to be the pre-SN mass loss in the late stages of the progenitor’s evolution. The rapid rise in radio flux density results from the shock overtaking progressively more of the ionized wind matter, leaving less of it along the line-of-sight to absorb the emission from the shock region. The slow decline in flux density at each wavelength after the peak is then due to the SN shock expanding into generally lower density regions of the now-optically thin CSM. This model has been parameterized by Weiler et al. (1986, 1990) as:

− ′ α β −τ 1−e τ S = K1(ν/5GHz) (t t0) e τ ′ mJy, − τ = K (ν/5GHz)−2.1(t t )δ , and 2 − 0 ′ τ ′ = K (ν/5GHz)−2.1(t t )δ . 3 − 0 K1, K2, and K3 then formally correspond to the unabsorbed ﬂux density (K1), uniform (K2) and non-uniform (K3) optical depths, respectively, at 5 GHz one day after −τ the explosion date t0. The term e describes the attenuation of a local medium with optical depth τ that uniformly covers the emitting source (“uniform external 128 CHAPTER 3. FORMATION AND EVOLUTION OF STARS

SN 1979C SN 1983N 20 cm 10 10

20 cm 6 cm

1 2 cm 1 6 cm

10 100 1000 Days Since Explosion Days Since Explosion

3.6 cm 6 cm

1 2 cm

20 cm SN 1988Z

1000 Days Since Explosion

Figure 3.23: Light curves at multiple radio frequencies for SN 1983N in M83, SN 1979C in M100, and SN 1988Z in MCG +03 28 022. − − 3.7. SUPERNOVAE 129

Figure 3.24: Light curves at multiple radio frequencies for a recent, nearby Type IIb, SN 1993J in M81 (NGC 3031), only 3.6 Mpc distant.

′ absorption”), and the (1 e−τ )τ ′−1 term describes the attenuation produced by an inhomogeneous medium with− optical depths distributed between 0 and τ ′ (“clumpy absorption”). All absorbing media are assumed to be purely thermal, ionized hydrogen with opacity ν−2.1. The parameters δ and δ′ describe the time dependence of the optical depths∝ for the local uniform and non-uniform media, respectively. ′ ′ Normally 0 >δ>δ , so that τ is the dominant opacity when (t t0) 1/(δ−δ′) − ≤ (K3/K2) days. At later times, the dominant opacity is τ until the CSM becomes optically thin and the radio emission is described by its characteristic power law decline with index β. In both Fig. 3.23 and 3.24 we show the model ﬁts to the observed data. As more radio information has become available, some interesting variations in this model have appeared, including clumpiness in the CSM, variations in mass- loss rates (and, thus, stellar evolution phase in the last few thousand years before explosion), and possibly synchrotron self-absorption (SSA) in the earliest phases of the SN evolution (Chevalier 1998) – the only example of SSA known outside of compact galactic nuclei and quasars. However, even with the considerable improvement in VLA sensitivity over the past 20 years, the ﬁeld of RSN studies is still very much sensitivity limited. More than 100 nearby SN events have been observed in the radio, with only a detection rate of 1/4, and we have only been able to develop even partial, multi-frequency radio light curves for fewer than half of those detected. 130 CHAPTER 3. FORMATION AND EVOLUTION OF STARS

With more than 1300 SNe which have been discovered optically since the first modern SN discovery, SN 1885A (S Andromeda) in M31, there is insufficient radio sensitivity, even with the VLA, to have a chance of detecting even a small fraction of them. Such sensitivity limitations restrict the scope of most RSN studies to distances smaller than the Virgo cluster, a cosmologically insignificant distance. Furthermore, because of sensitivity limitations, the statistics of radio emission from different types of optical SNe is very poor, with only 6 examples of SNe Ib/c and no examples of SNe Ia ever detected. Even the generally radio-brighter Type II SNe have only a dozen detections and fewer than half of that number have well measured, multi-frequency radio light curves.

3.7.2 New Observations Possible with the SKA Radio Emission from Type Ia SNe

With the SKA our RSNe studies would enter a new era. We would be able to monitor RSNe at a practical threshold up to distances ten times further than is currently possible. As a result, statistics for both the Type II and Type Ib/c RSNe would substantially increase. At the very least, we would establish detections for a far larger sample of objects. An aspect where the SKA would greatly advance RSNe studies and our general understanding of SNe is in the possible detection of radio emission from Type Ia SNe. These are the luminous objects currently serving as powerful cosmological probes out to z 1 and providing interesting constraints on Ω and Ω (Riess et al. 1998; ∼ M Λ Perlmutter et al. 1998). Yet, we still have no idea what stars are giving rise to Type Ia SNe, although we suspect theoretically that they involve the deﬂagration or detonation of a white dwarf in a mass-transfer binary system. Some scenarios for Type Ia SNe progenitor systems (see, e.g., Branch et al. 1995) would result in a CSM around the SN, which could produce faint radio emission, currently below the sensitivity limit for the VLA. The level of the SN shock/CSM interaction for Type Ia SNe, and its implication on the nature of the progenitor system, awaits a more sensitive radio array for these SNe to reach out to at least Virgo Cluster distances.

RSN Distance Determinations

Weiler et al. (1998) have presented evidence that the radio emission from SNe may have quantiﬁable properties which allow for distance determinations. Type Ib and Ic RSNe, based on a statistically very small sample of only four objects, may be approximate radio “standard candles,” with 6 cm peak luminosities of 19.9 1026 and 6.5 1026 erg s−1 Hz−1, respectively. Type II RSNe, based on a∼ small sample× of twelve∼ objects,× appear to obey a relation L 5.5 1023 (t t )1.4 6 cm peak ≃ × 6 cm peak − 0 erg s−1 Hz−1 (with time in days) (Fig. 3.25). Thus, measurement of the radio turn-on 3.7. SUPERNOVAE 131

Figure 3.25: Peak 6 cm luminosity, L6 cm peak of RSNe vs. time, in days, from explosion to peak 6 cm flux density (t6 cm peak t0). Type II SNe are plotted as filled triangles. The dashed line is the unweighted,− best fit to the 12 available Type II RSNe. The Type II SNe show a large range in times to 6 cm peak (t t ) 6 cm peak − 0 and in peak 6 cm luminosity L6 cm peak, but appear to obey the relation given in the text. Where no error or only a stub of a line is shown, the error in that direction is indeterminate.

time (t6 cm peak t0) and peak ﬂux density S6 cm peak can yield a luminosity estimate and therefore a− distance. The reality of the “standard radio candle” hypothesis for Type Ib and Type Ic RSNe may be tested simply through the study of more objects, but since Type Ib/c SNe occur relatively infrequently nearby, higher sensitivity is needed to study more objects at larger distances. While the same is true for Type II SNe, some examples of the class are bright enough that 1 per year can presently be detected in the radio to slowly increase the available statistics.∼ For the fainter Type II SNe, however, there exists a large gap in our knowledge between the very faint, somewhat oddball SN 1987A ( 3 x 1023 erg s−1 Hz−1 at 6 cm peak; which could only be detected in the radio because∼ it was extremely nearby in the LMC), and the faintest of the normal RSNe, such as SN 1980K, which can be observed to the VLA sensitivity and are more 132 CHAPTER 3. FORMATION AND EVOLUTION OF STARS

Figure 3.26: A plot of peak C band (6 cm in the observer’s frame) ﬂux density for various well-observed RSNe, if they were moved to higher redshift, assuming a −1 −1 Friedmann cosmology (ΩΛ = 0) with q0 = 0.50 and H0 = 75 km s Mpc . The prototypical RSNe are SN 1980K (open circles), SN 1979C (crosses), and SN 1988Z (asterisk). The 5σ limit of the SKA is shown by the horizontal red bar. The SKA will detect supernovae in galaxies at redshifts 1 to 4.

than two orders-of-magnitude brighter at 6 cm peak.

Additionally, the SKA holds the possibilities for detection of the very luminous RSNe IIn at large distances. Fig. 3.26 illustrates that, at a sensitivity level of 150 nanoJy (10σ in 8 hours), one can detect the brightest of RSNe, such as the Type IIn SN 1988Z and SN 1986J, at the z = 3. Even more normal Type II RSNe, such as SNe 1979C and 1980K can be studied at cosmologically interesting distances (z 1). ∼ If we can extend our horizons to observe SNe up to a redshift of 1 we will ﬁll in the gaps in our knowledge and improve the statistics for RSNe of all∼ types, such that RSNe may eventually provide a powerful and independent technique for investigating the long-standing problem of distance estimates in astrophysics. 3.7. SUPERNOVAE 133

The SN-SNR Connection

Old SNe, such as SN 1968D in NGC 6946, SN 1970G in M101, and SN 1923A in M83, provide a connection to young supernova remnants (SNRs). Currently, a large gap in time exists between the oldest RSNe, such as these, and the youngest radio SNRs such as Cas A (SN 1680). (See Fig. 3.27.) Bridging this gap and understanding the connection between∼ SNe into SNRs is vital for our understanding of the evolution of SNe, their interaction with the CSM, and their energy and chemical input into the ISM – with the resulting inﬂuence on star formation and galaxy evolution. The SKA would potentially allow detection of decades-old SNe which may still be radio emitters, but are currently well below the sensitivity limit of the VLA.

3.7.3 Summary and Conclusions

The present limit to our studies of RSNe are intimately tied to the present limitations of the VLA, such that: 1) a realistic study limit is 1 mJy peak ﬂux density; 2) we can only detect normal SNe to Virgo Cluster distances∼ ( 20 Mpc); 3) we can only ∼ ∼ study very luminous RSNe to 100 Mpc; 4) there is signiﬁcant delay from observing to mapping; and, 5) no realistic∼ SN search modes are possible. The current radio detection limit is to optical magnitude m 12 for normal v ∼ Type II SNe. 1300 SNe are known, with only 25 radio detections. 150 SNe are discovered each∼ year, with only 1–2 radio detections.∼ The SKA could∼ extend ∼ RSN detections to mv 19, such that 50 radio detections per year would become possible. The SKA could∼ provide better∼ SN statistics not limited by absorption and dust, and, as a result, could discover “hidden” SNe. The SKA could therefore provide better galaxy SN rates, which would provide improved chemical and dynamical galaxy evolution modeling. Radio data are vital for understanding the nature of SNe progenitor stars and stellar system, by probing the pre-SN mass loss in the late stages of the progenitor’s stellar evolution. As a result, radio data place important constraints on the SNe progenitor properties and masses. Improved radio data could bridge the SN-SNR time gap. With the SKA, normal SNe would be radio detectable to z 1. Bright SNe would ∼ be radio detectable to z 3. Radio distance estimates would then be available from ∼ the radio peak luminosity vs. turn-on time relation. H0 determinations could be made independent of optical limitations. Estimation of other cosmological parameters, such as q0 and Ω, might be possible. The current VLA is severely sensitivity limited for SN studies. The lack of on-line mapping at the VLA precludes RSN searches. The SKA would improve SN environment/progenitor studies, would improve SN statistics, would lead to improvements in our understanding of galactic chemical and dynamical evolution, and would provide independent distance and cosmological parameter estimates. 134 CHAPTER 3. FORMATION AND EVOLUTION OF STARS

SN 1986J

SN 1994I SN 1978K SN 1979C

SN 1980K SN 1957D Cas A SN 1950B (SN~1670) SN 1970G SN 1968D Crab SN 1983N SN 1923A (SN 1054)

Kepler SN 1993J (SN 1604) 3C58 (SN 1181) Tycho (SN 1572)

(SN 1006)

0.01 0.1 1 10 100 1000 Age [yr]

Figure 3.27: The radio luminosity vs. age for a number of young radio supernovae (curves), older “historical” supernovae which have been detected (squares), and the youngest known galactic SNRs (triangles). 3.8. THE RADIO AFTER-GLOWS OF GAMMA-RAY BURSTS 135

3.8 The Radio After-Glows of Gamma-ray Bursts

The unique status of the supernova explosion has now been challenged by a more energetic, bizarre, and still rather mysterious event, which has been observed for a quarter of a century, though its true distance only became apparent in 1997: the gamma-ray burst (Fishman and Meegan 1995). Discovered (accidentally) by military satellites looking for high-energy emission from nuclear tests (and, like radio emission from the sun 30 years before, first declared top secret), gamma-ray bursts (GRBs) are the brightest objects in the sky outside of the solar system. With a peak magnitude- equivalent at MeV to GeV energies brighter than v = 6m (for a few seconds), they would outshine the brightest stars if we could “see” at− such high energies. Whatever their exact nature, the short timescales involved (with some features lasting < 1 ms) betray an origin in a tiny region of space, and most of the models proposed involve neutron stars, in particular the merger of a neutron star pair. In 1997, a concerted effort by optical, gamma- and X-ray astronomers resulted in the first detection of a transient optical (van Paradijs et al. 1997) and X-ray (Costa et al. 1997) event associated with the burst GRB 970228. The key to success was rapid follow-up of the gamma-ray burst with a high-definition X-ray telescope (the BeppoSAX satellite), and the equally quick determination of a position to within a few arc minutes for immediate optical imaging. Optical transient events have now been observed from a handful of GRBs, all have been faint (v > 20), they may lie in distant galaxies, and in one case, GRB 970508, absorption lines indicate a redshift of z > 0.83 (Metzger et al. 1997). A weak radio transient was also observed in GRB 970508 (Frail et al. 1997), with peak flux densities of about 1 mJy. Unlike the X-ray and optical transients, which decayed on timescales of hours or days, the radio event lasted for weeks, turning on later at longer wavelengths. Besides its general behavior, rapid intensity variations attributed to scintillation set interesting limits to the GRB size. It’s also possible that many GRB events will be detectable in the radio even when the optical emission is highly obscured by dust as seems to be indicated in the case of GRB 980703 shown in Fig. 3.28. Radio observations provide critical and unique elements of the physics of gamma- ray bursts.

1. Rapid Subarcsecond Localization. At the present time the position of the gamma-ray and X-ray transient is only localized to arcminute accuracy. The role of providing sub-arcsecond positions of the GRBs relative to their host galaxies is the role of ground-based radio and optical instruments. Drawing upon historical analogy from the study of supernovae, we note that the location of GRBs within their hosts will provide important clues to one of the key unanswered questions: the identity of the GRB progenitors. For the moment, the detectionand morphology of the host galaxies and determining their star formation rates is exclusively an optical endeavour. With improvements in sensitivity the radio synchrotron from the host galaxy can be detected, allowing unique studies to be made of optically obscurred galaxies (see below). 136 CHAPTER 3. FORMATION AND EVOLUTION OF STARS

Figure 3.28: The broad-band radio/IR/optical/X-ray spectrum of GRB980703 ﬁve days after the burst. The synchrotron spectrum is self-absorbed at low frequency but rises as ν0 to ν1/3 in the optically thin regime, and falls as a power law at higher frequencies. See Bloom et al. (1998a) for more details. Figure 1b. (right) The radio light curve for GRB980703 at 8.46 GHz. Scintillation is strong at 5 GHz but is negligible at 8.46 GHz. The chief advantage of the radio band over the X-ray and optical bands is that one tracks the full evolution of the ﬁreball, from an optically thick source becoming optically thin, rising to a maximum and then undergoing the characteristic power law (t−1) decay. 3.8. THE RADIO AFTER-GLOWS OF GAMMA-RAY BURSTS 137

2. Dusty Hosts. The radio afterglow can penetrate the dust within the host galaxy, and thus GRBs can be seen both from unobscured and dusty galaxies. There is mounting evidence that at least one third of all GRBs occur in dusty environments (Ramaprakash et al. 1998, Reichart et al. 1998). For example, GRB980329, with an R K color of 5, might have escaped detection were it not for the initial VLA localization.− These findings are especially relevant for distinguishing GRB progenitor models since GRBs naturally arise in dusty regions in the “microquasar” model (Paczyński 1998) but not in the neutron star coalescence model (e.g. Bloom et al. 1998b). Furthermore, if GRBs trace the star formation history of the early Universe then the radio (and X-ray) detection rates relative to the optical detections offer insight into one of the key problems in observational cosmology - the relative amounts of unobscured and obscured star formation at large redshifts.

3. The Physics of Fireballs. The fireball model has been shown to be correct to first order but more detailed multi-wavelength studies are needed. By combining the radio data with that from optical, infrared and X-ray instruments we derive the instantaneous spectrum of the afterglow (Figure 3.28). Knowing the spectral shape it becomes possible to infer the total energetics, the ambient density and to probe the physics of relativistic shocks, addressing issues such as the degree of equipartition between the magnetic field and particle energy density and the efficiency of particle acceleration.

4. Expansion and Geometry of the Fireball. Radio observations provide unique diagnostics on the size and geometry of the ﬁreballs. This is possible thanks to the fact that the boundary between strong and weak scattering in the ISM is around 3 µas at centimeter wavelengths, and this is comparable to the angular size of young ﬁreballs (Narayan 1992). At later times when the scintillation is quenched (due to source expansion) we see the underlying light curve (and spectrum) whose behavior can be used to track the transition from Γ >> 1 to Γ 1 (Waxman, Kulkarni & Frail 1998). ∼ 5. New Phenomenology. It is tempting to assume that the GRB phenomenon is homogeneous but historically such simplistic assumptions have been shown to be erroneous. This is already apparent with the (Galactic) soft gamma- ray repeaters and at least one supernova SN1998bw (z =0.0085), which likely produced a burst of gamma-rays GRB 980425 (Galama et al. 1998, Kulkarni et al. 1998). Synchrotron emission at radio wavelengths will continue to be an important means to study these high energy transients discovered through their X-ray and gamma-ray emissions.

The eﬀect which the superior sensitivity of the SKA would have on observations of objects like GRB980703 is easy to imagine: the emission would be detected at an earlier stage, the “turn on” at several frequencies could be studied in greater 138 CHAPTER 3. FORMATION AND EVOLUTION OF STARS

detail, variations could be measured more accurately to see how well they conform to scintillation behavior, and it would be possible to detect such an event out to a redshift of 4 or more. In addition, of course, radio emission from many more GRBs would be detected, but in view of the limited statistics at the moment (one detection out of perhaps five well-observed events), it is difficult to estimate their numbers. We can, however, get some idea by considering the likely impact of the SKA upon radio observations of supernovae (SNe). With the SKA sensitivity, we will be able to extend such work to many tens of GRBs, and for any plausible luminosity function there should be hundreds more radio detections, since current detections must represent only the tip of the iceberg. Detections can be expected to moderately high redshifts (z = 4 to 5), and if GRBs occur in gas-rich galaxies, it should be possible to detect H I in absorption. This would provide radio limits to the source distance, as well as another probe of conditions along the line of sight to the GRB. This emission results, like the optical afterglow, from the interaction of burst-induced shocks with the ambient medium, and is thus peripheral to the gamma-ray event itself. There are, however, predictions, based upon rather general arguments, of prompt radio emission from the very burst. The detection of such emission, which is expected to be many times weaker than that observed from the afterglow, would contribute directly to our understanding of the burst process itself, and will only be possible with radio telescopes orders of magnitude more sensitive than what we now have. It has taken us almost a quarter of a century to come to grips with the gamma- ray burst problem. Now that we have pinned down the essentials required for their detection in other wavebands, and we have the tools for obtaining accurate positions quickly, further breakthroughs are likely to come sooner rather than later. Nonethe- less, it is clear that except for the intense gamma-radiation itself, observations of the associated phenomena require radio (and optical) telescopes of the greatest sensitivity and agility. Our understanding of these mysterious objects by the time a SKA sees first light will no doubt be better than it is now, and we may have even singled out their progenitors. Nonetheless, more sensitive radio observations will enable us to probe their environment, observe those enshrouded in dust, determine their distance and size, and detect the very low energy tail of the burst itself.

3.9 Pulsars

The discovery of pulsars three decades ago sparked the beginning of investigations and breakthroughs in an extraordinarily wide ﬁeld of physics and astronomy. Mea- surements of the time of arrival of the pulses and pulse characteristics themselves have allowed a deeper understanding of topics ranging from condensed matter with densities of about 1014 g-cm−3 and the extreme physical conditions in the environment of neutron stars, to the evolution of binary systems and mass transfer within them, supernova remnants, and the properties and characteristics of the interstellar medium. Observations of pulsars proved for the ﬁrst time the existence of planets 3.9. PULSARS 139

outside the solar system (Wolszczan and Frail 1992) and have demonstrated the existence of gravitational waves (Hulse and Taylor, 1975). Pulsars can be used, through a combination of accurate pulse timing and VLBI, to tie the earth-orbit oriented solar system reference frame to the quasi-inertial earth-spin oriented extragalactic reference frame (Bartel et al. 1996) and to allow, in the future, more accurate navigation of interplanetary spacecraft. As cosmological clocks they can rival in their long-term stability cesium clocks on earth. Finally pulsars are used in tests of general relativity and in searches for the cosmological background of gravity waves predicted as a remnant from the early universe. Nobel prizes have twice been awarded to pulsar researchers, to A. Hewish in 1974, and to J. Taylor and R. Hulse in 1993. Pulsar research will certainly remain one of the most exciting and productive scientiﬁc areas for years to come. The SKA has the sensitivity to revolutionize the ﬁeld by discovering many thousands of pulsars and making some of the most precise measurements possible in physics to constrain fundamental theories of gravitation, cosmology and nuclear matter.

3.9.1 Pulsar Searches with the SKA

From the density distributions of pulsars in luminosity, L, distance from the galactic plane, and radial distance from the galactic center, it is estimated that the total number of active pulsars is (3.5 1.5) 105 for L> 0.3 mJy-kpc2 of which the radio beams of about 20% sweep past± the earth.× That gives a total of about 70,000 observable pulsars in the Galaxy with L > 0.3 mJy-kpc2 which represents the luminosity limit of the most sensitive surveys with the (sky-limited) Arecibo telescope. The total number of pulsars detected so far, with L > 0.3 mJy-kpc2 is about 1000. Of these only about 30 are millisecond pulsars with periods from 1.6 to 20 milliseconds and a median value of a few milliseconds. Note that, due to sensitivity limitations, only about 4% of the detected pulsars have luminosities between 0.3 and 3 mJy-kpc2, although they represent the vast majority, 80% to 90%, of the galactic distribution. Most pulsars have been discovered in large scale surveys performed by single dishes such as the Arecibo 305m, Parkes 64m and Jodrell Bank 76m telescopes. There are two exciting ways that SKA could be used to discover pulsars in large numbers. The ﬁrst is in targeted searches of locations known to be rich in pulsars, such as globular clusters, the Magellanic Clouds and the Andromeda Galaxy. In radio images of these regions with good spectral and spatial resolution, pulsars will manifest themselves as steep-spectrum unresolved point sources. By using the SKA in tied array mode where up to one hundred pencil beams can be formed within the primary beam searches of the best pulsar candidates could be made at the full sensitivity of the instrument. We conservatively estimate that the SKA could ﬁnd of order 1000 pulsars in the SMC/LMC system in a survey lasting just one month. Globular clusters would be another rich source of pulsars, with their dense cores proven breeding grounds for the formation of relativistic binaries. The most interesting pulsars in these dense regions could be timed using the multiple beams available with the SKA. The potential to 140 CHAPTER 3. FORMATION AND EVOLUTION OF STARS

uncover black hole neutron star and other relativistic binaries is enormous due to the sheer numbers of new pulsars that would be found. The other less exotic mode for pulsar searching with the SKA would be to form an incoherent addition of the signals from each antenna element and use traditional search techniques. Although the effective area of the instrument only scales as the square root of the number of elements, the instrument would still be a very effective search instrument. This is because the area surveyed in each pointing is that of the primary beam of a single element, and the loss of sensitivity over a single dish is exactly made up by this increased search area, making the instrument as good (per square metre) for searching as a single dish with a square km of collecting area! In this mode SKA could survey 80% of the entire sky at 70cm in 6 months with a sensitivity 20 times that of the Parkes 70cm survey of the Southern sky which detected 298 pulsars, 19 of them with millisecond periods. Given that pulsars are disk population, we might expect a survey with SKA to detect 6000 pulsars in the Southern sky alone. In practice, interstellar dispersion and scattering∼ would probably prevent us from following the expected Log-N, Log-S relationship, but it is obvious that SKA would find pulsars by the thousands, and several hundred millisecond pulsars. The potential for the discovery of the first black hole and dual-line binary pulsars would be great. Such systems would be ideal gravitational laboratories. The number of relativistic neutron-star binaries would probably increase by at least an order of magnitude and better constrain the poorly-known coalescence rate of neutron star binaries. The best-studied planetary system outside the solar system is that orbiting the millisecond pulsar PSR B1257+12 (Wolzsczan 1994), and it is possible that several more instances of pulsar planetary systems would be found.

3.9.2 Pulsar Timing with the SKA

The tied array would have a much higher sensitivity than the incoherent search mode, and would be able to time any pulsars found in the all-sky survey in a small fraction of the time it took to discover them. Although the ﬁnal conﬁguration of the SKA is yet to be determined, with 625 45m dishes, the tied array could achieve the survey sensitivity in 1/625th of the time. It would be thus be possible to time thousands of pulsars in a day. The tremendous increase in the sensitivity of the tied array means that sub- microsecond timing residuals could be obtained on several hundred millisecond pulsars. This would make possible the formation of the ultimate pulsar timing array (Foster and Backer 1990). It has already been demonstrated that millisecond pulsars rival the best atomic clocks for stability on timescales of a year or so. Several hundred pulsars would establish a highly precise pulsar time standard. The Earth’s ephemeris is constantly being revised because of uncertainties in the position of the solar system barycenter. The pulsar timing array would be extremely sensitive to the masses and motion of the planets. Long-period gravitational waves are detectable in pulse timing experiments (Rajagopal and Romani 1995). The millisecond pulsar 3.9. PULSARS 141

timing array would place the strongest limit yet on the background of gravitational waves and make a sensitive detector for sources of long-period gravitational waves, such as ultra-massive black hole binaries. This new population of pulsars would contain at least twenty double neutron star systems suitable for exhaustive tests of general relativity. In particular a dual- line binary pulsar would allow the strictest tests yet of strong-ﬁeld gravity. All of the known double neutron-star binaries have only one visible pulsar. By detecting a dual- line pulsar, both the mass ratio and the sum of the masses could be determined to extraordinarily high accuracy, probably greater than one part in 106. This would, in conjunction with the three Keplerian parameters, completely specify the orientation and masses of the binary system. Therefore measurements of all of the remaining post- Keplerian parameters would allow tests of relativistic theories of gravity. A black hole neutron star system would also be of special interest. The relativistic eﬀects would be much larger than in a double neutron star binary. Millisecond pulsars with edge- on orbits could have their Shapiro delays measured, and this in turn would provide information on neutron star masses. The increase in computer power by the time of the survey would allow the data to be searched at a frequency resolution much greater than is currently possible, making sub-millisecond pulsars detectable. The discovery of such a pulsar would strongly constrain the equation of state of condensed nuclear matter (Phinney and Kulkarni 1994).

3.9.3 Radio Pulsar Timing and General Relativity Radio pulsars are highly-magnetized neutron stars which are formed during the core- collapse of massive stars in an event known as a supernova explosion. They represent the densest form of matter possible. Although weighing around 500,000 Earth masses, they are only 20-30 km in diameter making them ideal “point-masses” in physical experiments. They are formed approximately once every 250 years in our Galaxy, and the fastest known object rotates once every 1.55 milliseconds. Radio astronomers routinely make observations of these systems that enable them to determine the arrival time of the pulse to an accuracy of a microsecond or better. By making observations at several epochs astronomers have been able to establish that millisecond pulsars have rotational stability that rivals the Earth’s best atomic clocks (Kaspi, Taylor and Ryba 1994). In some special cases, pulsars orbit unseen companions, allowing tests of relativistic gravity. Using the powerful pulse timing method, astronomers have shown that the orbit of the binary pulsar PSR 1913+16 is shrinking at a rate of 3.1 mm per orbit, precisely the amount predicted by general relativity. At this rate the system will coalesce in 300 Million years in a spectacular death-spiral that will emit an enormous burst of gravitational radiation which may be detected by gravitational wave observatories. The stability of the pulsar rotation has been used to place important limits on the density of gravitational waves in the Universe and rule out some cosmological models 142 CHAPTER 3. FORMATION AND EVOLUTION OF STARS

(Kaspi, Taylor and Ryba 1994). Remarkably, an ensemble of millisecond pulsars with sub-microsecond timing accuracy could be used as a gravity wave telescope to detect ultra-massive black hole binaries anywhere in the Hubble volume (Rajagopal and Romani 1995). Above we argue that SKA will be capable of discovering thousands of pulsars and performing pulse timing observations with a precision far in excess of that currently possible. This will enable rigorous tests of relativistic theories of gravity, the establishment of a gravity-wave telescope suitable for detecting long- period gravitational radiation, and the determination of a large number of accurate neutron star masses. Other applications will include reference frame ties, studies of the interstellar medium, improvements in the solar system ephemeris, and advances in the theory of stellar evolution and pulsar emission mechanisms.

3.9.4 Summary The SKA would be capable of discovering many thousands of radio pulsars. Many of these pulsars would be members of binary systems, permitting exhaustive tests of general relativity. The millisecond pulsars found by the instrument would form a pulsar timing array capable of detecting long-period gravitational waves such as those emitted in the early formation of the Universe and by ultra-massive black- hole binaries. The timing array would form a new standard of time and furnish information about the solar system ephemeris. The mass of neutron stars in binaries and the minimum spin of millisecond pulsars are both paths to a mini holy grail - the equation of state of nuclear matter in the center of neutron stars which exceeds that in atomic nuclei even when collided at high energy. Surveys with SKA would provide the best opportunity to discover a sub-millisecond pulsar.

Chapter 4

Solar System Science

4.1 Thermal Emission from Small Solar System Bodies

The solar system oﬀers a wide variety of objects to study, most of which are very poorly understood. Other parts of this chapter focus on planetary atmospheres, the gaseous emissions of comets, radar astronomy and solar physics. This section will focus quite tightly on what might be learned from observations of thermal emission from the smaller, solid bodies of the solar system using the SKA. The SKA is proposed to have an angular resolution of 0.1′′ at 1.4 GHz. Table 4.1 shows how the resolution and sensitivity targets translate into the context of solar system astronomy.

2 GHz 20 GHz Angular resolution (mas) 70 7 Linear resolution (km at ∆ AU) 50∆ 5∆ −1 Sensitivity σF (µJy hr ) 0.056 0.048 −1/2 Thermal sensitivity (single pixel) σT (K hr ) 1.4 1.2 Smallest main-belt object (km) 12 1.2 Smallest Kuiper-belt object (km) 1300 120

Table 4.1: SKA Speciﬁcations for Imaging Small Solar System Bodies

Useful comparisons can be made with the HST, which deﬁnes the current standard for image quality, and with the projected performance of other proposed telescopes, such as the LMA and NGST. The HST currently provides images with roughly 100 mas resolution at visible wavelengths. The NGST will provide images with resolutions between 25 and 50 mas in the near infrared (1 - 2 µm). The LMA is intended to provide resolutions comparable to Hubble. Thus, at 20 GHz the SKA would provide resolutions better by a factor of 3 than the NGST and at least a factor of 10 better than the HST.

144 4.1. THERMAL EMISSION FROM SMALL SOLAR SYSTEM BODIES 145

In the next few subsections we will consider in more detail how the SKA might be used to study of these smallest, and least-understood members of the solar family.

4.1.1 Asteroids

In the asteroid belt the process which collected small protoplanets into the major terrestrial planets shut off prematurely because of strong gravitational perturbations from the early Jupiter. In the inner belt, some of the proto-asteroids melted and differentiated like the major planets. Others warmed enough to melt ice and show signs of aqueous alteration in their minerals. Still others, especially in the outer belt, seem to have remained frozen for their entire existence. Subsequently, the proto- asteroids evolved by collisional fragmentation into the population of asteroids we see today. The asteroid belt thus preserves for us samples from an ensemble of smaller protoplanets which froze in a variety of evolutionary states, broken up for our convenient inspection. We can never hope to obtain comparable samples from the mantle and core of the Earth! Bearing in mind that the SKA promises to provide images at 20 GHz significantly sharper than any competing instruments, the large asteroid Vesta will be used to illustrate the importance of high-quality imaging in the study of asteroids. With a diameter of 520 km, Vesta subtends 0.7 arcsec at a typical opposition distance of 1.3 AU, and is an appealing target for imaging with the best currently available telescopes.. Uniquely amongst the asteroids, Vesta has been imaged with the HST (e.g Zellner et al. 1997), from the ground with adaptive optics (Drummond et al. 1998) and at radio wavelengths with the VLA (Johnston et al. 1989). Prior to these images, several decades of lightcurve analysis had only just been able to establish that the true rotational period was 5.2 hours, not 10.4 hours, and it was widely believed that much of the shape of the lightcurve was due to albedo patches and not simply to an odd shape. Between November 28 and December 1, 1994, the HST was used to image Vesta (Zellner et al. 1997) at a distance of roughly 1.7 AU. With a linear resolution of roughly 80 km, the visible disk was roughly 6.5 resolution elements across. These images, using roughly 8 minutes of exposure time, completely determined the rotational pole of Vesta, showed that it rotated in the prograde sense, and unambiguously demonstrated that the rotation period was 5.2 hours. Compared to the lightcurve analyses, the data gathering and analysis of the HST images was faster by a thousand- fold and was free from most of the ambiguities which plagued the earlier efforts. The SKA at 20 GHz could have made equally valuable images of objects 10 times smaller than Vesta. There are only three main-belt asteroids with diameters as large as Vesta’s, but in the IRAS FP102 database of asteroid albedoes and diameters (Tedesco et al. 1992) there are 655 with estimated diameters larger than 50 km. For those studies which do not require well-resolved images, almost all of the 1884 asteroids in the FP102 database had estimated sizes larger than 10 km, the resolution limit for the SKA at 2 AU. The true number of resolvable asteroids will be larger 146 CHAPTER 4. SOLAR SYSTEM SCIENCE

Figure 4.1: HST Image of Vesta showing a large crater at its south pole Zellner (1997). The SKA will provide a ten-fold improvement in angular resolution, allow studies of the geological history of this and many other asteroids.

by one or two orders of magnitude because the FP102 database becomes seriously incomplete for asteroids smaller than 100 km. The importance of pushing the highest operating frequency of the SKA to at least 20 GHz starts to become clear at this point. If the highest frequency is restricted to 2 GHz, the SKA will only provide a linear resolution of 85 km, no better than that already possible with the HST. Only a dozen of the largest asteroids could possibly be mapped at 2 GHz. The study of radio images of asteroids could be complete within the ﬁrst year of operation, probably without any really surprising discoveries on these well-studied objects. Returning to Vesta, the HST images allow us to begin to study the geological history of this object. A later image (Fig. 4.1) shows a very large crater at the south pole of Vesta, and perhaps two other large craters near its equator. With a ten-fold improvement in the angular resolution of the raw image, we can expect that many smaller craters will be visible, allowing cratering-ages to be estimated for the larger units on the surface. Unlike Vesta, many smaller asteroids are suspected to be “rubble piles”, collections of fragments loosely held together by their own gravity. We might expect such objects to be lumpy, much like the radar images of the asteroid Toutatis in Figure 4.5. With images of hundreds of asteroids available, we should be able to test this model rigorously. Furthermore, the radio data will contain information not available at shorter wave- 4.1. THERMAL EMISSION FROM SMALL SOLAR SYSTEM BODIES 147

0.5

0.45 Vesta 2mm L Jy H 0.4

0.35 Flux Density 0.3

0.2 0.4 0.6 0.8 1 1.2 Rotational Phase

Figure 4.2: Radio light curves of asteroids provide information not available at shorter wavelengths. The light curve of Vesta at λ2mm shown above has twice the amplitude of the visible light curve (Redman et al. 1998b) and has changed from a single-peaked shape in less than 10 years. lengths. As reported at the Annual General Meeting of the Canadian Astronomical Society in May 1998, (CASCA’98) the rotational lightcurve of Vesta at 2 mm wavelength, measured with the SCUBA detector on the James Clerk Maxwell Telescope on October 14, 1997, had an amplitude twice as large as the visible lightcurve (Fig. 4.2 Redman et al., 1998b). At the time of writing, it is not clear why the amplitude of this lightcurve is so large, why it was single-peaked when our previous observations of Vesta in September 1989 showed a doubly peaked curve more in accordance with our theoretical expectations, nor how many other asteroids will have similarly informative radio lightcurves. On a longer timescale, the SKA might be used to prospect for mineral deposits on asteroids to be exploited by space-based mining ventures in the next century. Judging from the composition of iron and stony-iron meteorites, the M-type asteroids consist primarily of a Fe-Ni alloy with traces up to several hundred ppm of precious metals (Kargel, 1994). Ignoring the iron and nickel, the precious metal content of even a small (1 km) M-type asteroid should be worth roughly $5 1012 at current prices. × Our observations of the large M-type asteroids Psyche and Kleopatra with the JCMT (Redman et al., 1998a) show that the emissivity of these metallic bodies is roughly half that of stony or carbonaceous asteroids and that the emissivity drops rapidly with increasing wavelength. The emissivity at 20 GHz might be as little as a few percent. The tremendous sensitivity of the SKA should still make these objects easy to observe, with the ore bodies cleanly distinguished as dark patches amongst the brighter silicate rocks. For the nonmetallic asteroids, thermal models suggest that 20 GHz is a nearly 148 CHAPTER 4. SOLAR SYSTEM SCIENCE

optimal observing frequency. Following the discussion in Redman et al., 1998a the deviations of an asteroid’s spectral energy distribution (SED) from a Planck spectrum are dominated by different effects in different wavelength bands. At short infrared wavelengths the shape of the SED is determined by the temperature and area of the warm region around subsolar point. In the far infrared and submillimeter the shape of the SED is strongly affected by the thermal and dielectric properties of the regolith, a layer of shattered rock which is believed to cover the surfaces of most large asteroids. At centimeter wavelengths, starting around 20 GHz, the radiation emerges from layers in the regolith which do not participate in the diurnal heating and cooling cycle as the asteroid rotates. The temperature of this material is just the time average of the surface temperature, an easy quantity to model and interpret. At lower frequencies we get exactly the same information but with lower signal to noise. At higher frequencies we start to see the effects of the diurnal temperature wave which propagates down into the regolith, complicating the analysis of the data. Indeed, it is difficult to interpret the data at higher frequencies without having the centimeter data for comparison. In this sense, 20 GHz is exactly the correct frequency to observe an asteroid, providing the most easily modeled information at the highest possible signal to noise. We note that the NGST, LMA and SKA will each be able to observe asteroids in one of the key bands mentioned above. The NGST will be able to measure the surface temperatures in the infrared, the LMA will probe the structure of the regolith and the SKA will provide the mean asteroidal temperature from which thermal models may be constructed. These will be three complementary machines, and in an ideal world we would want access to all three. Complementing the studies of thermal emission from asteroids, radar observations offer valuable insights into the surface properties of asteroids. For example, radar might be used to prospect for metallic asteroids, since Ostro et al. (1985, 1991) have shown that they have much higher radar reflectivities at centimeter wavelengths than non-metallic asteroids. Also, the inventory of hazardous Earth-crossing asteroids is biased towards asteroids which reside primarily outside the Earth’s orbit so that they can be seen at night. Radar can easily search for nearby objects in the day-time sky, eliminating this bias. The large collecting area and multiple-beam capability proposed for the SKA may make it a valuable tool for radar studies throughout the inner solar system.

4.1.2 Planetary Satellites

The satellites of the giant planets provide a menagerie of bizarre worlds whose surfaces are unlike anything seen elsewhere in the solar system. Spacecraft have visited these places and have provided detailed maps of many of their surfaces, but have left many questions unanswered, some of which can be best addressed at radio wavelengths. The variety and complexity of each of these worlds precludes a detailed discussion here, but several interesting lines of research are worth mentioning. 4.1. THERMAL EMISSION FROM SMALL SOLAR SYSTEM BODIES 149

Io is so strongly heated by tidal interactions with Jupiter, Europa and Ganymede that it exhibits continuous and vigorous volcanic activity, with half a dozen large vents active at any given time. Although the striking colours on Io’s surface are believed to be due to sulphurous compounds the temperatures in some of the vents are so high that silicate lavas must also be involved. The vents show up brilliantly in the near infrared and are easy targets for even ground-based infrared telescopes using adaptive optics. However, to study the resulting lava flows, which run for hundreds of kilometers across the surface, will be more difficult since they are cool, faint features immediately adjacent to the bright vents. For this study the SKA will be nearly ideal. The linear resolution will be roughly 20 km at the distance of Jupiter, and in the radio the flux density will vary linearly with temperature, unlike the exponential variation in the infrared. The other large Jovian and Saturnian satellites have icy surfaces. The emissivities of these surfaces are quite low, probably because scattering and refraction at fractures and dense lumps within the ice prevents photons from escaping the surface. Many of these objects have been tectonically active and exhibit very different colours and large-scale textures in the different terranes comprising their surfaces. It is plausible to guess that in the different terranes the emissivity of the ice might also be different. Images at 20 GHz made with the SKA would allow us to study the interplay of ice tectonics, water volcanism and impact cratering under a variety of conditions on the Jovian and Saturnian satellites. Probably, the most mysterious world in the solar system is Titan, the only satellite with an extensive atmosphere. From images at infrared wavelengths chosen to penetrate the perpetual cloud cover, it is known that the surface of Titan is mostly dark-coloured, probably covered with an organic sludge created by photochemical reactions high in Titan’s atmosphere and precipitated onto the surface. There are, however, large light-coloured patches which may be icy highlands, swept clean of the organic material by some form of organic “rain”. It is a plausible guess that the emissivity of the light regions may be different from the dark regions. Since Titan’s atmosphere is transparent at centimeter wavelengths, it may be possible to map Ti- tan’s surface using the SKA, with a linear resolution of 40 km compared to Titan’s diameter of 5140 km. These maps may be combined with the radar maps which will be made by the Casini mission to better understand the chemistry and physical state of the materials on Titan’s surface.

4.1.3 Kuiper Belt Objects

Gas giant planets do not seem to have formed beyond Neptune. Instead, the original protoplanets continue to orbit in a vast, extended disk known as the Kuiper Belt. The largest known Kuiper Belt Object (KBO) is the planet Pluto. A few KBOs have been diverted inside the orbit of Neptune. Some of these objects occupy temporary orbits between Jupiter and Neptune and are known as Centaurs, at least one of which, Chiron, shows cometary behaviour. Others, whose orbits carry them into the inner 150 CHAPTER 4. SOLAR SYSTEM SCIENCE

Figure 4.3: Kuiper Belt like Phenomena Around Nearby Stars. The panel at left shows HST Images of Proplids around Newly Formings Stars in Orion (McCaugh- rean and O’Dell 1996). The Proplids are comparable in size to the Kuiper Belt of our Solar System. Similarly-sized dust rings have been detected using the SCUBA submillimeter-wavelength camera on the JCMT around several nearby main sequence stars, such as ǫ Eri (Greaves et al. 1998) (right). SKA observations of the constituents of the Solar System Kuiper Belt will complement exciting new observations in star formation and extra-solar planetary systems. solar system, constitute the population of short-period comets. The Kuiper Belt is comparable in size to the proplids which have been observed around newly forming stars in the Orion Nebula (Fig. 4.3 (left), McCaughrean and O’Dell, 1996) with the HST. At the least, the proplids are Kuiper Belts in formation, and may contain newly-forming planetary systems closer to their central stars. Similarly-sized dust rings have been detected using the SCUBA submillimeter- wavelength camera on the JCMT around several nearby main sequence stars, such as ǫ Eri (Fig. 4.3 (right), Greaves et al, 1998). Since the lifetime of these dust particles due to Pointing-Robertson drag is short compared with the lifetime of the central star, there must be a reservoir of larger objects just like the Sun’s Kuiper belt around each of these stars. Curiously, almost all of the dust rings seen with SCUBA are clumpy rather than axisymmetric around the central star. It is not clear whether the clumps are transient features or represent structures such as density waves or swarms of particles around large objects like planets embedded in the belts. Since it spreads so widely across the sky, we cannot see the emission from dust in the Sun’s Kuiper Belt, but we can see the larger objects which supply the dust. 4.1. THERMAL EMISSION FROM SMALL SOLAR SYSTEM BODIES 151

Observations of the solar Kuiper Belt therefore nicely complement and tie together some of the most exciting new observations in the fields of star formation and extrasolar planetary systems. At the time of writing the Minor Planet Center lists 68 KBOs, not including Pluto and Charon. Luu and Jewitt (1998) estimate there are 160,000 objects in the Kuiper Belt closer than 50 AU with diameters in excess of 100 km, hundreds of times more than similar sized objects in the main asteroid belt. Since the Kuiper Belt extends to about 100 AU, there may be comparably many objects again waiting to be discovered in the outer parts of the belt. In spite of their large numbers, KBOs are faint, difficult sources, and it will be especially troublesome at visible and infrared wavelengths to find members of the outer Kuiper belt, beyond 50 AU. The brightness of reflected sunlight drops off with distance d as d−4, making objects in the outer belt 5 magnitudes fainter than comparable inner belt objects. Similarly, the low temperature of the KBOs, dropping from around 40 K in the inner belt to only 20 K in the outer belt, means that the thermal emission does not start to dominate the reflected sunlight until far-infrared wavelengths. By contrast, the SKA operating at 20 GHz will be almost ideally suited to study KBOs. The linear scale which the SKA will resolve at a distance of 40 AU is 200 km. After an 8 hour integration it should be possible to achieve a 5 σ detection on a 120 km object, as shown in Table 4.1. Even at a distance of 100 AU, he smallest detectable object would be 350 km in diameter, and there may be many of these objects if their size distribution is similar to that which we have measured on the inner edge of the belt. If self-gravitating clusters of KBOs exist, capable of generating clumps of dust emission such as are seen around other stars, the SKA may be the tool required to find them. Very little is yet known about the physical properties of KBOs. Luu and Jewitt (1996) have shown that they have a wide range of colours, possibly indicating collisional resurfacing. The relatively slow orbital speeds in the outer solar system suggest that collisions in general will be slower and less violent than in the asteroid belt. This should be reflected in the shapes and surface features of the KBOs, but it is not yet known whether KBOs should be modeled as rubble piles, as solid, primordial bodies, or as some other kind of structure entirely, perhaps reflecting the original, gentle growth of the protoplanets. At the moment, nothing is known about their shapes, rotation periods, masses or regolithic structures. By imaging the larger objects it will be possible to address many of these issues with the SKA. Their shapes and rotational properties will be directly observable. With direct measurements of the thermal flux at radio wavelengths it will become possible to build reliable thermal models of the surfaces of these objects, giving insight into their densities and compositions. Curiously, with the strawman specifications for the SKA it will not be necessary to revive the old lightcurve-analysis techniques which have been used at shorter wavelengths to study asteroids. Asteroids are seen as bright point sources, but the surface brightness of KBOs will be so low that they will necessarily be extended sources. Reliable measurements will only be obtainable 152 CHAPTER 4. SOLAR SYSTEM SCIENCE

for objects large enough to ﬁll at least one pixel. To do this work it is essential that the SKA operate at 20 GHz. No machine less sensitive than this will even be able to detect these small, cold worlds. Restricting the top frequency to 2 GHz would prevent us from studying any but the largest few KBO’s, comparable in size to Pluto and Charon, as shown in Table 4.1. For solar system astronomy, the unique ability of the SKA at 20 GHz to observe KBOs throughout the full extent of the belt may be the most compelling reason to build it.

4.2 Radar Imaging of Near Earth Asteroids

Radar is already the most powerful ground-based technique for post-discovery in- vestigation of NEOs (near-Earth objects). Radar is uniquely capable of resolving NEOs spatially by measuring the distribution of echo power in time delay (range) and Doppler frequency (line-of-sight velocity) with very fine precision in each coordinate. Current capabilities utilizing the Goldstone, 70 m NASA Deep Space Network antenna, and the Arecibo, 305 m telescope, are for best resolutions of order 10 m in range and 0.1 mm/s in velocity. The fractional precision of delay-Doppler positional measurements plus their orthogonality to optical astrometry makes them invaluable for refining orbits and prediction ephemerides. A single radar detection secures the orbit well enough to prevent “loss” of the object, shrinking the object’s instantaneous positional uncertainties by orders of magnitude with respect to an optical only orbit and greatly improving the accuracy of long-term trajectory predictions. Additional advantages come from radar wavelength’s sensitivity to near surface bulk density and roughness. For comets, radar waves can penetrate optically opaque comas to examine the nucleus and can disclose the presence of macroscopic coma particles. One of the more spectacular NEO results of recent years has been a sequence of delay-Doppler images of 4769 Castalia (1989 PB) shown in Fig. 4.2 (Ostro et al. 1990) obtained two weeks after its Aug. 1989 discovery. These reveal it to consist of two kilometer-sized lobes in contact. Least-squares estimation of Castalia’s three- dimensional shape from the radar images supports the hypothesis that Castalia is a contact-binary asteroid formed from the gentle collision of the two lobes and also constrains the object’s surface morphology and pole direction. Another illustrative result is the delay-Doppler imaging of 4179 Toutatis in Dec. 1992 (Ostro et al. 1995). Resolutions as fine as 125 ns (19 m in range) and 8.3 mHz (0.15 mm/s in radial velocity) were obtained, placing thousands of pixels across the asteroid. The images shown in Fig. 4.5 reveal this asteroid to be in a highly unusual, nonprincipal -axis (NPA) spin state with several-day characteristic time scales. This data set provides physical and dynamical information that is unprecedented for an Earth-crossing object. Use of the SKA in a bistatic radar configuration with Arecibo or Goldstone as the transmitter would have a major impact on the study and detection of NEOs. The factor of 30–200 increase in receiver sensitivity with Arecibo and Golstone respectively) would permit objects of smaller diameter to be observed out to substantially 4.2. RADAR IMAGING OF NEAR EARTH ASTEROIDS 153

Figure 4.4: Radar images of asteroid 4769 Castalia (1989PB). This 64 frame Arecibo movie is to be read like a book (left to right in the top row, etc.). The radar lies toward the top of the ﬁgure in the image plane, which probably is about 35◦ from the asteroid’s equator. In each frame, echo power (i.e., the brightness seen by the radar) is plotted vs. time delay (increasing from top to bottom) and Doppler frequency (increasing from left to right). The radar illumination comes from the top of the ﬁgure, so parts of the asteroid facing toward the bottom are not seen in these images. The object, each of whose lobes is about a kilometer in diameter, is seen rotating through about 220◦ during the 2.5 h sequence. These images have been smoothed from the original resolution of 2 µs 1 Hz (0.3 0.17 km). (from Ostro et al. 1990). × × 154 CHAPTER 4. SOLAR SYSTEM SCIENCE

Figure 4.5: Radar images of asteroid 4179 Toutatis from 1992. Goldstone low- resolution images from Dec. 2–18 (top three rows) and Arecibo images from Dec. 14– 19 (bottom row) are plotted with time delay increasing toward the bottom and Doppler frequency increasing toward the left. On the vertical sides, ticks are 2 µs apart. Two horizontal sides have ticks separated by 1 Hz for Goldstone and 0.28 Hz for Arecibo; those intervals correspond to a radial velocity diﬀerence of 18 mm/s. (from Ostro et al. 1995). 4.3. THE ATMOSPHERE AND MAGNETOSPHERE OF JUPITER 155

larger distances. The number of objects accessible to radar study would therefore increase dramatically. In addition, the high angular resolution of SKA (as high as 5 mas) would permit simultaneous spatial imaging of the illuminated surface, greatly simplifying the physical modeling. The linear resolution aﬀorded by such a spatial resolution, 400 m at a distance of 0.1 AU, is such that a detailed three dimensional image of each object could be obtained instantaneously. A single observing session would then also provide extremely high precision in the orbital determination, since the three dimensional space velocity would follow from the proper motion observed in an interval of only a few hours.

4.3 The Atmosphere and Magnetosphere of Jupiter

Radio emission from Jupiter was first detected in 1955 at a frequency of 22.2 MHz. This emission was sporadic in character, and confined to frequencies less than 40 MHz. It is commonly referred to as decametric radiation. In subsequent years the planet’s thermal emission was detected at short centimeter wavelengths, and its synchrotron radiation at wavelengths between 2 cm and a few meters. The latter radiation is emitted by high energy electrons in∼ a Jovian Van Allen belt. At present, this radiation has been observed and imaged regularly at frequencies between 300 and 8000 MHz. Jupiter’s thermal emission becomes dominant at higher frequencies, and separation of the nonthermal component, although clearly present, becomes increasingly more difficult at frequencies above 8 GHz (λ < 3.6 cm). Jovian emission at frequencies below 40 MHz (decametric range) does have a synchrotron radiation component, we believe, but this is overwhelmed by sporadic bursts of radio emission closer to the planet (Jovian DAM emission) which can reach intensities up to 105 Jy. Due to strong background interference, very few, if any, observations have been made at frequencies below 300 MHz (45 <ν < 300), though some are planned with the VLA and WSRT arrays for the near future. With the SKA, however, the planet’s spectrum can be measured from 40 MHz up to 20 GHz, simultaneously; this gives interesting prospects to determine not only its spectrum at one particular time, but also to monitor any time variations in its spectrum. Such time variability has been measured at the higher frequencies (3 – 90 cm), for instance, following the impact of comet Shoemaker-Levy 9. If the spectrum changes regularly this would contain a lot of information on the physical processes going on in Jupiter’s magnetosphere, such as the interaction between electrons and dust (e.g. de Pater & Goertz 1990). Good brightness temperature measurements of the giant planet atmospheres at the longer wavelengths (ν < 1GHz) would provide invaluable information on the deeper layers of a planet’s atmosphere. Such information can be used to provide an accurate measurement of the ammonia abundance at deep levels, and possibly of any stratification within the ammonia abundance. Although it may push atmospheric models, it is not inconceivable that information on the water abundance can be obtained for Saturn and Jupiter. Nothing really is known about this. 156 CHAPTER 4. SOLAR SYSTEM SCIENCE

4.4 Comet Studies

Comets are interesting and diverse objects. At mm-cm wavelengths, however, comets are very weak, and one needs a lot of sensitivity to image molecular lines and detect continuum radio emission. Comet Hale-Bopp was an exception, observed widely at many frequencies. The continuum radio emission was easily detected at submm wavelengths, it was mapped at 3 mm wavelength, and just barely seen at 3.6 cm with the VLA. Continuum observations give information on the dust coma around the coma, such as the particle size distribution and halo mass. The first OH map was obtained by the VLA on comet Halley; only 2 more comets were mapped since comet Halley. Unfortunately, the VLA array configuration during the Hale-Bopp apparition was too spread out to image OH. This gets precisely at the point why it is so hard to image molecular species in comets: one needs a high sensitivity, as well as a sufficient number of short spacings to image the structure, since a cometary coma is many arcminutes on the sky. So to image a comet’s OH, one needs the VLA D configuration, and even then only a fraction of the emission is usually seen. Observing the coma simultaneously in autocorrelation mode, one actually knows how much flux is missed. Images of OH can be used to derive information on the dissociation process of water, the excitation mechanisms of OH (maser emission), which can be tied in with the physics in a comet’s immediate environment (e.g., interaction with the solar wind). Molecules such as water, ammonia and formaldehyde, which all have line transitions at short cm wavelengths, have only been detected for a few comets, and have never been imaged. Again, one needs a high sensitivity, and the right antenna spacings. We believe the SKA array might actually be the first array to detect and image such species.

4.5 Solar Radar

The El Campo radar built by the Lincoln Laboratory detected 38 MHz radar echos from the Sun in the 1960’s. Huge, rapidly-moving “targets” were occasionally observed but this was before the space-borne coronagraph discovery of coronal mass ejections (CMEs), and the physical nature of these targets was a mystery. We now recognize that CMEs were being observed and that a low frequency transmitter coupled with a high angular resolution receiving array would form an extremely eﬀective system to detect and track CMEs. The combination of Doppler and angular measurements would fully determine the velocity vector of the CME and should provide a reliable prediction of those CMEs that will strike the Earth’s magnetosphere. This topic is of great practical importance because the geophysical eﬀects caused by CMEs can cause hundreds of millions of dollars damage to spacecraft, communication, and electrical power systems. Appropriate transmitting facilities now exist in the form of Over-the-Horizon Radars that are no longer required for military purposes. However, appropriate re- 4.6. CORONAL SCATTERING 157 ceiving facilities are needed for such a project.

4.6 Coronal Scattering

In addition to providing detailed observations of solar nonthermal phenomena (eg. ﬂares) at unprecedented angular resolution, scattering measurements towards background sources oﬀer an excellent opportunity to study the structure of turbulence in the solar corona. Previous occultation studies below 100 MHz have been limited to using the Crab as a background source; the SKA instrument would open up many more sources as possible targets and achieve the angular resolution required to properly compliment higher frequency (eg. 330 MHz VLA) studies. Chapter 5

Formation and Evolution of Life

5.1 Detection of Extrasolar Planets

Recent reports of extrasolar planets have revolutionized the field of planetary formation because of the Jupiter-mass planets that appear to exist in close proximity to their stars (e.g Mayor and Queloz, 1995). Such planet-star configurations may be common. Unfortunately, no extrasolar candidate has yet to be confirmed with independent techniques. Most such candidates have been uncovered with the radial velocity technique. However, given the high angular resolution of the SKA such systems are ideal for astrometric detection. As noted in section 3.4.4, the SKA will detect the non-flaring Sun out to 50 pc. ∼ The binary motion of the Sun due to the orbit of Jupiter would be measurable by the SKA out to several 10’s of pc. The ability of the SKA to detect normal radio emission from solar-type stars would therefore allow the detection of Jovian planets by astrometric observations of the parent stars. Such studies would yield a census of Jovian planets in the Solar neighbourhood and statistical constraints on the range of orbital separations and masses.Independent confirmation of extrasolar planets will be a high priority for future Optical-IR arrays. The SKA will complement those efforts using radio selected stars. Jupiter emits intense, coherent bursts of radiation at decameter wavelengths (Carr and Wang 1990). It has often been suggested that Jupiter-like planets orbiting other stars might be detectable by decametric observations of similar emission. At a distance of 10 pc, Jupiter’s bursts would be at the 10 µJy level. A further consequence of the much smaller orbits in the systems detected to date is the high probability of strong interactions of stellar winds with the magnetospheres of the Jupiter-mass planets. The proximity of the planets to their stars should make low-frequency radio emission from the planets much stronger than what is currently observed from Jupiter. A search for steep-spectrum cyclotron masers from extra-solar planets is planned for the 74 MHz VLA. The SKA would increase the detectable volume by a factor of 1000.

158 5.1. DETECTION OF EXTRASOLAR PLANETS 159

Figure 5.1: The angular motion of a one solar-mass star due to the presence of Jupiter as a function of distance. The astrometric accuracies of the SKA at λ21 cm, λ6 cm and λ1.5 cm are shown by the horizontal lines. At λ1.5 cm the presence of Jovian planets would be detectable up to several tens of pc distance. 160 CHAPTER 5. FORMATION AND EVOLUTION OF LIFE

5.2 Pre-Biotic Interstellar Chemistry

The relationship between linear carbon chains such as carbenes (H2Cx), carbynes (HCx), and cyanopolyynes (HCxN) and the polyaromatic hydrocarbons (sheet-like molecules based on benzene rings) is likely to be the key to understanding the origin of complex organics in interstellar clouds. Extensive observations of the abundances of complex organics at various stages of pre-protostellar evolution are required to identify the specific astrochemical processes which determine the composition of the material from which protostars and their accompanying nebulae form. It is possible that current observations are at a significant sensitivity threshold. Laboratory measurements of the production of cyanopolyynes show that the relative abundance flattens when there are more than nine carbons, reaching at asymptotic value only factor of ten less for more than 15-17 carbons (McCarthy et al. 1998). However, because the energy is distributed over many more transitions, the intensity per line still decreases, requiring greater telescope sensitivity. Since the moment of inertia increases with more carbons, all transitions with any significant population will be at centimeter wavelengths. The origin of the diffuse interstellar bands (DIBs) has been a mystery since their discovery many decades ago. It has been suggested that ”cumulenes” are the carriers of the DIBs. Cumulenes are carbon chains with double bonds, instead of the alter- nating single/triple bonds of other interstellar carbon chains. The longest cumulene currently known to occur in pre-protostellar clouds is H2C6, in which it has an abundance of 4 10−11 (Langer et al. 1997). As shown in Fig. 5.2, H2C8 will probably have an abundance× of 4 10−12 and a brightness temperature of 1 mK, requiring ∼ × ∼ a very sensitive telescope. At the low temperatures of these clouds, such molecules radiate most of their energy at centimeter wavelengths. It is important therefore that the SKA be very sensitive to brightness temperature. The SKA offers significant advantages. Because of its higher angular resolution it will be better coupled to the small source regions where the emission originates, assuming that the SKA is not a dilute array. Also, we have no a priori knowledge of where these emissions peak. Observations have shown that because the regions are evolving physically and chemically on comparable timescales, distributions can vary significantly between species. The mapping capability of the SKA will dramatically reduce the time to find the positions of strongest emission and allow the different evolutionary stages of closely adjacent regions to be studied. For searches for new spectral lines whose frequencies may be uncertain, a resolution of 0.15 km/s is optimal. ∼

5.3 The Search for Extraterrestrial Intelligence

Today, many researchers believe we are on the brink of ﬁnding convincing evidence for extraterrestrial life a discovery that would generate widespread public interest and be 5.3. THE SEARCH FOR EXTRATERRESTRIAL INTELLIGENCE 161

10-6 C H C H 3 4

C H 6 10-7 HC N 3 C H H C 5 C H 2 2 8 -8 C H ) 7 2 10 C H 3 2 C H HC N 9 5 10-9 n(X)/n(H HC N 7

H C H C H C 2 5 2 6 -10 2 3 HC N 10 9

H C 2 4 10-11 H C H C 2 7 2 8 1 2 3 4 5 6 7 8 9 10 Number of Carbons (C ) n

Figure 5.2: The predicted and observed abundances of various carbon-chain species from Langer et al (1997) illustrates the need for greater sensitivity in the study of interstellar organic chemistry. Carbenes are colored blue with the solid triangles representing observed abundances. The black diamond is a measurement of cyclo- propenylidene, a ring-shaped molecule. Carbynes are colored red, and cyanopolyynes green. 162 CHAPTER 5. FORMATION AND EVOLUTION OF LIFE

Survey Freq. Range # of stars Detect EIRP EIRP mean span or Thresh. at 4 ly at 1000 ly (MHz) (MHz) % of sky (W m−2) (W) (W) Phoenix (SETI Inst.) Arecibo 2100 1800 280 8E-27 1E+08 9E+12 Parkes 2100 1800 156 2E-25 4E+09 2E+14 NRAO 2100 1800 600 3E-25 5E+09 3E+14 Serendip III(Berkeley) 29 12 28% 3E-25 5E+09 3E+14 Serendip IV (Berkeley) 1520 200 28% 1E-24 2E+10 1E+15 META(Harvard) 1420 3 68% 2E-23 3E+11 2E+16 BETA(Harvard) 1580 320 68% 1E-22 2E+12 1E+17 SKA 5500 9000 1E+05 8E-29 1E+06 9E+10

Table 5.1: Parameters of Selected SETI Surveys of profound philosophical and scientific importance. Several avenues of research could lead to this discovery. The first is the quest for the physical evidence of biology, living or remnant, on the planets and moons of our solar system. Mars, Europa, and Titan are nearby worlds where life may have sprung up, and all are the targets of current or planned space missions. Needless to say, such missions are costly, and substantial resources will be expended during the coming decades in the effort to learn if any of these cosmic neighbors has been, or is, a home to living things (NASA 1997). A second approach to finding extraterrestrial life is inspired by the recent discovery of planets around nearby stars (Mayor et al. 1997, Marcy and Butler 1998). We now know that at least a few percent of solar-type stars have planets, nd the fraction could be far higher. This discovery has spurred major new experiments that will use large, space-borne telescopes to find the small, Earth-like planets that are still beyond our reach. These telescopes would be able to collect the light of such planets, providing scientists with the opportunity to study their spectra. This leads to an intriguing possibility, the opportunity to detect primitive life at a great distance. Earths atmosphere contains a high proportion of oxygen and methane, a direct consequence of the presence of life on this planet. If studies of other planets also reveal the spectral signature of these gasses, we could infer that at least simple life exists on these distant worlds. The third approach to finding biology in the cosmos is, for both the public and much of the scientific community, the most interesting. By looking for artificially- generated radio or light signals from other star systems, we can hope to prove the existence of not just life, but intelligent, technically sophisticated life. This is the only feasible method we have to find out if we share the universe with other sentient beings. Nearly all searches for extraterrestrial intelligence (SETI) have been made using large radio telescopes operating at microwave frequencies. Table 5.3 summarizes the basic parameters of the most sensitive of recent SETI programs. 5.3. THE SEARCH FOR EXTRATERRESTRIAL INTELLIGENCE 163

Surveys that are not yet complete are shown in italics. The parameters for these searches are estimated. The Phoenix survey targets roughly one thousand nearby sun-like stars, and achieves high sensitivity and wide frequency coverage because it can spend many hours per star. The Berkeley and Harvard programs are scanning surveys, observing large portions of the sky but at the cost of lower sensitivity and narrower frequency coverage. The last columns of the table show the required equivalent isotropic power (EIRP) for a transmitter to be detected at the distance of the nearest star (4 light-years) and at 1,000 light-years. There are about one million solar-type stars within the latter distance. For comparison, note that a typical TV station has an EIRP of 106 watts, while Earths most powerful transmitter, the Arecibo planetary radar, has an EIRP of 1013 watts. Perusal of the table shows that current SETI efforts operate at sensitivity levels that leave much to be desired. The SKA would substantially improve our ability to detect signals, and consequently might lead to an early SETI success. The table lists the expected performance of a first cut SETI survey that could be undertaken with the SKA in a shared mode, using 20% of the beams over five years. This SKA survey is markedly better than current searches in several ways. To begin with, the sensitivity would be improved by two orders of magnitude. For the first time, we would be able to detect civilizations radiating signals similar to our own at interstellar distances, and be capable of finding modest radio beacons from a significant fraction of our galaxy. In addition, the availability of many beams and the large instantaneous bandwidth allow an improvement in search speed of about two orders of magnitude. This would allow an increase in the number of star systems examined from 103 to 105 or even 106, as well as permitting a search across a wider bandwidth. Finally, the expected interference mitigation schemes planned for the SKA, combined with the ability to make observations of many stars simultaneously, would likely permit SETI experiments to overcome the growing problem of terrestrial interference, a problem that might make such observations impossible with conventional radio telescopes. The SKA would enable an extraordinary advance for one of humankinds most exciting explorations, the effort to learn our place in the biological universe. Bibliography

[1] M.C. Anderson, L. Rudnick, 1995, Ap. J., 441, 307.

[2] Anderson, M. C. & Rudnick, L. 1996, “Sites of Relativistic Particle Acceleration in Supernova Remnant Cassiopeia A,” Astrophys. J., 456, 234–249.

[3] Antonucci & Ulvestad 1988, Ap.J., 330, l97.

[4] Allen, C. W. 1973, Astrophysical Quantities, 3rd Edition (London: Athlone).

[5] Allen, R. J., Goss, W. M., Ekers, R. D., & de Bruyn, A. G. 1976, A&A, 48, 253.

[6] Allen, S.W. & Fabian, A.C. 1994, MNRAS 269, 409.

[7] Anglada, G. 1996, in Radio Emission from the Stars and the Sun, Ed. A.R.Taylor & J.M.Paredes, p3.

[8] Arimoto N., Jablonka P. 1991, A&A, 249, 374.

[9] Arimoto N., Yoshii Y. 1987, A&A, 173, 23.

[10] Aschenbach, B., Egger, R., & Tr¨umper, J. 1995, Nature, 373, 587.

[11] Avery, L.W. 1991, in The Radio Schmidt Telescope, P. Dewdney and T. Lan- decker (eds.), p. 81.

[12] Baan, W.A. 1989, Ap.J., 338, 304.

[13] Baan, W.A. & Haschick, A.D. 1995, Ap.J., 454, 745.

[14] Baan, W.A. et al. 1992, Ap.J., 396, L99.

[15] Baan, W.A., Haschick, A.D., & Uglesich, R. 1993, Ap.J., 415, 140.

[16] Baan, W.A., Wood, P.A.D., & Haschick, A.D., 1982, ApJ 260, L49.

[17] Bahcall, J. N., & Ekers, R. D. 1969, Ap.J., 157, 1055

[18] Bartel, N., Chandler, J.F., Ratner, M.I., Shapiro, I.I., Pan, R., and Cappallo, R.J. 1996, A.J., 112, 1690.

164 BIBLIOGRAPHY 165

[19] Bastian, T. S. 1990, Solar Phys., 130, 265.

[20] Bastian, T. S., & Bookbinder, J. A. 1987, Nature, 326, 678.

[21] Bastian, T. S., & Bookbinder, J. A., Dulk, G. A., & Davis, M. 1990, ApJ, 353, 265.

[22] Bastian, T. S., Benz, A. O., & Gary, D. 1998, ARA&A, in press.

[23] Barvainis R., Tacconi L., Antonucci R., Coleman P. 1994, Nature, 371, 586.

[24] Beck R., Brandenburg A., Moss D.,Shukurov A., Sokoloﬀ D., 1996, Ann.Rev.Astron.Astrophys. 34, 155.

[25] Benz, A. O. 1985, Solar Phys., 96, 357.

[26] Benz, A. O., Conway, J., & G¨udel, M. 1998, A&A, 331, 596.

[27] Bester, M., Danchi, W. C., Hale, D., Townes, C. H., Degiacomi, C. G., M´ekarnia, D., & Geballe, T. 1996, ApJ, 463, 336.

[28] de Blok, W.J.G. & McGaugh, S.S. 1997, MNRAS

[29] Bloom, J. S., Frail, D. A., Kulkarni, S. R., Djorgovski, S. G. et al. 1998a, ApJ, in press (astro-ph/9808319).

[30] Bloom, J. S., Sigurdsson, S. & Pols, O. R. 1998b, MNRAS, in press, astro- ph/9805222.

[31] Bonnet, H., Mellier, Y. 1995, A&A 303, 331.

[32] Bonnet, H., Mellier, Y., Fort, B. 1994, ApJ 427, L83.

[33] Braatz, J. A., Wilson, A. S. & Henkel, C. 1996, ApJS 106, 51.

[34] Braatz, J. A., Wilson, A. S. & Henkel, C. 1997, ApJS 110, 321.

[35] Branch, D., Livio, M., Yungelson, L. R., Boﬃ, F. R., & Baron, E. 1995, PASP, 107, 1019.

[36] Broten, N.W., Macleod, J.M., Vall´ee, J.P. 1988, Astrophys. Space Sci., 141, 303.

[37] de Bruyn, A. G. 1973, A&A, 26, 105.

[38] Carr, T.D. and Wang, L. 1990, in Low Frequency Astrophysics from Space, eds. N.E. Kassim and K.W. Weilder (Heidelberg: Springer-Verlag) p. 113.

[39] Cannon, R. 1995, Astron. Soc. of Australia. Proc. 12, p258. 166 BIBLIOGRAPHY

[40] Cen, R., Miralda-Escud´e, J., Ostriker, J.P. & Rauch, M., 1994, ApJ. 437, L9.

[41] Chandler, C. J., Terebey, S., Barsony, M., Moore, T. J. T. and Gautier, T. N. 1996, Ap.J., 471, 308

[42] Chen et al 1995, Ap.J., 450, L71.

[43] Chevalier, R. A. 1982a, Ap.J., 259, 302.

[44] Chevalier, R. A. 1982b, Ap.J., 259, L85.

[45] Chevalier, R. A. 1998, ApJ, 499, 810.

[46] Claussen, M.J., Frail, D.A., Goss, W.M., Gaume, R.A. 1997, Astrophys. J., 489, 143.

[47] Condon, J.J., 1992, Ann.Rev.A&A 30, 575.

[48] Conti, P. S., et al. 1983, ApJ, 274, 302.

[49] Condon, 1992 AR A.&A., 30, 575.

[50] Conway J.E. 1996, in Second Workshop on Gigahertz Geaked Spectrum and Compact Steep Spectrum Sources, Ed. I.A.G.Snellen et al., Leiden University, p198.

[51] Couchman, H. M. P., Thomas, P. A., & Pearce, F. R. 1995, Ap.J., 452, 797.

[52] Costa, E. et al., 1997. Nature 387, 783.

[53] Dabrowski Y., Lasenby A., Saunders R., 1995, MNRAS 277, 753.

[54] Dekel, A., 1994 Ann. Rev. A&A 32, 371.

[55] Detweiler, S., 1979, ApJ 234, 1100.

[56] Dey, A., Spinrad, H., Stern, D., Graham, J. R., & Chaﬀee, F. H. 1998, Ap.J., 498, L93

[57] Dhawan V., Kellermann K.I.,Romney J.D. 1998, ApJ 498, L111.

[58] Dickey, J.M., 1997, Astrophys. J., 488, 258.

[59] Drake, S. A., Simon, T., & Linsky, J. L. 1993, ApJS, 71, 905.

[60] Drake, S. A., Simon, T., & Brown, A. 1993, ApJ, 406, 247.

[61] Drummond, J. D., Fugate, R. Q., Christou, J. C., and Hege, E. K. 1998, Icarus, 132, 80.

[62] Dulk, G. A., & Gary, D. E. 1983, A&A, 124, 103. BIBLIOGRAPHY 167

[63] Duncan, A.R., Haynes, R.F., Jones, K.L., Stewart, R.T. 1997, MNRAS, 291, 279.

[64] Duric, N., Gordon, S. M., Goss, W. M., Viallefond, F., & Lacey, C. 1995, “The Relativisitic ISM in M33: Role of the Supernova Remnants,” Astrophys. J., 445, 173.

[65] Eggen et al 1962, Ap.J., 136, 748.

[66] Elitzur,M., 1992, Astrophysical masers, Kluwer, Dordrecht.

[67] Ellison, D. C., et al. 1984, “Supernova Remnants and the Physics of Strong Shocks,” Pub. Astron. Soc. Pac., 701, 780.

[68] Erickson, W. C., McConnell, D., & Anantharamaiah, K. R., 1995, “Low- Frequency Carbon Recombination Lines in the Central Regions of Galaxy,” Astrophys. J., 454, 125.

[69] Evans A.S., Sanders D., Mazzarella J.M., Solomon P.M., Downes D., Kramer C., Radford S.J.E. 1996, Ap.J., 457, 658.

[70] Fahlman, G., Kaiser, N., Squires, G., & Woods, D. 1994, Ap.J., 431, L71.

[71] Field, G. B. 1958, Proc. I.R.E., 46, 240

[72] ——— . 1959, Ap.J., 129, 551.

[73] Filippenko, A. V. 1988, AJ, 96, 1941.

[74] Filippenko, A. V. 1997, AnnRevAstAp, 35, 309.

[75] Fischer, P., & Tyson, J.A. 1997, AJ 114, 14.

[76] Fishman, G.J. & Meegan, C.A. 1995. ARAA, 33, 415.

[77] Foster, R. S. & Backer, D. C., 1990, ApJ 361, 300.

[78] Frail, D.A., et al., 1997. Nature 389, 261.

[79] Frail, D.A., Goss, W.M., Slysh, V.I. 1994, Ap. J., 424, L111.

[80] Franx, M., Illingworth, G.D., Kelson, D.D., van Dokkum, P.G., & Tran, K.-V. 1997, Ap.J., 486, L75.

[81] Frayer T., Brown R., Vanden Bout P. 1994, Ap.J., 433, 5.

[82] Gaensler, B.M., Manchester, R.N., Green, A.J. 1998, MNRAS, 296, 813.

[83] Galama, T. et al. 1998, Nature (in press). 168 BIBLIOGRAPHY

[84] Gary, D. E., & Hurford, G. J. 1994, ApJ, 420, 903.

[85] Gilliland, R. & Dupree, A. K. 1996a, in IAU Symposium no. 176: Stellar Surface Structure, eds Strassmeier, K. G. & Linsky, J. L. (Kluwer:Dordrecht), 165.

[86] Gilliland, R. & Dupree, A. K. 1996b, Ap.J., 436, L29.

[87] Gliese, W., & Jahreiss, H. Third Catalogue of Nearby Stars, Astronomisches Rechen-Institut Heidelberg, Germany; software version.

[88] Gopalswamy, N., White, S. M., & Kundu, M. R. 1991, ApJ, 379, 366.

[89] van Gorkom, J.H,. Carilli, C.L., Stocke, J.T., Perlman, E.S. & Shull, J.M., 1996, AJ. 112,, 1397.

[90] Governato, F., Gardner, J. P., Stadel, J., Quinn, T., & Lake, G. 1997, preprint, astro-ph/9710140.

[91] Gray, A.D., Landecker, T.L., Dewdney,P.E., & Taylor, A.R. 1998, Nature, 393, 660.

[92] Gray, A.D., Landecker, T.L., Dewdney, P.E., Taylor, A.R., Willis, A.G., Nor- mandeau, M. 1999, Ap. J., in press.

[93] Greaves, J. S., Holland, W. S., Moriarty-Schieven, G., Jeness, T., Dent, W. R. F., Zuckerman, B., McCarthy, C., Webb, R. A., Butner, H. M., Gear, W. K., and Walker, H. J. 1998, Ap.J., 506, L133.

[94] Green 1984, MNRAS, 209, 449.

[95] Green, D. A. 1991, “Limitations Imposed on Statistical Studies of Galactic Supernova Remnants by Observational Selection Eﬀects,” Pub. Astron. Soc. Pac., 103, 209.

[96] G¨udel, M. 1994, ApJS, 90, 743.

[97] G¨uedel, M. et al. 1997 Ap.J., 483, 947.

[98] G¨udel, M., & Benz, A. O. 1989, A&A, 211, L5.

[99] G¨udel, M., & Zucker, A. 1999, in preparation.

[100] G¨udel, M., Benz, A. O., Bastian, T. S., F¨urst, E., Simnett, G. M., Davis, R. J. 1989, A&A, 220, L5.

[101] G¨udel, M., Benz, A. O., Schmitt, J. H. M. M., & Skinner, S. L. 1996, ApJ, 471, 1002. BIBLIOGRAPHY 169

[102] G¨udel, M., Guinan, E. F. & Skinner, S. L. 1998, in 10th Cambridge Workshop on Cool Stars, Stellar Systems, and the Sun. Ed. R. A. Donahue & J. A. Bookbinder (San Francisco: ASP), in press.

[103] Gunn, J.E. & Weinberg, D.H. 1995, in Wide-Field Spectroscopy and the Distant Universe, Ed. S.J.Maddox & A. Arag´on-Salamanca

[104] Gurvits L.I., Roland J., Demyk K. 1998, in Observational Cosmology with the New Radio Surveys, Ed. M.N.Bremer, N.Jackson, I.P´erez-Fournon, Kluwer Acad. Press, p227.

[105] Gurvitz, L., Kellermann, K. and Frey, S. 1999, A. & A., 342, 378. [106] Haardt, F. & Madau, P. 1996, Ap.J., 461, 20.

[107] Han, J.L., Manchester, R.N., Berjhuijsen, E.M., Beck, R. 1997, A. & A., 322, 98.

[108] Han J.L., Beck R., Berkhuijsen E.M., 1998, A. & A., 335, 1117. [109] Helou et al. 1985, Ap.J. Lett., 298, L7.

[110] Helou, G., 1991,in The Interpretation of Modern Synthesis Observations of Spi- ral Galaxies, ed. N. Duric and P. C. Crane, ASP Conference series No. 18, p. 125. [111] Henkel, C., Baan, W.A., & Mauersberger, R. 1993, Astron. Ap. Rev., 3, 47.

[112] Herbst, E. and Leung, C.M. 1989, Ap.J. Supp., 69, 271. [113] Hernquist, L., Katz, N., Weinberg, D.H. & Miralda-Escud´e, J., 1996, ApJ. 457, L51. [114] Herrnstein et al., 1997, ApJ 475, L17.

[115] Hewish, A., Bell, S. J., Pilkington, J. D. H., Scott, P. F. & Collins, R. A., 1968, Nat 217, 709.

[116] Hirahara, Y., Suzuki, H., Yamamoto, S., Kawaguchi, K., Kaifu, N., Ohishi, M., Takano, S., Ishakawa, S.-I. and Masuda, A. 1992, Ap.J., 394, 539. [117] Hjellming, R. M. & Newell, R. T. 1983, ApJ, 275, 704.

[118] Hoekstra, H., Franx, M., Kuijken, K. & Squires, G. 1998, ApJ, in press.

[119] Hopkins, A.M., Mobasher, B., Cram, L., & Rowan-Robinson, M., 1998 MNRAS 296, 839. [120] Hoyle, F. 1959, in Paris Symposium on Radio Astronomy, Ed. R.N.Bracewell, Stanford Univ. Press, p529. 170 BIBLIOGRAPHY

[121] Hudson, R. S., and Ostro, S. J. 1994, Science, 263, 940.

[122] Hughes, D., et al., 1998, preprint, astro-ph/9806297

[123] Hulse, R. A. & Taylor, J. H., 1975, 195, ApJ L51.

[124] Humphreys, R., Nichols, M., & Massey P. 1985, AJ, 90, 101.

[125] Ikuta C., Arimoto N., Sofue S., Taniguchi Y. 1997, PASJ, 49, 647.

[126] Jackson, P. D., Kundu, M.R., & White, S. M. 1987, ApJ, 316, L85.

[127] Johnston, K. J., Lamphear, E. J., Webster, Jr., W. J., Lowman, Jr., P. D., Seidelmann, P. K., Kaplan, G. H., Wade, C. M., and Hobbs, R. W. 1989, A.J., 98, 335.

[128] Jones, T. W., et al. 1998, “1051 Ergs: The Evolution of Shell Supernova Rem- nants,” Publ. Astron. Soc. Paciﬁc, 110, 125.

[129] Kaiser, N., & Squires, G. 1993, ApJ 404, 441.

[130] Kaiser, N., Squires, G., & Broadhurst, T. 1995, ApJ 449, 460.

[131] Kargel, J. S. 1994, J.G.R., 99. NO. E10, 21,129.

[132] Karovska, M., Hack, W., Raymond, J., & Guinan, E. 1997, ApJL, 482, 175.

[133] Kaspi, V. M., Taylor, J. H. & Ryba, M. 1994, ApJ 428, 713.

[134] Kassim, N. E., Weiler, K. W., Erickson, W. C., & Wilson, T. C. 1989, “Improved Estimates of Galactic HII Region Emission Measures and Filling Factors: Low- Frequency VLA Observations near Sharpless 53”, Astrophys. J., 338, 152.

[135] Kassim, N. E. 1990,“HII Regions in Absorption at Low Frequencies,” in Low Frequency Astrophysics from Space, eds. N. E. Kassim & K. W. Weiler (Hei- delberg: Springer-Verlag) p. 144.

[136] Kassim, N. E., Perley, R. A., Carilli, C. L., Harris, D. E., & Erickson, W. C. 1995, “Low Frequency Observations of Cygnus A,” in Cygnus A: Study of a Radio Galaxy, eds. C. L. Carilli & D. E. Harris, (Cambridge: Cambridge Uni- versity Press) p. 182.

[137] Kassim, N. E., Perley, R. A., Dwarakanath, K. S., & Erickson, W. C. 1995, “Evidence for Thermal Absorption inside Cassiopeia A,” Astrophys. J., 455, L59.

[138] Kellermann K.I., 1993, Nature 361, 134.

[139] Kemball, A. J., & Diamond, P. J. 1997, ApJL, 481, 111. BIBLIOGRAPHY 171

[140] Kim, K.-T., Kronberg, P.P., Landecker, T.L. 1988, A. J. 96, 704.

[141] Kim, K.-T., Kronberg, K.K., Dewdney, P.E. & Landecker, T.L. 1990, Ap. J., 355, 29.

[142] Krauss L.M., Schramm D.N., 1993, ApJ 405, L43.

[143] Kronberg, P.P, et al 1985, Ap.J., 291, 693.

[144] Kronberg, P. P. 1990, “Radio Emission from Intergalactic Gas and Its Implica- tions for Low Frequency Astronomy in Space,” in Low Frequency Astrophysics from Space, eds. N. E. Kassim and K. W. Weiler (Heidelberg: Springer-Verlag) p. 262.

[145] Kulkarni, S. R. et al. 1998, Nature (in press)

[146] Kundu, M. R.,Gopalswamy, N., White, S. M., Cargill, P., Schmahl, E. J., & Hildner, E. 1989, ApJ, 347, 505.

[147] Kuiper, T. B. H., Langer, W. D., Velusamy, T. 1996, Ap.J., 468, 761.

[148] Kwok, S. 1975, ApJ, 198, 583.

[149] Lahav, O. 1995, in Mapping, Measuring and Modelling the Universe, Ed. P. Coles & V. Martinez.

[150] Lang, K. R., et al. 1993, ApJ, 419, 398.

[151] Lang, K. R., Willson, R. F., Smith, K. L., & Strong, K. T. 1987, ApJ, 322, 1044.

[152] Langer, W.D., Velusamy, T., Kuiper, T.B.H., Peng, R., McCarthy, M.C., Travers, M.J., Kovacs, A., Gottlieb, C.A. and Thaddeus, P. 1997, Ap.J. Letters, 480, L63.

[153] Lanzetta, K., Wolfe, A. & Turnshek, D. 1995, ApJ 440, 435.

[154] Leahy,D.A., Roger, R.S., Ballantyne, D. 1997, A. J., 114, 2081.

[155] Lim, J., & White, S. M. 1995, ApJ, 453, 207.

[156] Lim, J., Carilli, C. L., White, S. M., Beasley, A. J., & Marson, R. G. 1998, Nature, 392, 575

[157] Linsky, J. L., Drake, S. A., & Bastian, T. S. 1992, ApJ, 393, 341.

[158] Lisenfeld, U., V¨olk, H.J. & Xu, C., 1996, A&A 306, 677. 172 BIBLIOGRAPHY

[159] Longair, M. S. 1990, “Cosmic Rays and the Galactic Radio Background Emis- sion,” in Low Frequency Astrophysics from Space, eds. N. E. Kassim & K. W. Weiler (Heidelberg: Springer-Verlag) p. 227.

[160] Lonsdale, C.J., Lonsdale, C.J., Diamond, P.J., & Smith, H.E. 1998, Ap.J., 493, L13.

[161] Luppino, G.A., Kaiser, N. 1997, ApJ 475, 20.

[162] Luu, J. X., and Jewitt, D. C. 1996, A.J., 112, 2310.

[163] Luu, J. X., and Jewitt, D. C. 1998, Ap.J., 502, L91.

[164] Lyne, A. G., et al. 1998, “The Parkes Southern Pulsar Survey—II. Final results and Population Analysis,” Mon. Not. R. Astron. Soc., 295, 743.

[165] Madau, P., Meiksin, A., & Rees, M. J. 1997, Ap.J., 475, 429.

[166] Madau, P., 1998, in The Hubble Deep Field and the Early Evolution of Galaxies, in press. (astro-ph/9707141).

[167] Madau, P., Pozzetti, P., & Dickinson, M. 1998, Ap.J., 498, 106.

[168] Marcy G. & Butler P., 1998, ARAA 36, in preparation.

[169] Marti, J. Rodriguez, L. and Reipurth 1998, Ap.J., 502, 337.

[170] Mather, J. C., et al., 1994, Ap.J., 420, 439.

[171] Mayor, M., Queloz, D., 1995, Nature, 378, 355.

[172] Mayor M., Queloz M., Udry S. & Halbwachs J.L, 1997, in Astronomical and Bio- chemical Origins and the Search for Life in the Universe, Ed. C.B. Cosmovici, S. Bowyer & D. Werthimer.

[173] McCarthy, M.C., Grabow, J.-U., Travers, M.J., CHen, W., Gottlieb, C.A. and Thaddeus, P. 1998, Ap.J. Lett., 498, L89.

[174] McCaughrean, M. J., and O’Dell, C. R. 1996, A.J., 111, 1977.

[175] Meiksin, A., & Madau, P. 1993, Ap.J., 412, 34.

[176] Men’shchikov & Henning 1997, A. & A., 318, 879.

[177] Metzger, M.R. et al., 1997. Nature 387, 879.

[178] Millar, T. J. and Herbst, E. 1990, A. & A., 231, 466.

[179] Miyoshi, M., Moran, J. M., Herrnstein, J., Greenhill, L. J., Nakai, N., Diamond, P. & Inoue, M. 1995, Nat 373, 127. BIBLIOGRAPHY 173

[180] Montes, M. J., Van Dyk, S. D., Weiler, K. W., Sramek, R. A., & Panagia, N. 1998a, ApJ, in press.

[181] Montes, M. J., Weiler, K. W., Van Dyk, S. D., Sramek, R. A., Panagia, N., & Park, R. 1998b, in preparation.

[182] Mundy et al. 1993, Icarus. 106, 11.

[183] Muxlow et al 1994 MNRAS, 266, 455.

[184] NASA, 1997, The Space Science Enterprise Strategic Plan

[185] Newell, R. T. & Hjellming, R. M. 1982, ApJL, 263, 85.

[186] Norris, R.P., Byleveld, S.E., Diamond, P.J., Ellingsen, S.P., Ferris, R.H., Gough, R.G., Kesteven, M.J., McCulloch,P.M., Phillips, C.J., Reynolds, J.E., Tzioumis, A.K., Takahashi, Y., Troup, E.R., & Wellington, K.J., 1998, Ap.J., in press.

[187] Norris, R.P., Whiteoak, J.B., Caswell, J.L., Wieringa, M.H., & Gough, R.G., 1993, Ap.J., 412, 222.

[188] Ohta K, Yamada T., Nakanishi K., Kohno K., Akiyama M., Kawabe R. 1996, Nature, 382, 426.

[189] Olano, C. A., Walmsley, C. M., and Wilson, T. L. 1988, A. & A., 164, 159.

[190] Ostro, S. J., Campbell, D. B., and Shapiro, I. I. 1985, Science, 229, 442.

[191] Ostro, S.J. et al. 1990, Science 248, 1523.

[192] Ostro, S. J., Campbell, D. B., Chandler, J. F., Hine, A. A., Hudson, R. S., Rosema, K. D., and Shapiro, I. I. 1991, Science, 252, 1399.

[193] Ostro, S.J. et al. 1995, Science 270, 80.

[194] Paczy´nski, B. 1998, ApJ, 484, L45.

[195] van Paradijs, J. et al., 1997, Nature 386, 686.

[196] Parker, E. N. 1988, ApJ, 330, 474

[197] de Pater, I. & C.K. Goertz, 1990, J.Geophys.Res. 95, 39.

[198] Payne, H. E., Anantharamaiah, K. R., & Erickson, W. C. 1994, “High Rydberg State Carbon Recombination Lines Towards Cas A: Physical Conditions and a New Class of Models,” Astrophys. J., 430, 690.

[199] Pearson T.J. 1996, in Energy Transport in Radio Galaxies and Quasars, Ed. P.E.Hardee, A.H.Bridle & J.A.Zensus, ASP, p97. 174 BIBLIOGRAPHY

[200] Pearson T.J. et al. 1994, in Compact Extragalactic Radio Sources, Ed. J.A.Zensus & K.I.Kellermann, NRAO/Socorro, p1.

[201] Pedlar et al 1998, MNRAS, (submitted).

[202] Perlmutter, S., et al. 1998, ApJ, in press.

[203] Phillips, C.J., Norris, R.P., Ellingsen, S.P., & Rayner, D. P. 1998, MNRAS in press.

[204] Phinney, E. S. & Kulkarni, S. R., 1994, ARAA 32, 591.

[205] Podsiadlowski, Ph., Joss, P. C, & Hsu, J.J.L. 1992, ApJ, 391, 246.

[206] Porter, A. C., & Filippenko, A. V. 1987, AJ, 93, 1372

[207] Rajagopal, M. & Romani, R. W., 1995, ApJ 446, 543.

[208] Ramaprakash, A. N., Kulkarni, S. R. et al. 1998, Nature, 393, 43.

[209] Reichart, D. E. et al. 1998, ApJ submitted (astro-ph/9806082)

[210] Redman, R. O., Feldman, P. A., and Matthews, H. E. 1998a, A.J., 116, 1478

[211] Redman, R. O., Feldman, P. A., and Matthews, H. E. 1998b, J.R.A.S.C., in press.

[212] Reid, M.J. & Menten, K.M. 1996, in Radio Emission from the Stars and the Sun, Ed. A.R.Taylor & J.M.Paredes, p87.

[213] Reid, M. J. & Menten, K. M. 1997, ApJ, 476, 327.

[214] Readhead A.CS. 1995, Proc. Nat Acad. Sci USA, 92, 11447.

[215] Reynolds, S. P. & Ellison, D. C. 1992, “Electron Acceleration in Tycho’s and Kepler’s Supernova Remnants. Spectral Evidence of Fermi Shock Acceleration,” Astrophys. J., 399, L75.

[216] Reynolds, S.P. & Chevalier, R.A. 1984, Ap. J., 281, L33.

[217] Riess, A. G., et al. 1998, AJ, 116, 1009.

[218] Rodriguez, P. 1996, “High Frequency Radar Detection of Coronal Mass Ejec- tions,” in Solar Drivers of Interplanetary and Terrestrial Disturbances, eds. K. S. Balasubramaniam, S. L. Keil, & R. N. Smartt (San Francisco: Astronom- ical Society of the Paciﬁc) p. 180.

[219] dos Santos, P.M., & Lepine, J.R.D, 1979, Nat 278, 34.

[220] Schneider, P. 1995, A&A 302, 639. BIBLIOGRAPHY 175

[221] Schneider, P. & Seitz, C. 1995, A&A 294, 411.

[222] Schneider, P., van Waerbeke, L., Mellier, Y., Seitz, S. & Fort, B. 1998, A&A 333, 767.

[223] Scoville, N.Z. et al. 1998, Ap.J., 492, L107.

[224] Seaquist, E.R. 1996, in High Sensitivity Radio Astronomy, Ed. N.Jackson & R.J.Davis, p23.

[225] Scott, D., & Rees, M. J. 1990, MNRAS, 247, 510.

[226] Seitz, C., Kneib, J.-P., Schneider, P., & Seitz, S. 1996, A&A 314, 707.

[227] Schlegel, E. M. 1990, MNRAS, 244, 269.

[228] Schlegel, E. M. 1996, AJ, 111, 1660.

[229] Schmitt, J. H. M. M. 1997, A&A, 318, 215.

[230] Shull, J.M., Penton, S.V., Stocke, J.T., Giroux, M.L., van Gorkom, J.H., Lee, Y.H. & Carilli, C.L., 1998, ApJ. in press.

[231] Sigad, Y.,Eldar, A.., DEKEL, AVISHAI; STRAUSS, M.A. & Yahil, A., 1998, ApJ. 495, 516.

[232] Silk J., Spaans M. 1997, Ap.J., 488, L79,

[233] Simard-Normandin, M., Kronberg, P.P. & Button, S. 1981, Ap. J. Suppl., 45, 97.

[234] Smail, I., Ellis, R. S., Dressler, A., Couch, W. J., Oemler, A. Jr., Sharples, R. M., Butcher, H. 1997, Ap.J., 497, 70.

[235] Smith, H.E., Lonsdale, C.J., Lonsdale, C.J., Diamond, P.J., & 1998, Ap.J., 493, L17.

[236] Squires, G., & Kaiser, N. 1996, ApJ 473, 65.

[237] Squires, G., Kaiser, N., Fahlman, G., Babul, A., & Woods, D. 1996, ApJ 469, 73.

[238] Soderblom, D. R., Stauﬀer, J. R., Hudon, J. D., & Jones, B. F. 1993, ApJS, 85, 315.

[239] Solomon, P.M., Downes, D., and Radford, S. J. E. 1992, Ap.J., 398, L29.

[240] Sokoloﬀ D.D., Bykov A.A., Shukurov A., Berkhuijsen E.M., Beck R. and Poezd A.D., 1998, MNRAS 299, 189. 176 BIBLIOGRAPHY

[241] Staveley-Smith, L, Allen, D.A, Chapman, J.M., Norris, R.P. & Whiteoak, J.B.,(1989), Nature, 337, 625.

[242] Stecklum, B., Kauﬂ,H.-U., Henning, Th., feldt, M., Eckart, A., & Nyman, L.-A., 1998, ESO press release 08/98.

[243] Stelmach J., 1994, ApJ 428, 61.

[244] Stepanas P.G. & Saha P., 1995, MNRAS 272, L13.

[245] Subrahmanyan, R. & Goss, W. M. 1996, “Electorn Temperatures in the Galactic HII Regions W43 and M17,” Mon. Not. R. Astron. Soc., 281, 239.

[246] Suzuki, H., Yamamoto, S., Ohishi, M., Kaifu, N., Ishikawa, S.-I., Hirahara, Y., and Takano, S. 1992, Ap.J., 392, 551.

[247] Taylor, G.B., Frail, D.A., Kulkarni, S.R., Shepherd, D.S., Feroci, M. & Frontera, F. 1998, Ap.J., in press.

[248] Taylor, J. H. & Weisberg, J., 1982, ApJ 253, 908.

[249] Taylor, J. H. & Cordes, J. M. 1993, “Pulsar Distances and the Galactic Distri- bution of Free Electrons,” Astrophys. J., 411, 674.

[250] Tedesco, E. F., Veeder, G. J., Fowler, J. W., and Chillemi, J. R. 1992, ”The IRAS Minor Planet Survey: Final Report”, PL-TR-92-2049.

[251] Tsuboi M., Nakai N. 1992, PASJ 44, 241.

[252] Tuthill, P. G., Haniﬀ, C. A., & Baldwin, J. E. 1997, MNRAS, 285, 529.

[253] Tyson, J., Valdes, F., & Wenk, R. 1990, ApJ 349, L1.

[254] Unger et al 1984, MNRAS, 211, 783.

[255] Uomoto, A. 1986, ApJ, 310, L35.

[256] Uyaniker, B., F¨urst, E., Reich, W., Reich, P., Wielebinski, R., 1998, A. & A. Suppl., in press.

[257] Uyaniker, R., Reich, W., F¨urst, E., Reich, P., Wielebinski, R. 1999, A. & A. Suppl., in press.

[258] Van Dyk, S. D., Weiler, K. W., Sramek, R. A., & Panagia, N. 1993a, ApJ, 419, L69.

[259] Van Dyk, S. D., Sramek, R. A., Weiler, K. W., & Panagia, N. 1993b, ApJ, 409, 162. BIBLIOGRAPHY 177

[260] Van Dyk, S. D., Weiler, K. W., Sramek, R. A., Rupen, M. P., & Panagia, N. 1994, ApJ, 432, L115.

[261] Vermeulen R.C. 1996, in Extragalactic Radio Sources, Ed. R.Ekers, C.Fanti & L.Padrielli, Kluwer Acad. Press, p57.

[262] van de Weygaert, R & van Albada T.S. 1997, in The Westerbork Observa- tory, Continuing Adventure in Radio Astronomy, Ed. E. Raimond & R. Gene´e, Kluwer ASSL Series 207, 225.

[263] Vourlidas, A., & Bastian, T. S. 1996, ApJ, 466, 1039.

[264] Vourlidas, A., Bastian, T. S., & Aschwanden, M. J. 1997, ApJ, 489, 403.

[265] Wall, J.V., & Jackson, C.A., 1997 MNRAS 290, L17.

[266] Warren, S. J., Hewett, P. C., & Osmer, P. S. 1994, Ap.J., 421, 412.

[267] Waxman, E., Kulkarni, S. R. & Frail, D. A. 1998, ApJ, 497, 288

[268] Weaver, H.F., Williams, D.R.W, Dieter, N.H., & Lum, W.T., 1965, Nature, 208, 29.

[269] Webber, W. R. 1990, “Cosmic rays in the Galaxy and Their Implications for very Low Frequency Radio Astronomy,” in Low Frequency Astrophysics from Space, eds. N. E. Kassim & K. W. Weiler (Heidelberg: Springer-Verlag) p. 217– 226

[270] Weiler, K. W., van der Hulst, J. M., Sramek, R. A., & Panagia, N. 1981, ApJ, 243, L151.

[271] Weiler, K. W., Sramek, R. A., Panagia, N., van der Hulst, J. M., & Salvati, M. 1986, ApJ, 301, 790.

[272] Weiler, K. W., Panagia, N., & Sramek, R. A. 1990, ApJ, 364, 611.

[273] Weiler, K. W., Van Dyk, S. D., Panagia, N., Sramek, R. A. & Discenna, J. L. 1991, ApJ, 380, 161.

[274] Weiler, K. W., Van Dyk, S. D., Panagia, N., & Sramek, R. A. 1992, ApJ, 398, 248.

[275] Weiler, K. W., Van Dyk, S. D., Montes, M. J., Panagia, N., & Sramek, R. A. 1998, ApJ, 500, 51.

[276] Wendker, H. J. 1995, A&AS, 109, 177.

[277] Wieringa, M., de Bruyn, A.G., Jansen, D., Brouw, W.N., Katgert, P. 1993, A. & A., 268, 215. 178 BIBLIOGRAPHY

[278] Williams, P.M., Dougherty, S.M., Davis, R.J., van der Hucht, K.A., Bode, M.F., Setia Gunawan, D.Y.A. 1997, MNRAS, 289, 10.

[279] Wilkinson & deBruyn 1990, MNRAS, 242, 529.

[280] Wilkinson P.N. et al. 1997, in Observational Cosmology with the New Radio Sur- veys, Ed. M.N.Bremer, N.Jackson & I.P´erez-Fournon, Kluwer Acad. Publishers, p221.

[281] Wills, K. A., Pedlar, A., Muxlow, T. W. B., & Wilkinson, P. N. 1997, “Low- Frequency Observations of Supernova Remnants in M82,” Mon. Not. R. Astron. Soc., 291, 517.

[282] Wills et al 1997, MNRAS, 291, 517.

[283] Wills et al 1998, MNRAS, 298, 347.

[284] Wilner, Ho, and Rodriguez, L. 1999, in preparation.

[285] Windhorst, R.A., Mathis, D.F., & Neuschaefer, L.W., 1990 in ASP Conf. Ser., Vol. 10, Evolution of the Universe of Galaxies, ed. R. G. Kron, 389.

[286] Windhorst, R.A., Fomalont, E.B., Partridge, R.B., & Lowenthal, J.D., 1993 ApJ, 405, 498.

[287] White, S. M., Jackson, P. D., & Kundu, M. R. 1989, Ap.J. Supp., 71, 895.

[288] White, S. M., Kundu, M. R., & Jackson, P. D., 1989, A&A, 225, 112.

[289] White, S. M., Kundu, M. R., & Gopalswamy, N. 1991, ApJ, 366, L43.

[290] White, S. M., Lim, J., & Kundu, M. R. 1994, ApJ, 422, 293.

[291] Wolszczan, A. and Frail, D. 1992, Nature, 355, 145.

[292] Wolszczan, A., 1994, Science 264, 538.

[293] Wouthuysen, S. A. 1952, A.J., 57, 31.

[294] Yamada T., Ohta K., Tomita A., Takata T. 1995, A.J., 110, 1564.

[295] Zellner, B. H., Albrecht, R., Binzel, R. P., Gaﬀey, M. J., Thomas, P. C., Stors, A. D., and Wells, E. N. 1997, Icarus, 128, 83.

[296] Zensus J.A., Diamond P.J. & Napier P.J 1995, in Very Long Baseline Interfer- ometry and the VLBA, Ed. J.A.Zensus, ASP Conf Series 82.

[297] Zhang, Y., Meiksin, A., Anninos, P., & Norman, M. L. 1998, Ap.J., 495, 63. BIBLIOGRAPHY 179

[298] Zheng, W., Kriss, G.A., Teﬀer, R.C., Grimes, J.P. & Davidsen, A.F., 1997, ApJ. 475, 469.

[299] Zwaan, M.A., Briggs, F.H., Sprayberry, D. & Sorar, E. 1997, Ap.J., 490, 173.