New Observational Insight on Shock Interactions Toward Supernovae and Supernova Remnants

Item Type text; Electronic Dissertation

Authors Kilpatrick, Charles Donald

Publisher The University of Arizona.

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Link to Item http://hdl.handle.net/10150/621574 NEW OBSERVATIONAL INSIGHT ON SHOCK INTERACTIONS TOWARD SUPERNOVAE AND SUPERNOVA REMNANTS

Charles Donald Kilpatrick

A Dissertation Submitted to the Faculty of the

DEPARTMENT OF ASTRONOMY

In Partial Fulﬁllment of the Requirements For the Degree of

DOCTOR OF PHILOSOPHY

In the Graduate College

THE UNIVERSITY OF ARIZONA

2016 2

THE UNIVERSITY OF ARIZONA GRADUATE COLLEGE

As members of the Dissertation Committee, we certify that we have read the dissertation prepared by Charles Donald Kilpatrick entitled New Observational Insight on Shock Inter- actions Toward Supernovae and Supernova Remnants and recommend that it be accepted as fulﬁlling the dissertation requirement for the Degree of Doctor of Philosophy.

Date: 14 June 2016 George Rieke

Date: 14 June 2016 David Arnett

Date: 14 June 2016 John Bieging

Date: 14 June 2016 Peter Milne

Date: 14 June 2016 Nathan Smith

Final approval and acceptance of this dissertation is contingent upon the candidate’s sub- mission of the ﬁnal copies of the dissertation to the Graduate College. I hereby certify that I have read this dissertation prepared under my direction and recommend that it be accepted as fulﬁlling the dissertation requirement.

Date: 14 June 2016 Dissertation Director: George Rieke 3

STATEMENT BY AUTHOR

This dissertation has been submitted in partial fulﬁllment of requirements for an advanced degree at the University of Arizona and is deposited in the University Library to be made available to borrowers under rules of the Library.

Brief quotations from this dissertation are allowable without special permission, provided that an accurate acknowledgment of the source is made. Requests for permission for extended quotation from or reproduction of this manuscript in whole or in part may be granted by the head of the major department or the Dean of the Graduate College when in his or her judgment the proposed use of the material is in the interests of scholarship. In all other instances, however, permission must be obtained from the author.

SIGNED: Charles Donald Kilpatrick 4

ACKNOWLEDGEMENTS

I would like to thank the many mentors and collaborators who have facilitated my development as a scientist. Foremost, my advisor George Rieke has constantly supported me in all my research interests. Every one of the opportunities I have had to explore infrared, optical, millimeter, and radio astronomy, to conduct research on supernova remnants, supernovae, and transients, and to start and complete all of the projects in this dissertation are owed to George. John Bieging taught me millimeter observing and had the patience to help me become a better writer. Nathan Smith guided me toward a much deeper understanding of supernovae and gave me several opportunities to conduct my own research as an observer. Thanks to the staff scientists and colleagues who have shown a willingness and ex- citement to support my research. I would like to acknowledge the staff at the Heinrich Hertz Submillimeter Telescope, the Bok Telescope, and the Kuiper Telescope, especially Joe Hoscheidt whose hard work and dedication to maintaining Arizona telescopes extends to winding mountain drives in the dead of night. Don McCarthy and Craig Kulesa de- voted considerable time to setting up PISCES infrared camera on the Bok 90” Telescope, guiding my data reduction efforts, and reading my manuscripts for publication. Members of the Arizona Transient Exploration and Characterization collaboration, especially Pe- ter Milne, Wen-fai Fong, and Jen Andrews, have contributed countless hours to this work through informal chats, brainstorming new ideas, and help in data reduction. Dave Arnett has had an enormous inﬂuence on my research and I am grateful for his encouragement and willingness to challenge my ideas about Cassiopeia A. I would like to express my gratitude to the many scientists whose data I solicited in writing this dissertation, including Crystal Brogan for providing her 90 cm data of Galactic supernova remnants, David Moffett for providing 20 cm radio continuum data of 3C 391, Jeonghee Rho and Kristoffer Eriksen for providing Ks band imaging of Cassiopeia A, and Alex Filippenko and the Berkeley transient group for providing optical spectroscopy of PS15si. I would also like to thank Felix Aharonian and Joachim Hahn for their helpful comments on gamma-ray emission from supernova remnants. Some of the observations reported in this dissertation were obtained using the Bart J. Bok 90” Telescope operated by Steward Observatory, the MMT Observatory operated jointly by the University of Arizona and the Smithsonian Institution, and the W. M. Keck Observatory operated as a scientiﬁc partnership among the California Institute of Technol- ogy, the University of California, and NASA. The Heinrich Hertz Submillimeter Telescope is operated by the Arizona Radio Observatory with partial support from National Science Foundation grant AST 1140030. The research in this dissertation was supported by NASA through Contract Number 1255094 issued by JPL/Caltech. 5

DEDICATION

To Dad and Zack: for providing the examples to which I aspire as a scientist. To Mom: for being the ﬁrst person to take me to the library, the museum, and school. 6

TABLE OF CONTENTS

LIST OF FIGURES ...... 9

LIST OF TABLES ...... 11

ABSTRACT ...... 12

CHAPTER 1 Introduction ...... 14 1.1 The Physics of Supernova Shocks ...... 15 1.2 Dynamical Evolution of Supernova Shocks ...... 17 1.2.1 Free-Expansion Phase ...... 17 1.2.2 Snowplow Phase ...... 19 1.3 Electromagnetic Signatures of Shocks From Supernovae and Supernova Remnants ...... 21 1.3.1 Supernovae with Circumstellar Interactions ...... 21 1.3.2 Nonthermal Emission from Young Supernova Remnants . . . . . 22 1.3.3 Interactions Between Supernova Remnants and Molecular Clouds 24 1.4 This Volume ...... 26

CHAPTER 2 PS15si: A Type Ia/IIn Supernova with Late-time Rebrightening . . 27 2.1 Introduction ...... 27 2.2 Observations ...... 30 2.3 Results ...... 34 2.3.1 Photometry ...... 34 2.3.2 Spectroscopy ...... 36 2.4 Discussion ...... 44 2.4.1 PS15si Explosion Date and Maximum R band Magnitude . . . . 44 2.4.2 Luminosity of the CSM and Underlying SN ...... 47 2.4.3 Late-Time Rebrightening and Spectral Fitting of SNe Ia/IIn . . . 49 2.5 Conclusion ...... 50

CHAPTER 3 Variability in the Near-Infrared Synchrotron Emission From Cas- siopeia A ...... 52 3.1 Introduction ...... 52 3.2 Observations ...... 54 3.3 Results and Analysis ...... 55 3.3.1 Fast-Moving Features ...... 55 7

TABLE OF CONTENTS – Continued

3.3.2 Are the Ks Band Knots Dominated by Synchrotron Emission? . . 57 3.3.3 IRAC Measurements ...... 59 3.3.4 Physical Parameters of Synchrotron-Emitting Knots ...... 62 3.4 Discussion ...... 64 3.5 Conclusions ...... 65

CHAPTER 4 Interaction Between Cassiopeia A and Nearby Molecular Clouds . 66 4.1 Introduction ...... 66 4.2 Observations ...... 68 4.3 CO Line Emission Around Cas A ...... 71 4.3.1 Kinematics of CO Line Emission as Observed in J = 2 − 1 Spectra 71 4.3.2 Possible Cloud/SNR Interactions ...... 74 4.3.3 CO Line Broadening ...... 74 4.4 Analysis of Cas A Infrared Spectroscopy ...... 79 4.4.1 Spectral Features ...... 79 4.4.2 [O IV] 25.94 µm line ...... 82 4.4.3 Analysis of Low-Resolution [O IV] Emission ...... 87 4.5 Discussion ...... 91 4.5.1 Possible Mechanisms for CO Line Broadening in Regions Beyond the Cas A Shock Front ...... 92 4.5.2 Broadening by Stellar Winds ...... 92 4.5.3 Broadening by Interaction with Outﬂows of Fast-Moving Knots . 93 4.5.4 Turbulence in Mid-Infrared Line Emission ...... 94 4.5.5 Comparison to Other Studies of SNR/MC Interactions ...... 97 4.6 Conclusion ...... 99

CHAPTER 5 A Systematic Survey for Broadened CO Lines Toward Galactic Super- nova Remnants ...... 101 5.1 Introduction ...... 101 5.2 Source Selection and Distances ...... 103 5.2.1 Distances ...... 106 5.3 Observations ...... 107 5.4 Algorithm for Selecting BML Regions from Observations ...... 109 5.4.1 Data ...... 109 5.4.2 Summary of BML Region Algorithm ...... 109 5.4.3 Criteria for Possible Interactions ...... 111 5.4.4 BML Region Identiﬁcation ...... 113 5.4.5 BML Region Examples: 3C 391 and Cas A ...... 115 5.4.6 Comparison to Other BML Region Indicators ...... 118 5.5 Results ...... 120 8

TABLE OF CONTENTS – Continued

5.5.1 SNRs with Previous Detections Supporting SNR-MC Interactions 121 5.5.2 Non-Detection Toward G54.1+0.3 ...... 126 5.5.3 Newly Identiﬁed SNR-MC Systems ...... 127 5.6 Discussion ...... 141 5.6.1 The SNR-MC Interaction Rate ...... 141 5.6.2 To What Extent Are SNR-MC Interactions Detectable in 12CO J = 2 − 1?...... 142 5.6.3 Suppressing the SNR-MC Interaction Rate ...... 145 5.6.4 Astrophysical Implications of a Low SNR-MC Interaction Rate . 146 5.6.5 Gamma-Ray Emission from SNR-MC Interactions ...... 147 5.7 Conclusion ...... 151

CHAPTER 6 Conclusion ...... 154 6.1 Results ...... 154 6.2 Future Direction ...... 156 6.2.1 Late-time Radio Emission from Supernovae ...... 156 6.2.2 Measuring the Physical Properties of Molecular Clouds Shocked by Supernova Remnants ...... 159

APPENDIX A Mid-Infrared Fine Structure Lines ...... 162 A.1 [S III] 18.71 µm line ...... 162 A.2 [S IV] 10.51 µm line ...... 162 A.3 [Ne II] 12.81 µm line ...... 165 A.4 Remaining emission lines ...... 167 A.5 Cloudy Models of Purely Photoionized Gas ...... 167

APPENDIX B Alternative Methods for CO Line Broadening ...... 172 B.1 Broadening by Radiative Acceleration ...... 172 B.2 Broadening by Gravitational Acceleration ...... 173

APPENDIX C Individual SNR Distance Determinations ...... 174

APPENDIX D Ancillary Data on Previously Detected BML Regions Toward SNRs 179

REFERENCES ...... 182 9

LIST OF FIGURES

2.1 R-band image of PS15si ...... 32 2.2 Photometry of PS15si ...... 33 2.3 Optical Spectra of PS15si and Comparison SNe ...... 35 2.4 Comparison between PS15si Spectroscopy and Chugai et al.(2004) Syn- thetic Spectrum ...... 39 2.5 Spectral Decomposition of PS15si ...... 45

3.1 2013 Ks band image of Cas A ...... 56 3.2 Multi-epoch Ks-band Imaging of Synchrotron-emitting Knots Toward Cas A 57 3.3 Radio and Infrared Imaging of Synchrotron-emitting Knots Toward Cas A 59 3.4 Spectral Energy Distribution of Synchrotron-emitting Knots Toward Cas A 61

4.1 Positions of Mid-Infrared Observations and Broadened Molecular Clouds Toward Cas A ...... 69 4.2 12CO J = 2 − 1 Spectrum Toward Cas A ...... 71 4.3 12CO J = 2 − 1 Channel Maps Toward Cas A ...... 72 4.4 Integrated Intensity of 12CO J = 2 − 1 in Regions of Line Broadening Toward Cas A ...... 73 4.5 Average 12CO J = 2 − 1 Spectra in Regions of Line Broadening Toward CasA ...... 77 4.6 Mid-Infrared Spectra to the North of Cas A ...... 78 4.7 Mid-Infrared Fine-Structure Lines to the North of Cas A ...... 80 4.8 Forbidden O IVEmission Lines from Cas A ...... 81 4.9 Evolution of Forbidden O IVSpectra from Cas A ...... 84 4.10 Convolution of Forbidden O IVEmission to Low-Resolution ...... 88 4.11 Low-Resolution Forbidden O IVmap of Cas A ...... 90

5.1 12CO Emission of 50 SNRs in Broad Line Survey ...... 108 5.2 Integrated 12CO J = 2 − 1 Emission Toward 3C 391 and Spectrum . . . 110 5.3 Integrated 12CO J = 2 − 1 Emission Toward Cas A and Spectrum . . . . 119 5.4 Integrated 12CO J = 2 − 1 Emission Toward G54.1+0.3 and Spectrum . 127 5.5 Integrated 12CO J = 2 − 1 Emission Toward G08.3−0.0 and Spectrum . 129 5.6 Integrated 12CO J = 2 − 1 Emission Toward G09.9−0.8 and Spectrum . 131 5.7 Integrated 12CO J = 2 − 1 Emission Toward G11.2−0.3 and Spectrum . 132 5.8 Integrated 12CO J = 2 − 1 Emission Toward G12.2+0.3 and Spectrum . 133 5.9 Integrated 12CO J = 2 − 1 Emission Toward G18.6−0.2 and Spectrum . 134 5.10 Integrated 12CO J = 2 − 1 Emission Toward G23.6+0.3 and Spectrum . 136 10

LIST OF FIGURES – Continued

5.11 Integrated 12CO J = 2 − 1 Emission Toward 4C−04.71 and Spectrum . . 137 5.12 Integrated 12CO J = 2 − 1 Emission Toward G29.6+0.1 and Spectrum . 139 5.13 Integrated 12CO J = 2 − 1 Emission Toward G32.4+0.1 and Spectrum . 140

6.1 Model Radio Light Curves of Type IIn SNe ...... 158 6.2 Comparison Between 12CO J = 2 − 1 and J = 3 − 2 Emission Toward 3C391...... 161

A.1 Forbidden S III Emission Lines from Cas A ...... 163 A.2 Forbidden S IV Emission Lines from Cas A ...... 164 A.3 Forbidden Ne II Emission Lines from Cas A ...... 166 A.4 Forbidden Ne V Emission Lines from Cas A ...... 168 A.5 Forbidden Ne III Emission Lines from Cas A ...... 169 A.6 Cloudy Synthesized Mid-Infrared Spectrum for Cas A ...... 171 11

LIST OF TABLES

2.1 Optical and UKIRT Photometry of PS15si ...... 31 2.2 Optical Spectroscopy of PS15si ...... 33 2.3 Best-Fit Parameters for Spectral Decomposition of PS15si ...... 41 2.4 RMS of Residuals for Each Template in Figure 2.5...... 44

3.1 Photometry of Synchrotron-Emitting Knots Toward Cas A ...... 60

4.1 Gaussian ﬁt values for CO lines ...... 75 4.2 Gaussian ﬁt values for convolved high-resolution [O IV] spectra . . . . . 89

5.1 Supernova Remnants in Survey ...... 104 5.2 Newly Detected BML Regions Toward SNRs ...... 113 5.3 SNRs with Detected OH Maser Emission ...... 143 5.4 SNRs with Possible TeV Gamma-Ray Emission at 5σ Signiﬁance . . . . 149 12

ABSTRACT

Supernovae (SNe) are energetic explosions that signal the end of a star’s life. These events and the supernova remnants (SNRs) they leave behind play a central role in stellar feedback by adding energy and momentum and metals to the interstellar medium (ISM). Emis- sion associated with these feedback processes, especially atomic and molecular line emission as well as thermal and nonthermal continuum emission is known to be enhanced in regions of high density, such as dense circumstellar matter (CSM) around SNe and molecular clouds (MCs). In this thesis, I begin with a brief overview of the physics of SN shocks in Chapter 1, focusing on a foundation for studying pan-chromatic signatures of interactions between SNe and dense environments. In Chapter 2, I examine an unusual SN with signatures of CSM interaction in the form of narrow lines of hydrogen (Type IIn) and thermal continuum emission. This SN appears to belong to a class of Type Ia SNe that shares spectroscopic features with Type IIn SNe. I discuss the difficulties of decomposing spectra in a regime where interaction occurs between SN ejecta and CSM, potentially confusing the underlying SN type. This is followed by a discussion of rebrightening that occurred at late-time in B and V band photometry of this SN, possibly associated with clumpy or dense CSM at large distances from the progenitor. In Chapter 3, I examine synchrotron emission from Cassiopeia A, observed in the Ks band over multiple epochs. The synchrotron emission is generally diffuse over the remnant, but there is one location in the southwest portion of the remnant where it appears to be enhanced and entrained as knots of emission in the SNR ejecta. I evaluate whether the Ks band knots are dominated by synchrotron emission by comparing them to other infrared and radio imaging that is known to be dominated by synchrotron emission. Concluding that they are likely synchrotron- emitting knots, I measure the magnetic field strength and electron density required for their evolution over the ∼ 10 yr baseline they were observed and find B ≈ 1.3 − 5.8 mG −3 and ne ≈ 1, 000 − 15, 000 cm . The magnetic field strengths appear enhanced beyond 13 values required by the adiabatic strong shock limit, arguing in favor of other forms of magnetic field amplification in the shock. In Chapter 4, I again discuss Cassiopeia A and interaction between the remnant and nearby MCs as seen at mid-infrared and millimeter wavelengths. I report detection of a SNR-MC interaction and analyze its signatures in broadened molecular lines. I extend this analysis in Chapter 5 to a large survey for SNR- MC interactions in the 12CO J = 2 − 1 line. Although broadened 12CO J = 2 − 1 line emission should be detectable toward virtually all SNR-MC interactions, I find relatively few examples; therefore, the number of interactions is low. This result favors mechanisms other than supernova feedback as the basic trigger for star formation. In addition, I find no significant association between TeV gamma-ray sources and MC interactions, contrary to predictions that SNR-MC interfaces are the primary venues for cosmic ray acceleration. I end this dissertation in Chapter 6 with a brief summary of my results and two extensions of this work: examining the late-time radio light curves of CSM-interacting SNe for signatures of radio synchrotron emission and dense or clumpy CSM at large distances from the progenitor and re-observing SNR-MC interactions in 12CO J = 3 − 2 in order to verify the presence of shock-heated molecular gas and perform a census on the densities and temperatures of post-shock molecular gas. 14

CHAPTER 1

Introduction

When a star reaches the end of its life and explodes as a supernova (SN), the interaction between the SN ejecta and surrounding environment triggers a wealth of energetic processes unlike anywhere else in the Universe. These processes are responsible for a large fraction of the nonthermal emission observed across the electromagnetic spectrum, and the supernova remnants (SNRs) left behind can be studied at radio, infrared, optical, X-ray, and gamma-ray wavelengths. Recent theoretical and observational studies emphasize that this emission is enhanced when the surrounding environment is dense (e.g., in a dense circumstellar environment or molecular cloud), which in turn can drive atomic and molecular line emission, high-energy particle acceleration, and unique nonthermal signatures. Indeed, targeted searches for narrow Hα emission, radio and near-infrared synchrotron emission, broadened molecular lines, and gamma-ray emission have all been used to search for strong shocks around SNe at various stages in their evolution (from days to tens of thousands of years old). These emission signatures all suggest a correlation between SN shocks and driving turbulence, mixing, and particle acceleration in their environments. Despite the success of the diffusive shock acceleration model (in which charged particles are accelerated across a shock in the presence of a magnetic ﬁeld) in explaining the mechanism that accelerates the electrons and protons that are thought to emit this nonthermal radiation, the physical conditions under which this process occurs are still largely uncertain (e.g., Slane et al., 2015). At what rate does particle acceleration occur? What fraction of accelerated particles are protons (as opposed to electrons)? What is the maximum energy accelerated particles can obtain? These questions are all directly related to the size, clumpiness, density, and magnetic ﬁeld strength of the dense medium responsible for particle acceleration and the mechanisms that cause relativistic particles to emit across the electromagnetic spectrum. Measurements of these quantities are critical in working models of SN-driven particle acceleration. 15

In this chapter, I briefly discuss the physics of shocks and the equations that describe physical parameters associated with SN and SNR shocks (Section 1.1). I then discuss the hydrodynamics of SNe as they evolve and drive shocks and the utility and difficulties in applying hydrodynamical models to the environments of SNe (Section 1.2). I briefly summarize the mechanisms by which we can identify SN and SNR shocks electromagnetically, especially in line emission and nonthermal radiation (Section 1.3) and go on to focus on three physical scenarios involving SN shocks and observational techniques used to probe these scenarios: the first several hundred days after a SN explodes into a dense circumstellar environment (Section 1.3.1), particle acceleration associated with nonthermal emission from young SNRs (Section 1.3.2), and the interaction between a SNR and molecular cloud (Section 1.3.3). I conclude this chapter with a brief overview of four observational programs I conducted in order to probe these physical scenarios and the physics of SN shocks generally (Section 1.4).

1.1 The Physics of Supernova Shocks q γkB T Sound propagates through fluids as pressure waves at the local sound speed, cs = , µmH where kB is the Boltzmann constant, T is temperature, µ is the mean molecular weight, and γ is the heat capacity ratio locally in the fluid. In general, the constituents of the fluid will have time to “get out of the way” as the longitudinal compression of the sound wave reaches different parts of the fluid, meaning that the fluid is incompressible. As disturbances propagate through fluid and reach the local sound speed, pressure waves due to the disturbance will not have time to propagate before other sound waves begin to build up. Above the local sound speed, the cascade of sound waves will compress the fluid such that sharp changes in density, pressure, and temperature occur on either side of the compression. This phenomenon is called a “shock” and is observed in many astrophysical systems. Given that they involve compression of a fluid, changes in the fluid due to heating, acceleration, hydrodynamic instabilities, and dissipation of energy are commonly associated with shocks. These changes in the physical parameters from one side of a shock to the other (or the pre-shock to post-shock region for a propagating shock wave) can involve 16

extremely sharp discontinuities. Although they occur under extreme conditions and are often described as a mathemat- ical discontinuity in certain physical parameters, shocks satisfy and are often described using typical conservation relations. The most fundamental of these relations is derived using the one-dimensional Euler equations describing conservation of mass, momentum, and energy in terms of density (ρ), velocity (v), thermal pressure (p), and magnetic ﬁeld (B) in the frame of the discontinuity (i.e., where δ/δt = 0, δ/δy = 0, δ/δz = 0, as

in Kennel et al., 1989; Somov, 2012; Anderl et al., 2013). Assuming vk,1 = vk,2 = 0 and Bk,1 = Bk,2 = 0 where I use 1 and 2 to indicate the upstream and downstream (i.e., pre-shock and post-shock) values of each physical quantity and ⊥ and k to indicate the components of vector quantities normal and tangential to the shock discontinuity, the four jump conditions that describe magneto-hydrodynamical properties across the shock are,

ρv⊥ = 0 (1.1) 2 ρv⊥ + p = 0 (1.2) γ 1 ρv3 + v p + v B2 = 0 (1.3) ⊥ γ − 1 ⊥ 8π ⊥ ⊥ v⊥B⊥ = 0. (1.4)

Here A indicates the difference between the upstream and downstream values of A.

Henceforth, I will refer to B⊥ as B and v⊥ as v while acknowledging the assumptions made above. These equations immediately lead to useful relations between physical parameters on either side of the shock, such as ρ2 = B2 = v1 . Given the simple relation between these ρ1 B1 v2 quantities, the previous set of equations is typically solved for a unitless parameter called the “compression ratio”, that is, r ≡ ρ2 . The compression ratio in this model is obtained ρ1 (as in Draine, 2011; Anderl et al., 2013) assuming the Mach number M = v such that cs

2(γ + 1) r = , where (1.5) q 2 −2 D + D + 4(γ + 1)(2 − γ)MA D ≡ (γ − 1) + 2M −2 + γM −2, and (1.6) √ 1 A v1 4πρ1 MA ≡ . (1.7) B1 17

These equations yield dependencies between the pre- and post-shock conditions in the

strong shock limit where M1 1,MA 1. In this case and assuming γ = 5/3 (monoatomic gas), then D = 2/3 and r = 4. The non-magnetic case is similar to this −1 γ+1 −2 γ+1 treatment where MA → 0 such that r = D and D = (γ − 1) + 2M1 , or r = γ−1 in the strong shock limit. From these relations, we can apply the physics of SN shocks to a wide range of scenarios with differing physical conditions. Below, I discuss the evolution of SN shocks and apply the prescription outlined above to several different regimes.

1.2 Dynamical Evolution of Supernova Shocks

As a SNR expands, the size of the remnant and the shock speed are often described in differing regimes with certain simplifying assumptions. Two of these are the velocity- conserving free-expansion phase and the adiabatic “snowplow” phase. In the following discussion, I derive equations that describe each phase and discuss their application to open problems examining the evolution of SN and SNR shocks in dense environments. In particular, the treatment of these phases sometimes involves a number of assumptions about the geometry of the SN and surrounding medium, which can obscure their application to observations of SN-shocked environments.

1.2.1 Free-Expansion Phase

The total kinetic energy in a SN explosion is a critical parameter that yields information about the progenitor system and evolution of the SNR. For thermonuclear explosion mod- 56 els, a white dwarf produces 1.4 M of ejecta and 0.4 − 0.8 M of Ni with a kinetic energy of ∼ 1051 ergs (Hoyle and Fowler, 1960; Arnett, 1968; Nomoto, 1986; Nady- ozhin and Imshennik, 2005; Wang et al., 2008). In one of the biggest coincidences in astrophysics, measurements of core-collapse SNe such as SN 1987A (e.g., Imshennik and

Nadezhin, 1989; Bethe and Pizzochero, 1990) suggest that the SN produces about 2 M of ejecta and also carries a kinetic energy of roughly 1051 erg. There may be a wide variety of core-collapse SNe with varying ejecta masses and kinetic energies to match the variety 18

of spectral types (Filippenko, 1997), but these values are fiducial numbers for SNe of all types. These values imply that SN ejecta will typically explode with an initial velocity of q 2Ek −1 v0 = = 10, 000 km s . As sound speeds in most astrophysical systems are Mejecta 100 km s−1 or less, the initial ejecta velocities are clearly supersonic. For example, the −1 sound speed in molecular clouds with T = 20 K, γ = 5/3µ = 2 is cs = 0.3 km s and 5 −1 the sound speed in H II regions with T = 10 K, γ = 5/3, µ = 0.62 is cs = 47 km s . However, we have implicitly assumed that the kinetic energy is evenly distributed throughout all of the ejecta. This assumption is typically justified by arguing that SN explosions should follow spherical geometry and pointing to the general spherical (or cir- cular, in projection) shape of young SNRs (e.g., SN 1006, Cassiopeia A, Tycho, Kepler, 0509-67.5, 0519-69.0 van den Bergh, 1976; Patnaude and Fesen, 2014; Williams et al., 2016; Burkey et al., 2013; Hovey et al., 2015; Kosenko et al., 2010). But many of these systems are clearly somewhat asymmetric, such as along the northeast to southwest bipolar axis of Cassiopeia A (Hwang et al., 2004; Grefenstette et al., 2014) or simply in polarimetry and models of core-collapse SNe (Leonard et al., 2005; Couch and Ott, 2013). On average, however, optical spectroscopy confirms that 10, 000 km s−1 is a good approximation for the expansion velocity for ejecta in young SNe (Fesen et al., 1988; Foley and Kasen, 2011; Foley, 2012). This velocity is largely (though not entirely) conserved until the ejecta are significantly decelerated by material swept up by the expanding SN, which is why this phase is referred to as “free-expansion.” Cooling due to expansion and radiative losses is also relatively insigificant. In spite of the fact that SNe are rapidly expanding and among the most luminous observed objects, energy from the decay of radioactive isotopes such as 56Ni keeps the ejecta hot and produces a characteristic light curve as the SN fades (Arnett, 1982; Arnett et al., 1985). Eventually, the SN becomes optically thin and is no longer able to efficiently absorb radiation from these isotopes. The physical regime in which radiative cooling dominates the evolution of the remnant occurs much later than the free-expansion phase. Although this description matches the majority of SNe, Type IIn (IIn for narrow lines of hydrogen in their spectra) SNe break the trends described in the free-expansion phase. 19

I discuss these objects in detail below in Section 1.3.1.

1.2.2 Snowplow Phase

Although SN ejecta carry significant kinetic energy and generally stream freely away from the point of explosion for decades or centuries, the medium in which the SN explodes will significantly decelerate the ejecta. This deceleration is analogous to a hot bubble expanding adiabatically in a low temperature fluid or a snowplow sweeping up material; eventually, the mass of the swept up material will contribute significantly to the overall mass of the expanding remnant even as momentum and internal energy are conserved. From this observation, it is reasonable to infer that deceleration will become significant when the mass of the swept up material is equal to the ejecta mass. Analytically, we can approximate how long the free-expansion model will be a good approximation of SN evolution. For a SN expanding into a homogeneous medium with 1/3 3Mejecta density n and mass density ρ = µmH n, the SNR will have a size RST = 4πρ when it has swept up material with mass equal to the ejecta mass, which will take tST ≈ RST . Assuming the same ejecta mass and initial velocity from the free-expansion phase v0 and that the environment surrounding the SN is homogeneous with a density of about −3 0.5 cm (for an H I cloud with µ ≈ 1.3), then RST ≈ 3.1 pc and tST ≈ 300 yr. The latter assumption in this calculation is clearly incorrect. The interstellar medium is not homogeneous but consists of clumps of high density atomic and molecular gas and low density H II regions created by OB associations and other SNRs. While young SNRs often bear some imprint of spherical symmetry, many older SNRs are not spherical (e.g., Cygnus Loop, 3C391, IC 443, W44 Green, 1990; Becker and Kundu, 1975; Erickson and Mahoney, 1985; Kundu and Velusamy, 1972) while some SNRs retain spherical symmetry beyond the 300 yr horizon typically assumed for the Sedov-Taylor phase (e.g., SN 1006 van den Bergh, 1976). Many SNRs bear evidence of significant deceleration in regions where they have impacted high density material, such as a molecular cloud, with “blow outs” of high velocity ejecta in regions where the surrounding medium has a much lower density. Indeed, the interaction between SNRs and their interstellar environments plays a significant role in shaping both. Hydrodynamical simulations (e.g., FIRE Hopkins et al., 20

2014) of heterogeneous interstellar matter will yield a more accurate picture of the evolution of SNRs and deposition of energy and momentum as they expand. Nevertheless, the Sedov-Taylor solution (discovered independently by L. I. Sedov and G. I. Taylor; Sedov, 1946; Taylor, 1950) to the expansion of the SNR shock in a homogeneous medium provides insight as to the evolution of the remnant shock velocity and temperature over time. Radiative losses are relatively insigniﬁcant during this phase, and so it is reasonable to assume energy conservation in the mass of ejecta and swept up material. This phase is usually described by a self-similarity solution in which R ∼ Eχρψtω. Di- mensional analysis indicates that the solution to this relation must be χ = 1/5, ρ = −1/5, ω = 2/5. A similar result can be obtained by explicitly assuming energy conservation and

Z R 2 ! 1 2 dr Ek = 4πr ρ dr , that is (1.8) 2 0 dt 2π ρR5 E ≈ (1.9) k 5 t2 E 1/5 n −1/5 t 2/5 R ≈ 3.1 pc k , and (1.10) 1051 erg 0.5 cm−3 300 yr E 1/5 n −1/5 t −3/5 v ≈ 10, 000 km s−1 k (1.11) 1051 erg 0.5 cm−3 300 yr where I have assumed a homogeneous density profile (i.e., ρ(r) = constant). This relation yields several interesting results when combined with the jump conditions in Section 1.1. For example, in the strong shock limit and assuming γ = 5/3, it can be shown that the compression ratio is r = 4 and the post-shock medium will be slowed to 1/4 of the velocity of the pre-shock gas in the shock frame. In the non-magnetic case, the jump conditions yield a post-shock temperature T2 = 2 p2 ρ1 2(γ−1) µ2mH 2 3 µ2mH v1 T1 = 2 v1 = (for γ = 5/3). Combining this rela- p1 ρ2 (γ+1) kB 16 kB tion with the above equation for velocity, the post-shock temperature should be of order 2/5 −2/5 −6/5 T = 3 × 108 K Ek n t . 2 1051 erg 0.5 cm−3 1000 yr This result provides a number of useful insights about the nature of the post-shock gas. Thermal X-ray emission should be a dominant component of the electromagnetic signature, the distribution of particle momenta in a thermal (Maxwellian) particle distri- 21 bution will facilitate significant particle acceleration (i.e., via diffusive shock acceleration Kang et al., 2002), and the magnetic field strength corresponding to an Alfven wave with √ q 4πγkB T velocity equal to γkBT µmH (i.e., B = n ) is about B = 0.5 mG. I discuss more implications of this result in the context of nonthermal emission and interactions with molecular clouds in Section 1.3.2 and Section 1.3.3.

1.3 Electromagnetic Signatures of Shocks From Supernovae and Supernova Remnants

Current observational studies probe strong shocks from SNe and SNRs across a wide range of the electromagnetic spectrum, from gamma-ray signatures of highly relativistic particles to thermal and nonthermal X-ray emission to line emission in optical and infrared wavelengths and dust emission in the mid-infrared. Below, I discuss three physical scenarios and observational signatures associated with them, including CSM interactions around SNe, nonthermal emission from young SNRs, and interactions between SNRs and molecular clouds (SNR-MC interactions).

1.3.1 Supernovae with Circumstellar Interactions

Stars of all masses exhibit some form of mass-loss through winds, ﬂares, ejections, eruptions, or mass transfer. These processes all contribute to pollution of the immediate environments of stars with circumstellar matter (CSM). When stars explode as SNe, they shock and interact with the material in the circumstellar medium producing radiation that can be observed in the radio (e.g., SN 1988Z Williams et al., 2002), optical (e.g., SN 2006gy Ofek et al., 2007), and X-ray (e.g., SN 2010jl Ofek et al., 2014; Chandra et al., 2015). One of the most ubiquitous signatures of strong CSM interactions is narrow lines of hydrogen in SN spectra (Type IIn SN or SN IIn). SNe IIn are usually interpreted as evidence for a mass of slow-moving material from a pre-SN wind or other mass-loss event. Fast-moving ejecta will interact with this material creating a shock, but the ejecta are decelerated as a result. The shock in turn heats the CSM, which is optically thick and generates blackbody continuum emission in addition to narrow lines of hydrogen. An interesting consequence of the interaction between the SN and CSM is that the total 22

amount of CSM in the immediate environments of these events can easily equal the ejecta mass causing the early onset of the snowplow phase. As a result, SNe IIn can efﬁciently convert the kinetic energy into radiation. For example, the equations derived above for the post-shock temperature in the strong shock approximation of the Sedov-Taylor solution can be used to derive the luminosity due to Hα emission observed in a SN IIn (as in, e.g., Chevalier and Fransson, 1994)

2 2 3 µ2mH v1 9 v1 T2 = = 3 × 10 K −1 (1.12) 16 kB 10, 000 km s −1 37 −1 T2 L(Hα) = j(Hα)/a = 4.4 × 10 erg s L40 (1.13) 107 K

where j(Hα) is the Hα emissivity, aB is the recombination coefﬁcient of hydrogen, and 40 −1 L40 is the ionizing luminosity in the shocked region in units of 10 erg s . Similar approximations can be made for the total amount of radio emission (see Section 6.2.1) and X-ray emission. Another consequence of SN and CSM interactions and the emission arising from the interaction region is that the intrinsic SN emission is obscured by continuum emission from the interaction region. In optical spectra, absorption and emission features lose contrast relative to the continuum level and are obscured by emission from the shock interaction (e.g., in SN 2006gy Smith et al., 2007a). How can we interpret the underlying SN spectrum when optical spectra must be decomposed into emission from the SN and interaction region? This task becomes non-trivial when one considers that feedback from ultraviolet/X-ray radiation in the interaction radiation may change the ionization state of the outermost ejecta layers where visible emission from the SN originates.

1.3.2 Nonthermal Emission from Young Supernova Remnants

SNRs are thought to be the primary source of cosmic rays below the “knee” of the cosmic ray spectrum at ∼ 1 PeV. In the shells of young SNRs such as Cassiopeia A, ﬁlaments observable in X-ray and radio wavelengths exhibit distinct nonthermal emission consistent with emission processes from shock accelerated electrons (Allen et al., 1997; Vink 23 and Laming, 2003a). These ﬁlaments have been detected in Chandra ACIS observations from 4.2–6 keV where the nonthermal X-ray emission appears to decline at 1.5% per year (Patnaude et al., 2011). Recent observations of gamma-rays from hadronic particle acceleration in SNRs (Aharonian et al., 2006a,b,c; Acciari et al., 2009; Abdo et al., 2010) have also contributed to the study of cosmic ray production. The physics of this nonthermal emission from high energy particles is also well-documented at near-infrared wavelengths

(Gerardy and Fesen, 2001; Rho et al., 2003) where Ks-band emission closely matches the morphology of nonthermal emission at other wavelengths that has been demonstrated to exhibit a typical synchrotron power-law spectrum. The mechanism by which particles are accelerated in SNRs is called diffusive shock acceleration and derives from first-order Fermi acceleration (Fermi, 1949; Bell, 1978) in which particles gain energy linearly with the ratio of the velocity of the accelerated particle to the velocity of post-shock gas (as opposed to quadratically in second-order acceleration). Although particle acceleration occurs throughout the lifetime of a SNR, particles are accelerated most efficiently when the shock speed is high, which decreases the time needed for particles to be accelerated up to TeV and PeV energies. Thus, young SNRs such as Cassiopeia A and Tycho are strong X-ray and radio sources; indeed, the former is the brightest radio source in the sky. While the diffusive shock acceleration model accurately describes the radio synchrotron spectrum in SNRs, there is still ambiguity as to the source of nonthermal X-ray emission in SNRs. Predictions of the magnetic field from nonthermal X-ray filaments and assuming a synchrotron spectrum disagree with those predicted from diffusive shock acceleration (Vink, 2012), and other mechanisms are known to produce radiation in SNRs in the 10–100 keV range such as inverse Compton scattering (Porter et al., 2006) and nonthermal bremsstrahlung emission (Castro et al., 2013). The confusion between various sources of emission at these energies has generated new problems as high energy astrophysics extends further into the GeV and TeV gamma-ray observations. Well-understood emission processes in the radio cannot be linked to new observations in gamma-ray astronomy without some bridge to connect these regimes. Given the rapid expansion of SN shock fronts into heterogeneous interstellar medium, 24

it is reasonable to expect corresponding variations in synchrotron emission generated from in situ particle acceleration near the shock front. Temporal variation in the synchrotron emission at these locations can help to constrain the density of the medium into which the SN is expanding as well as the physics of electron acceleration at the sites of strong

Ks-band emission. Thus, when the SN ejecta shock regions of high-density such as a molecular cloud, one might expect a signiﬁcant increase in synchrotron emission. Can

the physical parameters associated with Ks emission be accurately linked to synchrotron emission and particle acceleration as observed at other wavelengths? What can studies of this emission tell us about the physical conditions in SNR shocks?

1.3.3 Interactions Between Supernova Remnants and Molecular Clouds

Expanding SNRs play a central role in the turbulent acceleration of interstellar gas, which in turn regulates important processes in the interstellar medium, for example, in star formation where turbulent molecular gas can delay the collapse of star forming cloud cores. For a SNR in the Sedov-Taylor phase of expansion, the total momentum deposited into a cloud can be calculated from the total mass and velocity in the shock, that is,

4/5 1/5 −1 Ek n Mv ≈ 20, 000 M km s (1.14) 1051 erg 0.5 cm−3 This value is invariant under the assumptions above due to the conservation of momentum implicit in the Sedov-Taylor approximation. Suppose a SNR with size of R ∼ 3 pc encounters a molecular cloud with comparable size and density ∼ 1000 cm−3. If the SNR transfers all of its momentum to the cloud as turbulence, then the cloud will have an aver- −1 age turbulent velocity of vturb ≈ 4 km s . This value is comparable to what is found for SNR-MC interactions where molecular clouds are turbulently disrupted by SNR shocks (e.g., Ziurys et al., 1989; van Dishoeck et al., 1993). Indeed, in extreme cases where a SNR shocks and transfers a large part of its momentum to a small cloud core, turbulent velocities in molecular clouds can reach > 30 km s−1. This physical scenario is an important one for feedback and energetic processes in the interstellar medium. A correlation is predicted between the power spectrum of turbulent 25 velocity and the magnetic fields observed in molecular clouds, which suggests some connection between these two mechanisms (McKee and Ostriker, 2007). From some seed magnetic field, dynamo action driven by turbulence from SNe may be largely responsible for the amplified magnetic field observed in the galaxy (Ferriere, 2003). The locations of Galactic SNRs and the amount of turbulent kinetic energy they inject into interstellar gas are important parameters to understand the origin of Galactic magnetic fields. From this association between Galactic magnetic fields and the engines of particle acceleration and magnetic field amplification, it is reasonable to predict that SNRs may be one of the dominant sources of cosmic rays (Fermi, 1949; Ginzburg and Syrovatskii, 1964; Zinnecker and Yorke, 2007; Cardillo et al., 2014). This complex system of cosmic rays, magnetic fields, and SN-driven turbulence affects virtually every process in the interstellar medium. Enrichment from remnants also plays a critical role in the chemical evolution of the interstellar medium, especially where dust condenses from outflowing SN ejecta rich in iron, magnesium, silicon, and oxygen (Bianchi and Schneider, 2007; Rho et al., 2008). The boundary between a SN shock and the interstellar medium is thus a pivotal point in the evolution of the interstellar medium. Emission from shocked molecular gas yields insight to the morphology of young SNRs, the energy and velocity of their ejecta, and the efficiency of feedback and enrichment mechanisms. Shocked emission from molecular species is one of the most reliable ways to detect interaction between SNe and the surrounding interstellar medium. In particular for young SNRs the shock velocity can be too high to excite strong molecular hydrogen emission (Burton et al., 1988) or the 1720 MHz OH maser, both of which are often used to trace shocks (Claussen et al., 1997). In such cases, molecular line broadening associated with turbulence in molecular clouds serves as the most useful tracer of the interaction of the SN ejecta with molecular clouds. How can we extrapolate our knowledge of where and how often these interactions occur to processes in the Milky Way Galaxy at large? What does the incidence of SNR-MC interactions in our galaxy tell us about star formation and cosmic ray production? 26

1.4 This Volume

I have described several physical scenarios associated with SN shocks and emission signatures that might be used to constrain the conditions associated with them. In the following chapters, I describe observational programs I conducted in order to probe these emission signatures. Chapter 2 describes a SN with narrow lines of hydrogen which suggest an interaction with dense CSM. Unlike most SNe IIn, optical spectroscopy from this SN indicates that it shares morphological signatures with Type Ia SNe. I attempt to decompose this spectrum in order to probe the SN obscured by emission from the CSM interaction.

In Chapter 3, I use multi-epoch Ks band imaging of the SNR Cassiopeia A to verify that emission at this wavelength is dominated by synchrotron radiation and then use the variation in synchrotron emission to examine the magnetic field strength and electron density in the SNR. I re-examine Cassiopeia A at mid-infrared and millimeter wavelengths in Chapter 4 where I find evidence for interaction between the SNR and nearby molecular clouds. The SNR appears to be interacting with molecular gas well beyond the shock front of the SNR, which I argue arises from high-velocity ejecta along the northeast-southwest bipolar outflow from the SNR. Finally, in Chapter 5, I extend my millimeter observations to a survey involving 50 Galactic SNRs. I find that the fraction of SNRs with signs of molecular cloud interactions is much lower than expected from the assumption that all remnants from core-collapse SNe should exhibit these interactions, and I use these results to interpret predictions related to the effect SNRs have on sequential star formation and particle acceleration associated with TeV gamma-rays. Finally, I summarize my findings in Chapter 6 and conclude this volume with a description of ongoing projects that extend my work to radio observations of SNe and sub-millimeter observations of interactions between SNRs and molecular clouds. 27

CHAPTER 2

PS15si: A Type Ia/IIn Supernova with Late-time Rebrightening

We present optical/near-infrared spectroscopy and photometry of the supernova (SN) PS15si. This object was originally identified as a Type IIn SN, but here we argue that it should be reclassified as a Type Ia SN with narrow hydrogen lines originating from interaction with circumstellar matter (CSM; i.e., SN Ia/IIn or SN Ia-CSM). Based on deep nondetections 27 days before discovery, we infer that this SN was discovered around or slightly before optical maximum, and we estimate the approximate time that it reached R-band maximum based on comparison with other SNe Ia/IIn. In terms of spectral morphology, we find that PS15si can be matched to a range of SN Ia spectral types, although SN 1991T-like SNe Ia provides the most self-consistent match. While this observation agrees with analysis of most other SNe Ia/IIn, we find that the implied CSM luminosity is too low to account for the overall luminosity of the SN at a time when the CSM should outshine the underlying SN by a few magnitudes. We infer that the similarity between PS15si and the hot, overluminous, high-ionisation spectrum of SN 1991T is a consequence of a spectrum that originates in ejecta layers that are heated by ultraviolet/X-ray radiation from CSM interaction. In addition, PS15si may have rebrightened over a short timescale in the B and V bands around 85 days after discovery, perhaps indicating that the SN ejecta are interacting with a local enhancement in CSM produced by clumps or a shell at large radii.

2.1 Introduction

The increasing number of targets from high-cadence surveys have revealed enormous variety in spectroscopic and photometric signatures of supernovae (SNe) including the subclass of Type Ia SNe (SNe Ia). While some variations in SN Ia spectra and light curves have long been recognised, such as the width-luminosity or Phillips relation (Phillips, 1993), recent focus on ejecta velocities (Wang et al., 2009), spectral variation with time 28

(Patat et al., 2007), and polarization signals (Kasen et al., 2003; Patat et al., 2009; Porter et al., 2016) suggest a wide variety of SN Ia subtypes. Peculiar trends in these measurements contribute to the ambiguity in the underlying mechanism(s) for SNe Ia. Despite the success of the C/O white dwarf (WD) thermonuclear explosion model in explaining the explosion mechanism of SNe Ia (Hoyle and Fowler, 1960; Arnett, 1968; Nomoto, 1986), it is still unclear whether a dominant evolutionary channel exists to ignite these explosions. Two models predict different channels through which ignition might occur — a single-degenerate model, involving a WD that accretes from a nondegenerate companion star, and a double-degenerate model, involving the merger of two WDs. Various lines of evidence from spectroscopic and photometric signatures of various SNe Ia sub-types seem to argue in favor of both models. For example, it has been suggested that helium in peculiar SNe Iax (Foley et al., 2013) originates from nondegenerate companion stars, and signatures of circumstellar matter (CSM) in thermal X-ray emission and spectropolarimetry in some SNe Ia also point toward mass loss from a nondegenerate companion (Wang et al., 1996; Immler et al., 2006; Patat et al., 2012). At the same time, so-called super-Chandrasekhar SNe exhibit high luminosity and low ejecta velocities, and point toward massive (∼ 2 M ) WD progenitors, possibly from WD mergers (Howell et al., 2006; Hicken et al., 2007; Silverman et al., 2011). While illuminating the diversity of SNe Ia, these systems generate additional questions. Do they represent the extremes along a continuum of SN Ia explosion scenarios? What similarities do these SNe Ia share with each other and the more common classifications? In particular, one new class of SN Ia-like events exhibits the spectroscopic signatures of both SNe Ia and SNe IIn (IIn for narrow lines of hydrogen) in the form of broad Fe, Ca, S, and Si absorption combined with strong, narrow Hα emission, consistent with a SN Ia explosion encountering dense CSM. The dominant hydrogen feature seen in this subclass contrasts with the majority of SNe Ia where Hα searches have yielded null results to deep limits of ∼ 0.001 − 0.01 M (e.g., SNe 2005am, 2005cf, and 2011fe as in Leonard, 2007; Shappee et al., 2013). The first SN identified with both SN Ia and SN IIn features, SN 2002ic (Hamuy et al., 2003), exhibited absorption features characteristic of SN 1991T-like SNe Ia (Filippenko et al., 1992; Phillips et al., 1992; Filippenko, 1997), but 29

was identified as having both broad and narrow Hα profiles as early as 6 days after maximum light. SN 2002ic and SNe with similar spectroscopic signatures have been denoted as SNe Ia/IIn, IIa, and sometimes Ia-CSM1 (Deng et al., 2004; Kotak et al., 2004; Silverman et al., 2013a) owing to their spectroscopic overlap with both SNe Ia and IIn. Subsequent to the discovery of SN 2002ic, two SNe IIn — SNe 1997cy and 1999E (Hamuy et al., 2003; Wood-Vasey et al., 2004) — were reclassified as SN 2002ic-like. More recently, several SNe Ia/IIn, such as SN 2005gj and PTF11kx (Aldering et al., 2006; Prieto et al., 2007; Dilday et al., 2012), were spectroscopically identified soon after explosion. It is now estimated that this class represents as many as 0.1 − 1% of all SNe Ia (Dilday et al., 2012). In addition to their unique spectroscopic properties, the light curves of SNe Ia/IIn also exhibit traits that overlap with both SNe Ia and SNe IIn. SNe Ia/IIn consistently have R-band peak absolute magnitudes brighter than −19 mag, and their light curves evolve slowly with a linear decay that can last for several weeks. For the first 25 days after peak, SNe Ia/IIn light curves are generally consistent with “normal” SNe Ia (Hamuy et al., 2003; Prieto et al., 2007). After this point, SNe Ia/IIn decline much more slowly than SNe Ia, and the residual flux is assumed to be due to CSM interaction. In this model, the freely expanding, 56Ni-powered SN can be observed mostly at early times and the optically thick CSM interaction is dominant at later times (Chugai and Yungelson, 2004), perhaps followed by a return to the 56Ni-powered decay line. While the general trend of slow decay is observed, it is still an open question why some SNe Ia/IIn exhibit strong SN Ia-like absorption in early-time spectra (e.g., Hamuy et al., 2003; Silverman et al., 2013a) while others appear to be diluted by a thermal continuum (e.g., Aldering et al., 2006). During the period when Hα emission fades (generally ∼70– 300 days), the spectrum further evolves as the thermal continuum becomes less apparent and the SN Ia-like component enters the nebular phase. This epoch in SN Ia/IIn evolution is critical for determining the characteristics of both the CSM and the intrinsic SN emission where competing models of the progenitor system can be evaluated (Chugai et al., 2004). This phase also probes the CSM out to large radii where observations may reveal complex

1Hereafter, we refer to this class of objects as SNe Ia/IIn, since the name SN Ia-CSM mixes an interpretation of a physical mechanism with a spectral classification. 30 structure in the surrounding medium. In this paper we discuss PS15si, discovered by the Pan-STARRS Survey for Transients (Smith et al., 2015a) on 2015 Mar. 23 (all dates presented herein are UT) and located 100.6 from the center of the galaxy 2dFGRS N166Z116 (hereafter N166Z116). The target was originally classified 5 days after discovery as a SN IIn, with narrow hydrogen lines indicating redshift z = 0.053 (Walton et al., 2015). In this paper, we argue that PS15si is a SN Ia/IIn. We adopt a Milky Way line-of-sight reddening of E(B − V ) = 0.046 mag (Schlafly and Finkbeiner, 2011), a distance to N166Z116 of 219.4 ± 15.4 Mpc, and m − M = 36.71 ± 0.15 mag (Colless et al., 2003). From the light curve and prediscovery images2, we infer that the SN was discovered around or slightly before optical maximum given the discovery magnitude w = 17.16 (Sloan g + r + i) on 2015 Mar. 23 (see Huber et al., 2015). Assuming a line-of-sight Aw = 0.125 mag, the peak apparent magnitude indicates PS15si had an absolute w-band magnitude of Mw ≈ −19.7 mag.

2.2 Observations

We took optical imaging photometry of PS15si using the Super-LOTIS (Livermore Optical Transient Imaging System; Williams et al., 2008) 0.6 m telescope at Kitt Peak National Observatory between 2015 May 8 and Jun. 20. Standard reductions, including flat-fielding and bias subtraction, were carried out using a semi-automatic routine. We then performed differential photometry using stars in the field and APASS standards. A Super-LOTIS image including the SN from 2015 May 21 is shown in Figure 2.1. We transformed R and I-band standards from APASS r0 and i0 magnitudes as described by Jester et al.(2005). We took near-infrared (NIR) observations with the 3.8 m United Kingdom Infrared Telescope (UKIRT) on Mauna Kea using WFCAM2 (2015 May 3, Jun. 21, Jul. 5)3. JHK observations were pipeline reduced by the Cambridge Astronomical Survey Unit (CASU). We performed aperture photometry on all images using a 100.3 diameter aperture to reduce contamination from the host galaxy, and we calibrated the instrumental magni-

2The target was provided from the Pan-STARRS NEO survey and made public via star.pst.qub. ac.uk/ps1threepi/. 3Our observing campaign for PS15si ended at this point due to solar conjunction. 31

Table 2.1: Optical and UKIRT Photometry of PS15si

UT Date day4 BVRI (y-m-d) 2015-05-08 46 17.62±0.14 17.04±0.14 17.31±0.04 16.58±0.06 2015-05-10 48 17.78±0.12 17.09±0.13 17.37±0.04 16.62±0.06 2015-05-19 57 ... 17.23±0.12 17.50±0.04 16.83±0.08 2015-05-21 59 ... 17.38±0.13 17.51±0.03 16.82±0.07 2015-05-31 69 ... 17.28±0.17 17.73±0.07 16.82±0.07 2015-06-14 83 ...... 17.79±0.04 17.17±0.08 2015-06-15 84 18.16±0.16 17.62±0.12 17.75±0.04 ... 2015-06-16 85 18.32±0.18 17.63±0.14 17.83±0.04 ... 2015-06-17 86 ... 17.51±0.12 17.83±0.04 ... 2015-06-20 89 18.11±0.19 17.31±0.12 ...... UT Date day JHKs (y-m-d) 2015-05-03 41 16.65±0.06 16.14±0.08 16.08±0.06 2015-06-21 90 17.68±0.09 17.16±0.09 16.72±0.07 2015-07-05 104 17.97±0.07 17.40±0.08 17.36±0.08

tudes using 2MASS JHKs NIR standard stars present in the field. For both optical and NIR magnitudes, we calculated uncertainties by adding in quadrature photon statistics and zero-point deviation of the standard stars for each epoch. We detail the optical and NIR magnitudes for PS15si in Table 2.1 and Figure 2.2. We obtained 2 epochs of moderate-resolution optical spectra with the Bluechannel spectrograph on the Multiple Mirror Telescope (MMT) on 2015 Apr. 30 and 2015 Jun. 12. Each MMT Bluechannel observation was taken with a 1.000 slit and the 1200 l mm−1 grating with a central wavelength of 6350 A˚ and 3 × 1200 s exposures. The spectral range we used covers approximately 5800–7000 A.˚ Standard reductions were carried out using IRAF5 including bias subtraction, flat fielding, and optimal extraction of the spectra. Flux calibration was achieved using spectrophotometric standards observed at an airmass similar to that of each science frame, and the resulting spectra were median combined into a single 1D spectrum for each epoch. In addition, we retrieved the spectrum used for the original spectroscopic identification of PS15si by the Public ESO Spectroscopic Survey of Transient Objects (PESSTO; see

4Since discovery on 2015 Mar. 23. 5IRAF, the Image Reduction and Analysis Facility, is distributed by the National Optical Astronomy Observatory, which is operated by the Association of Universities for Research in Astronomy (AURA) under cooperative agreement with the National Science Foundation (NSF). 32

Figure 2.1: R-band image of PS15si obtained on 2015 May 21 by Super-LOTIS. The position of the SN is indicated along with stars used to perform differential photometry (as described in Section 2.2). The SN was at R = 17.51 mag at this time. 33

Figure 2.2: (Left) R-band photometry of PS15si is plotted as crosses along with the discovery magnitude (in the w band; i.e., g + r + i). We show the R-band light curve of a SN 1991T-like (SN Ia) template in red derived from models in Stern et al.(2004). We also plot R-band photometry of the known SN Ia/IIn 2005gj in blue, which has shifted, stretched by a factor of 1.2 and dimmed by 0.56 mag for comparison to the R-band decline of PS15si. The phase of the SN 2005gj and SN 1991T-like light curves (i.e., in time since explosion) are matched to each other. (Right) Our BVRIJHKs photometry of PS15si.

Table 2.2: Optical Spectroscopy of PS15si

UT Date day7 Telescope/Instrument Slit Width Res. Range λ ˚ (y-m-d) (arcsec) ∆λ (A) 2015-03-28 5 ESO-NTT/EFOSC2 1.0 355 3640–9235 2015-05-01 39 MMT/Bluechannel 1.0 4500 5727–7012 2015-06-13 82 MMT/Bluechannel 1.0 4500 5727–7012 2015-06-16 85 Keck/LRIS 1.0 600 3200–10200

Smartt et al., 2013; Valenti et al., 2014) from WISeREP6 (Yaron and Gal-Yam, 2012); it had been obtained with the Faint Object Spectrograph and Camera (EFOSC2) on the European Southern Observatory’s New Technology Telescope (ESO-NTT) on 2015 Mar. 28 (Walton et al., 2015). The slit width, observed spectral range, and resolution are given in Table 2.2. We obtained the ﬁnal spectral epoch of PS15si on 2015 Jun. 16 with the Low Resolu- tion Imaging Spectrometer (LRIS; Oke et al., 1995) at the Keck Observatory. We used the 1.000 slit rotated to the parallactic angle to minimise the effects of atmospheric dispersion (Filippenko, 1982, in addition, LRIS has an atmospheric dispersion corrector and the ob-

6wiserep.weizmann.ac.il/ 7Since discovery on 2015 Mar. 23. 34 ject was at low airmass, ∼ 1.3). In our LRIS conﬁguration, coverage in the blue with the 600/4000 grism extends over 3200–5600 A.˚ We used the 5600 A˚ dichroic and our coverage in the red with the 400/8500 grating extends over 5600–10,200 A.˚ We obtained one 200 s exposure, and reduced it using routines written speciﬁcally for LRIS in the Carnegie

PYTHON (CARPY) package. We performed standard reductions on the two-dimensional (2D) images including flat-fielding, correction for distortion along the slit axis, wavelength calibration with arc-lamp spectra, and cosmic ray cleaning before extracting the 1D spectrum of PS15si. We flux calibrated the extracted spectrum using a sensitivity function derived from a standard star obtained the same night in the same instrument configuration. We also used the standard-star spectrum to remove the telluric sky absorption features. The final spectra are shown in Figure 2.3. The resolution at Hα in each spectrum (in chronological order) is roughly 840, 66, 66, and 500 km s−1. These data have been dereddened for Milky Way Galaxy reddening assuming E(B − V ) = 0.046 mag as reported by Schlafly and Finkbeiner(2011). We have also removed the recession velocity v = 15, 747 km s−1 given the observed wavelength of narrow Hα emission, which is consistent with the host-galaxy redshift.

2.3 Results

2.3.1 Photometry

To analyze our photometry of PS15si, we compare our data to light curves of other SNe Ia in the literature, especially those of SNe Ia/IIn after optical maximum. In Figure 2.2 (left), we overplot light curves from an overluminous SN 1991T-like template (Stern et al., 2004) and SN 2005gj as a SN Ia/IIn template (Prieto et al., 2007) along with our R-band observations and the discovery w-band magnitude of PS15si. The light curve of SN 2005gj is ﬁt to our data of PS15si assuming that both light curves follow the same shape, and that PS15si is 0.56 mag fainter than SN 2005gj (i.e., shifted in magnitude without any stretch). This brightness allows us to match the shape of the SN 2005gj light curve to prediscovery constraints detailed in Section 2.4.1. As we demonstrate in Figure 2.2 (right), the R-band magnitude of PS15si appears to 35

Figure 2.3: Four spectral epochs of PS15si with the day relative to discovery (d##) of observation given in black. Overplotted at each epoch is a comparison SN Ia/IIn spectrum in red taken from Silverman et al.(2013a) and corrected for redshift and Galactic reddening using the values in Tables 1 and 2 of that publication. The phase (i.e., d##) for each comparison spectrum is the time in days since R-band maximum for CSS120327 and since discovery for SN 2011jb. Features labeled in grey (Mg II, Fe II) are in absorption while features labeled in blue (Hα,Hβ, He I, [O I], [O III], [Fe II], [Fe III], He I λ7281/[Ca II] λ7324 blend, Ca II, [Co III] λλ5890, 5908,6197) are in emission. We also indicate a telluric feature (O2 A) in the ﬁrst spectral epoch with a ⊕ symbol. 36 be in a phase of linear decay and at a rate of 0.011 mag day−1, although later epochs in the B and V bands may be decaying more slowly or even rebrightening. We explore this behaviour in Section 2.4.3.

2.3.2 Spectroscopy

Comparison to Other SNe Ia/IIn and Spectral Line Identiﬁcation

In Figure 2.3 we compare spectra from each available epoch of PS15si to spectra of known SNe Ia/IIn. We note several features in absorption (grey) and in emission (blue) present in each spectral epoch. The comparison spectra were sampled from all of the SNe Ia/IIn spectra presented in Silverman et al.(2013a). For each epoch, the time (in days) since discovery is indicated for PS15si and the comparison spectra, including CSS120327:110520015205 (hereafter, CSS120327; identified as a SN Ia/IIn in Drake et al., 2012), SN 2008cg (identified as a SN Ia/IIn in Filippenko et al., 2008), SN 1999E (identified as a SN Ia/IIn in Deng et al., 2004), and PTF10iuf (identified as a SN Ia/IIn in Silverman et al., 2013a). The

PS15si spectra are scaled such that FHα = 1, and the comparison spectra are corrected for redshift and Galactic reddening using values from Silverman et al.(2013a) (Tables 1, 2, 4, and 5 therein) and scaled to minimise the root-mean square (RMS) of the difference spectrum except for Hα (i.e., 6500–6600 A).˚

The correspondence between Hβ, He I λ5876, Mg II around 4400 A,˚ and Fe II 4700– 5100 A˚ in the first spectral epoch of PS15si and the spectrum of CSS120327 indicates the similarities between these objects. In our two moderate resolution epochs (Figure 2.3, right panels), Hα is narrower in the PS15si spectra than in the SN 2008cg and SN 1999E spectra. This discrepancy could be due to the difference in spectral resolution of the PS15si epochs and comparison spectra, especially as the comparison spectra appear to peak at a lower flux density relative to the continuum level. If the discrepancy is real, however, it could also be due to the relative difference in epochs between PS15si and the comparison spectra or an intrinsic difference in the strength of post-shock CSM. That is, the intermediate- width component of Hα, which begins to appear 82 days after discovery in PS15si, is significantly weaker in PS15si than in the comparison spectra. This feature presumably 37 arises from the shell of post-shock CSM entrained in the ejecta (Chevalier and Fransson, 1994; Zhang et al., 2012), and therefore indicates the relative epoch of PS15si and the comparison spectra assuming similar CSM profiles. In the final epoch, there are additional differences in the comparison to PTF10iuf. Broad emission features are generally well-matched to PTF10iuf, although the continuum level appears slightly higher in the comparison spectrum. Emission around 5900 A˚ and especially 6200 A˚ is also stronger in PS15si than in PTF10iuf. These features are of particular interest because they correspond to known Co features in the nebular spectra of SNe Ia (i.e., [Co III] λλ5890, 5908 and [Co III] λ6197 as described in Bowers et al.

(1997); Liu et al.(1997)). In addition to [Co III] λ6578 (which is obscured by Hα in SNe Ia/IIn), these Co lines are generally the strongest Co features in optical spectra of SNe Ia. We interpret this line identification as independent confirmation that PS15si is likely a SN Ia/IIn and results from a thermonuclear SN; SNe IIn generally do not have strong Co emission in their late-time spectra but are instead dominated by narrow lines of H, O, and He as well as a quasicontinuum of forbidden and permitted Fe lines (as in SNe 1988Z, 1998S, 2006gy Turatto et al., 1993; Mauerhan and Smith, 2012; Kawabata et al., 2009). The presence of strong Co emission in addition to the quasicontinuum of Fe lines at bluer wavelengths strongly suggests that the underlying SN is Type Ia. One strong similarity between PS15si and other SNe Ia/IIn in the final epoch is the presence of narrow [O III] λ5007 emission, which is also accompanied by narrow [O I]

λ6364. PS15si represents the first time [O I] emission has been identified from a SN Ia/IIn. The emission from these features cannot be due to host galaxy contamination given the lack of any such emission in the first spectral epoch. Therefore, we infer that the [O I] and [O III] emission arise from either the SN ejecta or the surrounding medium. In spectra from SN 2005gj, Aldering et al.(2006) found that [O III] around 71 days past explosion exhibited an inverted P Cygni profile with emission around −100 km s−1 and absorption around +450 km s−1. Our spectral resolution in the final epoch (∼ 500 km s−1) cannot resolve any P Cygni feature in [O III] or [O I] if it exists. This line may also be obscured by strong Fe emission at similar wavelengths8.

8 We mark [Fe II] λ4890 and [Fe III] λ5011 in Figure 2.3 for reference. 38

The Ca II IR triplet feature appears to be one of the dominant spectral lines in the ﬁnal spectral epoch of PS15si. For SN 2002ic, Chugai et al.(2004) argued that Ca II emission around 8500 A˚ points to a Ca-rich layer in the shocked SN ejecta arising from an outer layer of Fe-poor material. In Figure 2.4, we show that their model of a Ca II line-emitting shell closely matches the proﬁle observed toward PS15si in the latest epoch. However, that quasicontinuum, Fe-dominated model around 4300–6000 A˚ is poorly matched to PS15si.

Strong Ca II emission is a sign that incomplete burning up to Fe-peak products may ac- company some SNe Ia/IIn (Marion et al., 2003), which can occur in thermonuclear SNe in 9 general . This prominent, broad Ca II proﬁle is common for SNe Ia/IIn; indeed, it appears in virtually all SNe of this type with spectra available ∼ 70 − 400 days after optical maximum (e.g., SNe 1999E, 2002ic, 2008J, 2011jb, PTF11kx, 2012ca in Rigon et al., 2003; Wang et al., 2004; Silverman et al., 2013a,b; Fox et al., 2015). This feature contrasts with spectra of “normal” SNe Ia in which Ca II is generally weak during the nebular phase (Bowers et al., 1997), although it has been observed at very late times (e.g., > 500 days in SN 2011fe as in Graham et al., 2015).

Fitting PS15si to Spectral Templates with Added Blackbody Emission

Apart from strong, narrow Hα emission, the presence of luminous, optically thick continuum is one of the deﬁning characteristics of SNe Ia/IIn that led to their interpretation as CSM interaction events (Hamuy et al., 2003; Aldering et al., 2006; Leloudas et al., 2015). Strong continuum emission tends to “dilute” the SN Ia spectrum, as emission and absorption features lose contrast relative to the continuum level. This is seen in core-collapse SNe IIn with strong CSM interaction as well (e.g., SN 2006gy as in Smith et al., 2010). As- suming the continuum emission is dominated by thermal emission from CSM interaction, we can constrain the temperature of the ejecta-CSM region and subtract this component to examine the underlying SN Ia-like features. This analysis is based on the assumption that the SN itself is typical of SNe without signatures of CSM interaction at varying epochs. In the case of PS15si, we simultaneously ﬁt a blackbody continuum and SNe templates to

9However, note that Mazzali et al.(2005) discuss abundance enhancement versus density enhancement in combination with SN-CSM interaction to explain Ca II IR absorption features in some normal SNe Ia. 39

Figure 2.4: Final spectral epoch of PS15si (2015 June 16) in black. In red, we show a comparison, synthetic spectrum taken from Chugai et al.(2004, Figure 10(c) therein). 40

the spectrum at each epoch. These templates include SN 1994D-like (SN Ia-norm; 0–90 days after explosion), SN 1991T-like (overluminous SN Ia; 0–93 days), SN 1991bg-like (underluminous SN Ia; 0–113 days) and SN 1999ex-like (SN Ib/c; 0–85 days) spectra (Nu- gent et al., 2002; Hamuy et al., 2002; Stern et al., 2004)10. While we have established that PS15si is spectroscopically similar to other SNe Ia/IIn and has spectral features typical of SNe Ia, we use a SN Ib/c template in our analysis as an additional check on the hypothesis that PS15si is more similar to SNe Ia than other types.

The template spectra (P0(λ)) we obtained were in arbitrary units, and we normal-

P0(λ) ized them by taking Pλ = . For each PS15si spectrum Fλ, we simultaneously P0(5000 A)˚ ﬁt a template spectrum Pλ and a blackbody continuum spectrum normalized to unity at 5 −1 hc hc hc its peak (λpeak = ) such that Bλ(T ) = exp − 1 with 4.965kB T 1.842kB T λ λkB T variable temperature T and the relative weights of the template and blackbody C and D, respectively. Thus, we searched for parameters, C,D,T to minimise the RMS of the difference spectrum

Zλ = Fλ − C × Pλ − D × Bλ(T ) (2.1)

Our results from spectral fits at all four epochs of PS15si are presented in red, blue, orange, and green for each template in Figure 2.5. The template spectra are labeled with the epoch in days relative to explosion along with the temperature of the blackbody continuum added to the fit. Below the best-fit spectrum from each template, we plot the residuals for all four templates (i.e., Zλ). The RMS of Zλ for each spectral fit is given in Table 2.4. In calculating the RMS, we ignore wavelength ranges that may be fit to features not present in the template spectra, such as Hα (6400–6700 A)˚ as well as the telluric feature in the first epoch (7200–7300 A).˚ We consider other features to be too weak to skew our fitting routine significantly and thus we include, for example, the wavelength range around Hβ, which also contains Fe II features. In this way, the statistic that we use to fit template spectra to epochs of PS15si, the RMS of Zλ, is useful primarily to compare different templates and fitting parameters within a single epoch and not as an overall “quality of fit” parameter. In Section 2.3.2, we assess the overall quality of our spectral fitting and alternative statistics

10c3.lbl.gov/nugent/nugent_templates.html 41

Table 2.3: Best-Fit Parameters C,D,T , and fV used to generate spectra in Figure 2.5.

Template Par.11 2015-03-28 2015-05-01 2015-06-13 2015-06-16 +3.9 +1.7 +3.7 +8.0 91bg-like C 0.9−0.9 1.4−1.4 3.5−3.5 5.9−5.9 +7.8 +2.3 +4.5 +6.8 D 27−7.7 7.5−2.3 12.5−4.6 10.8−6.6 +2300 +18000 +1500 +7700 T 6400−1600 3200 −1100 2700−1100 4500−2100 fV 0.07 0.33 0.48 0.54

+7.0 +1.3 +2.5 +6.5 91T-like C 6.1−6.1 1.0−1.0 2.8−2.7 11.1−6.2 +6.3 +2.4 +4.6 +14.4 D 21.4−6.4 7.1−2.4 11.0−4.7 12.1−12.1 +2000 +26000 +19000 +1700 T 5600−1700 3400 −1300 2800 −1300 2100−2100 fV 0.32 0.31 0.62 12.5

+7.4 +1.8 +3.0 +5.1 99ex-like C 4.8−4.8 1.2−1.2 0.1−0.1 8.3−5.2 +7.7 +2.3 +4.0 +5.6 D 24.9−7.6 6.8−2.2 13.7−4.4 7.4−5.7 +2700 +42000 +17000 +5400 T 6500−1700 3700 −1400 3500 −1400 4000−2400 fV 5.4 6.6 141 1.2

+7.1 +1.9 +3.1 +7.6 94D-like C 7.8−7.1 1.1−1.1 3.8−3.4 4.7−4.7 +6.3 +2.2 +4.0 +7.0 D 19.7−6.2 7.4−2.3 10.8−4.1 11.3−6.9 +2200 +42000 +15000 +7700 T 5400−1700 3500 −1200 3000 −1300 4500−2200 fV 0.44 0.21 0.53 0.43

in order to determine which spectral template best matches PS15si. We report the best-ﬁt values C,D, and T for each epoch of PS15si and set of template spectra (SN Ia 1991bg-like, SN Ia 1991T-like, SN Ib/c 1999ex-like, and SN Ia-norm 1994D-like) in Table 2.3. We also include the ratio of the weighted template spectrum to

blackbody emission (C × Pλ to D × Bλ(T )) in the V band (i.e., 5028–5868 A),˚ which we refer to as fV (see Leloudas et al., 2015). For each of the three fitting parameters, we approximate an uncertainty by fixing two parameters at their best-fit values and varying the third parameter until the RMS of Zλ (given in Table 2.4) increases by a factor of 2.

Quality of Spectral Fitting and the SN Ia Spectrum Underlying PS15si

As Figure 2.5 and Table 2.3 demonstrate, the statistical leverage in identifying PS15si comes mostly from the ﬁrst and last epochs given the much wider spectral range. Indeed, the spectral identiﬁcation, including type and epoch, as well as the intensity and temperature of blackbody emission in the other two (high-resolution) epochs is highly uncertain. While we include these spectral epochs in most of our analysis, they do not contribute

11C and D in units of 10−17 erg s−1 cm−2 A˚ −1, T in units of K. 42

significantly to our conclusions regarding spectral fitting. The analysis described above is similar to the Monte Carlo simulation used by Leloudas et al.(2015) to fit SNe Ia/IIn spectra. Similarities between these analyses are

that the template spectrum Pλ spans a range of types including SNe Ia-norm, SN 1991T- like, SN 1991bg-like, and SNe Ib/c, the variables D and T account for the blackbody radius and temperature, respectively, and we fix the spectral resolution of the template and blackbody to that of the input PS15si spectrum. Unlike the Leloudas et al.(2015) analysis, we do not explicitly fit the line spectrum associated with CSM interaction or artificially inject noise into our comparison spectra, we fix the reddening to the value noted above for PS15si (i.e., we use a dereddened PS15si spectrum as input), and we do not use human classifiers but rather rely solely on statistical analysis to identify the best-fit spectra.

A key parameter in the Leloudas et al.(2015) analysis is fV , which we describe in Section 2.3.2 and include for each PS15si epoch and template spectrum in Table 2.3. One

of the central arguments in Leloudas et al.(2015) is that, when fV is large (e.g., > 1), dilution from the CSM emission is low enough that spectral identiﬁcation is more accurate. Given that CSM emission is dominant at late times after explosion, this analysis would suggest that earlier spectra are more reliable when identifying the underlying SN. Our analysis reveals some curious trends in matching PS15si to SN templates of varying types. In general, the SN Ib/c 1999ex-like spectra yield poor ﬁts to PS15si with the

SN emission dominant over the CSM emission (fV > 1) in all epochs. This trend and the large uncertainties on fitting parameters for the SN Ib/c template suggest this spectrum is unlikely to describe the underlying SN emission from PS15si. In addition to our analysis in Section 2.3.2, from the fact that PS15si does not exhibit a clearly broad Hα component from emission in its ejecta and the poor fit between PS15si and SN Ib/c templates, we infer that PS15si must be a better match to SN Ia templates. This appears to be the case, as in each spectral epoch the SN Ia templates yield lower RMS values and the best fits to PS15si overall. While there is no clear preference in the RMS values for any particular template, the fV values and uncertainties on the fitting parameters suggest that the SN 1991bg-like template is not physically plausible. For example, in the first spectral epoch when emission from the underlying SN should be most dominant, 43

fV = 0.07 for the SN 1991bg-like template makes it clear that the template is fitted poorly to the PS15si spectrum. The RMS value for this spectral fit supports the hypothesis that the SN 1991bg-like template is the poorest fit to PS15si in this epoch. In the final epoch, however, PS15si is better fit to the SN 1991bg-like template, although that may be due to the relative lack of continuum emission in underluminous SN 1991bg-like SNe Ia or else the fact that the available epochs of the SN 1991T template (0–93 days) do not cover as large a range as the SN 1991bg templates (0–113 days). As we discuss below for the SN 1991T-like template, it is difficult to decompose a SNe Ia/IIn into its constituent SN and CSM spectrum given the degeneracy between these components. One might expect that if SN 1991bg-like SNe Ia are a better match to SNe Ia/IIn than SN 1991T-like SNe Ia (or vice versa), then SNe Ia-norm should fall between the two in terms of quality of fit. This trend appears to be the case in both the first and last epochs when PS15si matches the 1994D-like template about as well as the SN 1991T-like and

SN 1991bg-like templates, respectively. The value of fV is comparable to what Leloudas et al.(2015) find for post-maximum SNe Ia/IIn templates, although it barely decreases between the best-fit values for the first and last epochs of the 1994D-like template, contrary to the observation that SNe Ia fade much more rapidly than CSM emission. Again, there are problems in that the decomposition does not fully recover the relative contribution of these components in a physically consistent manner. Based on the uncertainties and the self-consistent evolution of the temperature across all four epochs, SN 1991T appears well-fit to PS15si. However, in the final epoch there are some differences on the blue end of the spectrum that suggest there is extra CSM emission diluting the continuum between emission features as well as much weaker emission around

5800–6000 A,˚ which in SN 1991T was identified in nebular spectra as [Co III](Schmidt et al., 1994). Overall, SN 1991T is the poorest fit to PS15si in this epoch and the fit to

SN 1991T involves an extremely large value of fV (= 12.5) for this late epoch, although both of these effects might be explained by the poor approximations involved in decomposing a SN Ia/IIn spectrum into a SN Ia and CSM emission. For example, SN 1991T-like SN Ia spectra are associated with overluminous and hot SNe with signiﬁcant 56Ni production and whose spectra exhibit high-ionisation lines (Filippenko et al., 1992). Perhaps the 44

Table 2.4: RMS of Residuals for Each Template in Figure 2.5

UT Date 91bg-like 91T-like 99ex-like 94D-like (y-m-d) 2015-03-28 0.067 0.062 0.065 0.063 2015-05-01 0.048 0.051 0.052 0.050 2015-06-13 0.039 0.043 0.045 0.041 2015-06-16 0.049 0.081 0.065 0.049

high luminosity that is associated with 56Ni production in SN 1991T-like SNe Ia instead comes almost entirely from CSM interaction in SN Ia/IIn, which would explain the poor

fit to [Co III] in the final epoch of PS15si as well as the poor fit to continuum emission

between spectral lines and the unusually large fV ratio. We discuss further ramiﬁcations of this idea in Section 2.4.2.

2.4 Discussion

2.4.1 PS15si Explosion Date and Maximum R band Magnitude

Comparison between our photometry and light curves from other SNe Ia/IIn provides a method to infer the phase of PS15si. We acknowledge that this analysis assumes that the shape of the PS15si light curve near maximum, which is not well-constrained by our data, is similar to that of other SNe Ia/IIn. Recent analysis comparing several SNe Ia/IIn (SNe 2005gj, PTF11kx, 2002ic, 1999E, 1997cy in Inserra et al., 2016, and Figure 4 therein) found that light curves from most SNe Ia/IIn have a similar shape and bolometric luminosity from early times to as late as ∼ 200 days after optical maximum except for the unusual case of PTF11kx, whose bolometric luminosity fades much more quickly and is very similar to SN 1991T. For the majority of SNe Ia/IIn, variations in CSM profiles for each object may cause bolometric light curves to diverge significantly after ∼ 200 days. SN Ia/IIn absolute magnitudes also span at least 1.8 mag in the r band at maximum brightness (Sil- verman et al., 2013a), so calibrating a light curve to these objects will have similarly large systematic uncertainties. In Figure 2.2, we compare the available R-band light curve of PS15si to the exponential decay (i.e., in flux versus time) phase of SN 2005gj (Prieto et al., 2007). Estimating 45

Figure 2.5: Four spectral epochs of PS15si with the day relative to discovery (d##) given in black. We also plot comparison SN Ia 1991bg-like templates (red), SN Ia 1991T-like templates (blue), SN Ib/c 1999ex-like templates (orange), and SN Ia-norm 1994D-like templates (green) plus blackbody continuum. These template spectra are derived from those available in Nugent et al.(2002, 1991bg-like and 1994D-like, also c3.lbl.gov/ nugent/nugent_templates.html), Stern et al.(2004, 1991T-like), and Hamuy et al.(2002, 1999ex-like). We label the template spectrum with the epoch (d##) in days relative to explosion as well as the temperature of the added blackbody continuum. We also plot the residuals for each comparison spectrum below the stacked spectra. 46

the R-band maximum brightness of PS15si involves a degeneracy between the maximum brightness, epoch of observation, and the “stretch” or rise and decay time. From our R-band photometry in Figure 2.2 and comparison to the SN 2005gj light curve, we find that both objects are declining at roughly 0.011 mag day−1. This rate is also significantly slower than the 0.025 mag day−1 observed for 56Co decay in normal SNe Ia during the epoch 30 − 90 days after explosion (see, e.g., Phillips et al., 1999;F orster¨ et al., 2013), although it is consistent with the decline rates observed in the R band for other interaction- powered SN light curves (e.g., SNe 2005cp and 2005db as in Kiewe et al., 2012). Therefore, we fit the SN 2005gj R-band light curve to our PS15si data assuming some variation in peak absolute magnitude, relative epoch, and “stretch” factor in the temporal axis. That is, for SN 2005gj R-band band magnitudes corrected for distance modulus

µ = 36.71 mag and extinction AR = 0.306 mag (Prieto et al., 2007) and normalized in

time such that the discovery date corresponds to tSN 2005gj = 0, we ﬁt our PS15si data to this light curve a correction in absolute magnitude δM, relative epoch δt, and temporal stretch factor η. We found that the best-ﬁt values corresponded to δM = 0.56, δt = −14.2

and η = 1.2 with the magnitudes of the light curves given by mPS15si = mSN 2005gj + δM

and the temporal data given by tPS15si = η × tSN 2005gj + δt. The PS15si data and stretched and shifted SN 2005gj light curve are depicted in Figure 2.2. Pre-discovery detections of PS15si in w-band on 23 Feb. 201512, 27 days before discovery, provide a strong constraint on the rise of PS15si. Assuming PS15si is similar to SN 2005gj and follows the model given in Figure 2.2, the SN exploded 42 days before discovery and peaked approximately 6 days before discovery with an absolute magnitude

of mR = −19.80. This model implies that PS15si had a rise time slightly longer than SN 2005gj of ∼ 36 days in R band, which agrees with the range of rise times observed for SNe Ia/IIn (18 and 45 days in Silverman et al., 2013a).

12These data are available from the Pan-STARRS Survey for Transients at star.pst.qub.ac.uk/ ps1threepi/ and described in Huber et al.(2015) 47

2.4.2 Luminosity of the CSM and Underlying SN

For CSM-interacting SNe at late times and for a high wind-density parameter w = 4πr2ρ > 1016 g cm−1, the overall luminosity of the radiative forward shock wave with 3 velocity vs will be L ∝ wvs (Chugai and Yungelson, 2004); this is the rate at which the 2 2 shock sweeps up mass (4πr ρvs) times the energy per unit mass in the shock (vs ). Following the treatment presented in Ofek et al.(2014), we assume that ρ = br−k, 2k−7 2−k 3 − k−4 which implies L ∝ r vs ∝ t , where we have implicitly assumed that ρej ∝ −3 r −m 13 t ( t ) , and m = 4, which holds for a momentum-conserving “snowplough” SN (Svirski et al., 2012). Assuming the CSM component of the R band data in Table 2.1 and Figure 2.2 follows a a power-law with LR ∝ t , i.e., MR ≈ −2.5a log t (with t in days since discovery), we subtract the R band luminosity from a SN 1991T-like template to determine the contri-

bution to the luminosity from CSM emission (i.e., L = LCSM + LSN ) and find that the −0.727t residual magnitude is well-fit by MR ≈ 1.82 log t, which implies LR ∝ t . This trend in luminosity implies k = 3.5. The inferred power-law is relatively insensitive to the explosion time. For example, if we add 26 days to the time to model the luminosity versus time since explosion according to our analysis in Section 2.4.1, then we find −1.02t MR ≈ 2.56 log t, which implies LR ∝ t and k = 3.7. Clearly, the density profile implied by our photometry and this model is very steep. Indeed, the self-similar CSM treatment of Chevalier(1982) breaks down for k > 3, although similar results have been reported in the much later (t ≈ 300 days) optical light curve of SN 2010jl (Ofek et al., 2014). From our spectral fitting to SN 1991T-like templates, we find that a thermal continuum component fit to the spectrum of PS15si at our first and last epochs implies best fits with T = 5600 K and 2100 K, although as we discuss below, the latter temperature may be an underestimate. The second and third (higher-resolution epochs) have highly uncertain temperatures which we do not include in this analysis. Assuming the SN has expanded to ∼ 1016 cm by 86 days after discovery (i.e., 100−120 days after explosion given the model

13 α (2−k)(m−3)+3(k−3) 2k−7 (m−3)/(m−k) Ofek et al.(2014) find L = L0t , where α = m−k = − k−4 , r ∝ t = t−1/(k−4) for m = 4. 48 in Section 2.4.1 and with an expansion velocity of 104 km s−1), a spherical shell emitting at 2100 K will have a luminosity of about 1.4 × 1042 erg s−1 or M ≈ −16.7 mag. This value is far too small to account for the amount of luminosity observed at the final epoch (M ≈ −19) when emission from the CSM should dominate the supernova. There is some inconsistency in the overall shape of the spectrum at late times, especially as features from the SN are easily visible at a time when the CSM should outshine the SN emission by a few magnitudes (e.g., Figure 2.2 herein, and Figure 7 in Prieto et al., 2007). However, ultraviolet/X-ray light from the reverse shock may contribute significantly to emission from the ejecta, producing a quasicontinuum with distinct spectral features. We have found that PS15si can be fit by overluminous 1991T-like SNe Ia spectra. This underlying SN type is consistent with other SN Ia/IIn, but there are some remaining inconsistencies in terms of the continuum level and the relative strength of the [Co III] feature in the final epoch. Perhaps some of the continuum emission that comes from CSM interaction in SNe Ia/IIn is misinterpreted as the hot, luminous continuum emission observed in 1991T-like SNe Ia. At the same time, some of the light from the CSM interaction in SNe Ia/IIn could be reprocessed by the ejecta, changing the ionisation state near the interaction region and heating new layers of ejecta as the SN evolves. This hypothesis would also explain why the spectral continuum of some SNe Ia/IIn appears to evolve from thermal to nonthermal, as in SN 2002ic (Chugai et al., 2004) and SN 2005gj (Aldering et al., 2006). In the former, the continuum emission was well-fit around 244 days after optical maximum by a quasicontinuum model composed of cool shocked ejecta in lines of Fe- peak elements. Varying the mixture of elements in models of the line-emitting shocked ejecta layer may simultaneously reproduce the observed continuum level, SN Ia-like absorption features, and anomalous line strengths such as the Ca II emission we noted in Section 2.3.2. One consequence of this hypothesis would be that, with an added source of ultraviolet/X-ray emission, we would expect to observe an increased fraction of high- ionisation lines in the ejecta. Indeed, the correlation between most SNe Ia/IIn and SN 1991T-like spectra may largely be due to the added radiation from the CSM interaction, as SN 1991T exhibited a blue spectrum dominated by Fe III absorption around 49 maximum (Filippenko et al., 1992; Ruiz-Lapuente et al., 1992) and [Fe III] and [Co III] emission dominant in its nebular spectrum (Schmidt et al., 1994). This hypothesis may explain the spectral morphology of SNe Ia/IIn, although the dif- ficulty in identifying the characteristics of the underlying SN remain. For example, one might expect that the underlying SN Ia spectrum should be correlated with the amount of radiation produced by the CSM interaction and thus the peak luminosity of the SN. Stud- ies suggest that SNe Ia/IIn with relatively low and high peak luminosities (e.g., SN 2008J and PTF10iuf in Taddia et al., 2012; Leloudas et al., 2015) have been associated with SN 1991T-like spectra. Are there added effects due to the distribution of the CSM, view- ing angle, or dilution of the underlying spectrum from the CSM quasicontinuum? In addition, the outer layers of the SN ejecta will receive most of the added radiation from the CSM interaction, and so the composition is largely dependent on the underlying thermonuclear explosion. While it is worth speculating on these questions, it may be impossible to disentangle the relative importance of these effects without detailed hydrodynamic and radiative-transfer models of CSM-interacting SNe.

2.4.3 Late-Time Rebrightening and Spectral Fitting of SNe Ia/IIn

PS15si appeared to be rebrightening in B band (increase of 0.31 mag with 1.6σ significance) and V band (increase of 0.32 mag with 2.3σ significance) between 85 to 90 days after discovery. This trend was observed in the last four epochs of our optical photometry (Table 2.1) in B and V . Either the rebrightening was not as apparent in NIR observations or the trend occurred over a relatively short period of time and the SN was no longer increasing in brightness by 104 days after discovery. Our final spectral epoch corresponds to the beginning of this trend in the photometry. The flux from a CSM-interacting SN light curve may look more like a power law at late times as opposed to exponential decay. This may cause the slope of the decay curve to change where the light curve appears to be leveling off as the CSM cools radiatively as in Section 2.4.2. This hypothesis does not fully explain the behavior in V , however, where the light curve is systematically brightening over the final three epochs by as much as ∼ 30% in luminosity (0.32 ± 0.13 mag). 50

This behavior supports the conjecture that the CSM density proﬁle around SNe Ia/IIn can be dense and clumpy, perhaps with steep density gradients associated with these clumps (e.g., as suggested by spectropolarimetry of Hα toward SN 2002ic; Wang et al., 2004). Several other SNe Ia/IIn, such as SN 1997cy (Inserra et al., 2016), have exhibited variations in their bolometric light curves over short timescales. This behavior contrasts with rebrightening seen in the mid-IR toward SN 2005gj (Fox and Filippenko, 2013) where the increase in luminosity was observed predominantly at longer wavelengths (3.6−5 µm), over longer timescales (∼ 600 days), and starting at least a year after the SN was discovered. In this latter case, rebrightening was thought to be due to reprocessed light from a dust shell surrounding the SN. For PS15si, interaction between the SN ejecta and clumps or a thin shell in the CSM can account for the short timescale over which the rebrightening occurred. As with core-collapse SNe IIn (e.g., SNe 2001em and 2006gy as in Schinzel et al.(2009) and Smith et al.(2008a), respectively), late-time X-ray and radio-wavelength studies of SNe Ia/IIn may reveal much about the CSM proﬁle.

2.5 Conclusion

Spectra of PS15si are well ﬁt by spectra of other SNe Ia/IIn observed previously (Silver- man et al., 2013a). Comparisons to these examples and some SN Ia subtypes, however, reveal new underlying diversity in this class of SNe, as follows.

1. PS15si is best ﬁt by spectra of overluminous SNe Ia (e.g., SN 1991T) if we add extra thermal continuum, but there are inconsistencies in the continuum level at different

epochs as well as the strength of [Co III] emission in the ﬁnal epoch. This spectral morphology matches the interpretation of other SNe Ia/IIn, although we interpret this similarity as an indication of the added continuum emission and changes in the ionisation state of the visible ejecta layers brought on by CSM interaction. De- tailed radiative-transfer models of shocked ejecta illuminated by CSM interaction are needed in order to satisfactorily reproduce SN Ia/IIn spectra, but PS15si illus- trates that the similarity to overluminous SNe Ia is incidental to spectra of SNe Ia/IIn and not indicative of the underlying SN. 51

2. PS15si appears to have rebrightened over a short timescale at around 85 days after discovery. Similar behavior was also observed by Inserra et al.(2016) for SN 1997cy and in other SNe Ia/IIn where the mechanism was assumed to be clumpiness or steep density gradients in the CSM. Especially at late times, SN Ia/IIn environments are poorly ﬁt by models assuming uniform CSM. SNe Ia/IIn may be good candidates for late-time radio and X-ray observations, which can probe the properties of the CSM out to large distances from the progenitor. 52

CHAPTER 3

Variability in the Near-Infrared Synchrotron Emission From Cassiopeia A

We present multi-epoch Ks band imaging of the supernova remnant Cassiopeia A (Cas A). The morphology of the emission in this band is generally diffuse and filamentary, consistent with synchrotron radiation observed at radio wavelengths. However, in one region to the southwest of the remnant, compact knots of emission appear to be entrained in the ejecta and have the same proper motion as ejecta observed at similar projected radii. The presence of these knots suggests that material with high magnetic field strength contributes significantly to synchrotron emission at these wavelengths. We analyze these knots at J, H, and Ks bands as well as in 3.5 − 8 micron emission and at 6 cm where synchrotron emission is dominant and we find that the Ks band emission falls along the expected synchrotron spectrum. Using multi-epoch data, we calculate the magnetic field strength and electron density for a population of near-infrared synchrotron-emitting electrons. We find electron densities from 1, 000 − 15, 000 cm−3 and magnetic field strengths from 1.3 − 5.8 mG. These magnetic field strengths are an order of magnitude higher than inferred from the much lower angular resolution gamma-ray observations toward Cas A. We conclude that dense knots of post-shock material behind the Cas A shock front are emitting synchrotron emission in a compressed and enhanced magnetic field.

3.1 Introduction

Acceleration of cosmic rays (CRs) in supernova remnants (SNRs) contributes signiﬁcantly to the population of high energy electrons up to the “knee” of the CR spectrum around 1015 eV. This population has been observed toward SNRs via synchrotron radiation and other nonthermal emission from radio through gamma ray energies. The diffusive shock acceleration model has largely succeeded in tying the radio synchrotron spectrum observed toward SNRs to GeV electrons (Bell, 1978). At higher energies, electrons with E > 1 TeV 53 have been associated with X-ray emission in the form of synchrotron radiation (Reynolds, 1998), inverse Compton scattering (Tanimori et al., 1998; Porter et al., 2006), and nonthermal brehmsstralung radiation for the hardest X-rays above 100 keV (Vink, 2008a). GeV gamma-rays from SNRs likely originate from leptonic CRs, especially toward SNRs where particle acceleration is enhanced due to interaction with a molecular cloud (e.g., Castro et al., 2013). Even for gamma-rays with E ≥ 1 TeV, it is ambiguous to what extent the emission originates from leptonic as opposed to hadronic particle acceleration (Ellison et al., 2010; Inoue et al., 2012; Slane et al., 2015).

Near-infrared Ks band emission from SNRs is thought to be dominated by synchrotron emission. At these energies, synchrotron radiation requires electrons with higher energies (∼ 0.2 TeV) than those that emit predominantly in the radio. In support of this hypothesis, α spectroscopic measurements of the Ks band spectral index α (where Fν ∼ ν ) toward Cassiopeia A (Cas A) with −0.80 < α < −0.67 are consistent with a synchrotron spectrum (Wright et al., 1999; Rho et al., 2003; Eriksen et al., 2009). In addition, the fractional polarization of Ks band emission (5-10%) is consistent with measurements of synchrotron emission around 6 cm (Jones et al., 2003). Measurements of the emission in this region therefore provide a unique probe of the synchrotron spectrum while offering a test against which radio, X-ray, and gamma-ray synchrotron measurements can be compared. One of the most significant advantages of near-infrared measurements of Cas A involves the timescales of synchrotron losses. Given the magnetic field strengths and electron energies involved, infrared-emitting electrons are likely to have been accelerated no more than ∼ 80 yr ago while shocks can produce infrared-emitting electrons on timescales of ∼ 1 yr (Jones et al., 2003; Rho et al., 2003). Therefore, multi-epoch measurements of near-infrared synchrotron emitting material on timescales of several years are likely to be sensitive to variability in the acceleration of electrons over baselines of 1 − 10 yrs. In turn, this variability is a direct probe of the electron density and magnetic field accelerating these electrons. Additionally, current magnetic field strength estimates for the radio-emitting plasma in Cas A are based on gamma-ray fluxes measured with poor angular resolution and thus averaged over multiple acceleration sites. The magnetic field strength over this region 54 is related to the distribution of relativistic particle energies and the bremsstrahlung flux emitted by those particles, which is thought to be the dominant emission process at GeV energies. Cas A is detected by Fermi with a beam size of 0◦.1 as a single GeV source and this emission is well-fit by a leptonic model with a magnetic field strength of B ≈ 0.12 mG (Abdo et al., 2010). Alternative analyses based on radio, infrared, X-ray, and Fermi gamma-ray data suggest B ≈ 0.05−0.30 mG (Araya and Cui, 2010), 0.23−0.51 mG (Saha et al., 2014), and 0.2 − 0.4 mG (Zirakashvili et al., 2014). Although the production of GeV gamma-ray emission via the leptonic process predicts short cooling timescales and thus emission originating near the Cas A forward shock (e.g., Esposito et al., 1996; Gotthelf et al., 2001; Abdo et al., 2010), these measurements cannot be localized to specific regions where significant magnetic field amplification might occur. In all of these studies, in effect the magnetic field strengths are averaged over the entire relativistic electron population of Cas A. On small scales, compression and turbulent amplification may lead to significant magnetic field amplification. This amplification should be measurable in the near-infrared when nonthermal emission is observed over multiple epochs. In this paper we use multi-epoch Ks band imaging to constrain the proper motions and variability in knots of synchrotron emission toward Cas A. We look for any features that are well-resolved in the imaging and compare them to archival data in other wavelengths to verify that the emission from these features falls along a synchrotron spectrum. Finally, we derive magnetic field strengths and electron densities for these features and compare them to values calculated over the entire remnant.

3.2 Observations

We obtained near-infrared Ks band imaging of Cas A using PISCES on the Bok 2.3m telescope on 11 Nov. 2013. The PISCES wide-ﬁeld camera (McCarthy et al., 2001) has a ﬁeld of view of 80.5 and a pixel size of roughly 000.5 and we were able to observe the entire

SNR in a single pointing. We employed a Ks band filter centered at 2.22 microns with a width of approximately 0.51 microns. We used 12 s individual exposures and alternated sky exposures in an on (source)- off-off-on pattern. For each individual exposure, we added a random 3000“wobble” vector 55 in order to observe the source at a random position on the array. The total on-source 00 exposure time in Ks band was roughly 32 minutes. We had 1 .2 seeing for the entirety of our observation. Standard image reductions were performed using IRAF1 including bias and dark subtraction, bad-pixel removal, flat-fielding, sky subtraction, distortion correction, image stacking, and registration. Flux calibration was achieved using 2MASS (Skrutskie et al., 2006) stars in the same field as Cas A. We correct for extinction toward the remnant assuming AV = 6.2 (as in Eriksen et al., 2009).

We present the ﬁnal Ks band image in Figure 3.1. In the subsequent analysis, we employ additional Ks band imaging from the literature, including epochs from 2002 (Rho et al., 2003) and 2003 (Eriksen et al., 2009).

3.3 Results and Analysis

3.3.1 Fast-Moving Features

The diffuse and ﬁlamentary structure at near-infrared wavelengths in Figure 3.1 highlights the presence of synchrotron continuum. However, there is at least one location where the

Ks emission appears as knots of emission (Figure 3.2). These knots appear to be entrained in the ejecta and with proper motions of at least 000.34 yr−1, although the second knot is below the level of detectability or obscured by a star approximately 4 mag brighter than the knots in the ﬁnal epoch. At the distance of Cas A (3.4 kpc Fesen et al., 2006), this proper motion corresponds to a velocity of ∼ 5500 km s−1, indicating that these knots of emission are likely associated with fast-moving knots (FMKs; e.g., van den Bergh and Dodd, 1970) in the ejecta. In the following sections, we consider emission processes in

Ks band that could account for compact knots of emission entrained in the ejecta.

1IRAF, the Image Reduction and Analysis Facility, is distributed by the National Optical Astronomy Observatory, which is operated by the Association of Universities for Research in Astronomy (AURA) under cooperative agreement with the National Science Foundation (NSF). 56

Figure 3.1: Bok 2.3m PISCES Ks band image of Cas A from 11 Nov. 2013 showing smooth synchrotron continuum emission. Individual knots of Ks band emission are resolved to the southwest of the remnant. 57

Figure 3.2: Imaging of the two synchrotron-emitting knots (SEKs) from 2002 (left), 2003 (right), and 2013 (right). The northwestern SEK appears as extended emission near two reference stars whereas the southeastern SEK has either disappeared or moved behind the stars. The knots are indicated with a red 100.5 circle in each image.

3.3.2 Are the Ks Band Knots Dominated by Synchrotron Emission?

A central question to this study is whether the observed Ks band knots represent synchrotron continuum, as has been argued in the past for the underlying emission, or are they dominated by line emission from the ejecta? Near-infrared spectroscopy of FMKs near the infrared knots reveals that there is very little line emission in this waveband, with only some contribution from Brγ, He I, [Fe II], and [Si IV](Gerardy and Fesen, 2001).

The bright [Si IV] line at 1.965 µm is outside the spectral range deﬁned by our Ks band ﬁlter and would not contribute to the emission observed in our images. The lack of any

other strong source of line emission in typical Ks band spectra of FMKs toward Cas A implies that the photometry is dominated by continuum emission, that is the nonthermal synchrotron spectrum. Line emission may dominate in J and H bands. In the 2003 epoch, knots 1 and 2

exhibit mJ = 15.81 ± 0.06 and 15.49 ± 0.04 and mH = 15.66 ± 0.07 and 15.41 ± 0.06, respectively. The ratio between these near-infrared bands is typical given that most near- infrared spectra toward Cas A reveal that J band is strongly dominated by forbidden line

emission, especially from [S II] and [P II] emission. Gerardy and Fesen(2001) ﬁnd that R the ﬂux (F = Fνdν) from J band line emission is at least an order of magnitude greater than line emission observed in H band. However, our measurements suggest that for 58

knots 1 and 2, F = νFν is about FJ = 1840, 2480 and FH = 1010, 1260 (in units of 10−15 erg s−1 cm−2), which is atypical for FMKs. Perhaps some other source of line emission, such as the [Fe II] lines typically observed in slow-moving quasi-stationary ﬂoculi, can account for the unusual ratio J to H band ratio for knots 1 and 2, or else the nonthermal continuum must account for a signiﬁcant enhancement in H band.

Comparison to radio continuum suggests that the Ks band is morphologically similar to wavebands dominated by synchrotron emission. Gerardy and Fesen(2001) found that the Ks band images have no clear optical, X-ray, or mid-infrared counterpart and the diffuse emission in these bands is most similar to radio continuum images. Jones et al. (2003) made a similar argument based on predictions of the polarization angle of synchrotron radiation from 2.2 µm, which closely matches the polarization angle observed at 6 cm.

Perhaps the most convincing argument that the Ks band features to the southwest of

Cas A are dominated by synchrotron emission is the comparison between the Ks band spectral indices and those observed in the radio. Rho et al.(2003) performed “spectral tomography” by matching the Ks band brightness to radio emission across Cas A and sub- α tracting the Ks band emission in proportion to (νradio/νIR) . This method simultaneously provides a check against the hypothesis that the near-infrared emission is well-fit by synchrotron continuum as observed in radio emission and a way to measure spectral indices across the remnant. We performed the same analysis using archival 6 cm VLA imaging toward Cas A obtained from 2000-2001 (see Roy et al., 2009). As we demonstrate in Fig- ure 3.3, the Ks band knots correspond to the position of a local enhancement in the radio synchrotron emission. At 6 GHz, knot 1 and knot 2 have flux densities of 0.55 ± 0.02 and 0.53 ± 0.02 Jy, respectively (see Table 3.1). The radio spectral index is ∼ −0.76 (Onic´ and Uroseviˇ c´, 2015), so extrapolating to 2.2 µm predicts flux densities of 0.27 and 0.26 mJy respectively, close to the observed values of 0.37 and 0.35 mJy (taking the averages for 2002 and 2003 - see

Table 3.1). It is therefore plausible that the Ks band emission is dominated by synchrotron radiation. There are indications of a ﬂattening of the synchrotron spectrum toward high radio frequencies (Onic´ and Uroseviˇ c´, 2015, see also Rho et al.(2003)), which would make 59

Figure 3.3: Images of the synchrotron-emitting knots at (top; left to right) 6 cm, IRAC Band 4, IRAC Band 3, IRAC Band 2, (bottom; left to right) IRAC Band 1, Ks, H, and J bands. The epoch is indicated in the upper-right of each panel. The red rectangle indicates the approximate position of the SEKs.

the extrapolation from 6 GHz to 2.2 µm even closer to the measured values and suggest the Ks band emission is almost entirely due to synchrotron radiation.

3.3.3 IRAC Measurements

In addition to radio emission at 6 cm and Ks band emission, Ennis et al.(2006) argued that some of the Spitzer Infrared Array Camera (IRAC) Bands (3.6 − 8 µm) toward Cas A are dominated by synchrotron emission. In particular, this study argued that Cas A is morphologically similar in 6 cm, IRAC Band 1 (3.6 µm), and Ks band emission. Therefore, we examine the knots in all four IRAC Bands in addition to radio emission and near-infrared emission. Full analysis (given below) indicates that IRAC Bands 2-4 are dominated by non-synchrotron emission while Band 1 is dominated by synchrotron emission.2

2This analysis is based on observations over a four year baseline, spanning from 2001 (date of 6 cm radio imaging) to 2003 (date of near-infrared imaging) to 2005 (date of IRAC imaging). Synchrotron losses or the injection of fresh electrons into the synchrotron-emitting population may account for small deviations from a reasonable synchrotron spectrum across the entire SED. 60

Table 3.1. Photometry of Synchrotron-Emitting Knots Toward Cas A

Knot (#) 6 cm 8.0 µm 5.7 µm 4.5 µm 3.6 µm Ks Ks Ks HJ (51794) (53388) (53388) (53388) (53388) (52289) (52833) (56607) (52833) (52833) (mJy) (mJy) (mJy) (mJy) (mJy) (mJy) (mJy) (mJy) (mJy) (mJy)

1 551±20 54.1±0.1 1.15±0.03 7.94±0.04 0.53±0.02 0.37±0.06 0.37±0.04 0.31±0.04 0.56±0.04 0.76±0.04 2 530±20 52.7±0.1 1.15±0.03 5.37±0.04 0.41±0.02 0.39±0.06 0.31±0.04 – 0.70±0.04 1.02±0.04

Note. — Flux densities for knots 1 and 2 from VLA 6 cm, IRAC Band 4-1 (8.0, 5.7, 4.5, 3.6 µm), and Ks (2002, 2003, 2013), H, and J band emission. We indicate the approximate Modified Julian Date (MJD) of each observation below each band. All flux densities are in terms of mJy. For the 2002, 2003, and 2013 Ks band epochs, the flux densities for SEK 1 correspond to magnitudes of 15.63, 15.64, and 15.82, respectively, while those for the 2002 and 2003 epochs for SEK 2 are 15.59 and 15.82, respectively.

We plot the SED of knots 1 and 2 across all eight wavebands depicted in Figure 3.3 in Figure 3.4. We have fit a synchrotron spectrum to the radio flux densities measured at each knot and assuming a spectral index of α = 0.78. The Ks and IRAC Band 1 flux densities are a good match to the assumed synchrotron spectrum, although every other point lies above the inferred power-law.

IRAC Band 2 (4.5 µm) is anomalously bright toward the Ks band knots with emission an order of magnitude brighter than the inferred synchrotron continuum level in both knots. Ennis et al.(2006) noted this brightness excess and inferred that line emission must dominate emission toward Cas A between 4 and 5 µm. We suggest that the most likely candidate is the CO fundamental vibrational band, which has been spectroscopically identiﬁed toward Cas A at 4.5 µm (Rho et al., 2009a, 2012). No spectra of the 4.5 µm emission are available toward the southwest shell where these knots occur, although the CO emission observed from AKARI at these wavelengths is generally ∼ 60 MJy sr−1 (1.4 mJy arcsec−2)(Rho et al., 2012), roughly in agreement with the 7.94 ± 0.04 mJy we observe over a 100.5 radius aperture toward knot 1.

It is known that IRAC Band 4 (8.0 µm) is dominated by [Ar II] (6.99 µm) emission (Ennis et al., 2006; Smith et al., 2009) while IRAC Band 3 (5.7 µm) may be dominated by dust continuum or possibly by forbidden line emission from [Fe II] (5.3 µm). The IRAC 61

Figure 3.4: SED of SEK 1 (left) and SEK 2 (right) including 6 cm radio emission, IRAC Bands 4-1, and Ks, H, and J band emission (from left to right in each panel). We overplot a synchrotron spectrum (dashed line) corresponding to the ﬂux density of the 6 cm emission and a spectral index of −0.71 and −0.73 for SEKs 1 and 2, respectively. 62

Band 3 emission from both knots is only ∼ 1.5 times the inferred synchrotron continuum level, which may imply a combination of dust continuum and synchrotron emission in Band 3.

Given the good ﬁt between the synchrotron spectrum from Ks band and the radio emission, we infer that the knots are both dominated by synchrotron emission in this band.

Henceforth, we refer to these features as synchrotron-emitting knots (SEKs) in Ks band. One intriguing hypothesis regarding the synchrotron emission is that features dominated by synchrotron radiation should be concentrated in compressed magnetic ﬁeld downstream from the shock front (see, e.g., Reynolds, 1998; Bleeker et al., 2001). It has been argued that this hypothesis is contradicted by the ubiquity of diffuse emission observed in

both Ks band and radio emission (Jones et al., 2003). However, the presence of compact

synchrotron-emitting knots in Ks would argue in favor of such a scenario indicating the magnetic ﬁeld enhancements occur in some locations.

3.3.4 Physical Parameters of Synchrotron-Emitting Knots

Given the multi-epoch imaging of the SEKs in Figure 3.2, we can use their measured ﬂux densities to estimate the physical conditions required to account for the observed decay rate. A single electron with Lorentz factor γ and emitting synchrotron radiation in a uniform magnetic ﬁeld with strength B emits power

1 P = σ c(γ2 − 1)B2. (3.1) 6π T 2 A high-energy (γ 1) electron with E = γmec emitting synchrotron radiation will therefore have a mean lifetime

8 E 3πmec 5.6 × 10 s τ = = 2 = 2 . (3.2) 2 ln(2)P ln(2)σT γB γB We assume that a SEK will decay in brightness at roughly the same rate. That is, for a SEK with luminosity L such that m ∼ −2.5 log(L), the magnitude of the synchrotron −t/τ t emission will have a dependence m ∼ −2.5 log(L0e ) mag ∼ 2.5 log(e) τ mag = 63

0.06γB2t mag yr−1. Thus, the magnitude of the synchrotron emission will decay at a rate dm 2 −1 dt = 0.06γB (B in gauss). The measured decay rate is 0.017 mag yr for SEK 1 across all three epochs and 0.27 mag yr−1 for SEK 2 from the first to second epoch (see note in Table 3.1). Assuming both SEKs are losing energy in a magnetic field of constant strength 3 2 eB and all electrons are emitting as a δ-function at their characteristic frequency νc = γ 4π mec 14 −1/2 (where νc = 1.4 × 10 Hz), then the electrons will have γ = 5800B (B in gauss) and 3/2 −1 the SEKs will decay at a rate 0.011BmG mag yr . From our measured decay rate, the magnetic fields toward SEK 1 and 2 are approximately 1.3 mG and 5.8 mG, respectively. To calculate the electron density of each SEK, we make the standard assumption that the number density of relativistic electrons follows a power-law distribution of energies −p N(γ) = N0γ (where γ ∈ [1, ∞)) where the spectral index is given by α = −(p − 1)/2.

We can integrate this distribution to obtain the total number density of electrons ne =

R ∞ −p N0 −p 1 N0γ dγ = p−1 , such that N(γ) = (p − 1)neγ . The total power per unit volume per unit frequency of emitted synchrotron radiation (see, e.g., Rybicki and Lightman, 1986; Ghisellini, 2013) from this population of electrons will be

p 2 − 2 σT cneB ν ν = 3/2 f(p) (3.3) 32π νL νL 2.25 f(p) ≈ 3p/2(p − 1) + 0.105 (3.4) p2.2 eB νL = (3.5) 2πmec where we have assumed the magnetic field lines in the SEK are tangled, i.e. we have integrated the pitch angle between the path of the electron and the magnetic field over the possible solid angle that synchrotron radiation can be emitted. At these wavelengths, we assume each SEK is optically thin. It follows that the total luminosity emitted in near- K L = 1 πD3 ν infrared synchrotron emission at s band and for each SEK is Ks 6 SEK Ks . K L = 7.2 × 1032 −1 The observed s band luminosities of SEK 1 and SEK 2 are Ks erg s 32 −1 2 and 7.6 × 10 erg s , respectively, defined as L = 4πd νFν using the flux density from the 2002 epoch and assuming a distance of d = 3.4 kpc. The angular size of both SEKs is 64

00 17 approximately θ = 3 , corresponding to a size of DSEK = 1.5 × 10 cm at 3.4 kpc. For p = 2.42, 2.46 (corresponding to α = −0.71, −0.73 for SEKs 1 and 2, respectively) the electron densities toward SEK 1 and SEK 2 are 15, 000 cm−3 and 1, 000 cm−3.

3.4 Discussion

The magnetic field strength in SEKs 1 and 2 is clearly enhanced in a small knot compared to values calculated for acceleration sites distributed across the entire remnant from gamma-ray fluxes. Moreover, the location of the knots places them approximately 12300 from the expansion center of the remnant given their location in Figure 3.2 and the expansion center calculated in Thorstensen et al.(2001). It is therefore reasonable to assume that these knots are behind the location of the forward shock at this position angle (e.g., in Fesen et al., 2001; Gotthelf et al., 2001). The presence of compact knots of synchrotron emission closely matches the emission model derived in Reynolds(1998) where the magnetic field can be compressed and amplified downstream of the forward shock. For a compression ratio r = 4 (i.e., B1 = rB2), which is typical for the adiabatic strong-shock limit with γ = 5/3, the magnetic field strengths calculated for SEKs 1 and 2 (B1) imply an upstream magnetic field strength

(B2) of 0.33 − 1.5 mG, which is of the same order, but somewhat larger than the values inferred from gamma-ray measurements. This rough agreement supports the overall validity of the Reynolds(1998) emission model. Additional observations and analysis are needed to determine if the field values above the predictions can be explained within this model. Otherwise, some additional mechanism may be amplifying the magnetic field in this region, such as turbulent amplification brought on by Rayleigh-Taylor instabilities in the shock (as in, e.g., Jun and Jones, 1999). Moreover, it is curious that SEK 2 is fading so quickly (Figure 3.2) and that this behavior implies a large magnetic field strength and low electron density compared to SEK 1. There must be a sharp change in the properties of the remnant over a ∼ 0.05 pc distance such that, despite being entrained in the other synchrotron-emitting ejecta, the SEKs have their own distinct properties. Along with the additional magnetic field amplification, this behavior also supports a scenario in which a strong shock interaction and instabilities in 65 the ejecta have compressed knots of material. Indeed, the presence of a recent shock interaction to the southwest of Cas A would also explain enhancement in hard X-ray and radio emission (Anderson and Rudnick, 1995; Grefenstette et al., 2015) as well as the presence of shocked molecular gas (Kilpatrick et al., 2014) toward this part of the remnant.

3.5 Conclusions

We have presented multi-epoch Ks band imaging of Cas A demonstrating that two knots of emission to the southwest of the remnant appear entrained in the ejecta. Using radio and Spitzer IRAC data from the literature, we have argued that the Ks band emission from these knots is dominated by synchrotron radiation. Our main conclusions from our analysis of these knots are:

1. The two synchrotron-emitting knots are fast moving (∼ 5000 km s−1) and appear to have magnetic ﬁelds of 1.3 and 5.8 mG.

2. These ﬁelds are an order of magnitude higher than deduced by gamma ray measurements that average electron acceleration sites distributed over the entire remnant.

3. This behavior appears to be consistent with the compression and ampliﬁcation of the magnetic ﬁeld downstream of the forward shock as proposed by Reynolds(1998). 66

CHAPTER 4

Interaction Between Cassiopeia A and Nearby Molecular Clouds

We present spectroscopy of the supernova remnant Cassiopeia A (Cas A) observed at infrared wavelengths from 10 − 40 µm with the Spitzer Space Telescope and at millimeter wavelengths in 12CO and 13CO J = 2 − 1 (230 and 220 GHz) with the Heinrich Hertz Sub- millimeter Telescope. The IR spectra demonstrate high-velocity features toward a molecular cloud coincident with a region of bright radio continuum emission along the northern shock front of Cas A. The millimeter observations indicate that CO emission is broadened by a factor of two in some clouds toward Cas A, particularly to the south and west. We believe that these features trace interactions between the Cas A shock front and nearby molecular clouds. In addition, some of the molecular clouds that exhibit broadening in CO lie 1 − 20 away from the furthest extent of the supernova remnant shock front. We propose that this material may be accelerated by ejecta with velocity significantly larger than the observed free-expansion velocity of the Cas A shock front. These observations may trace cloud interactions with fast-moving outflows such as the bipolar outflow along the southwest to northeast axis of the Cas A supernova remnant, as well as fast-moving knots seen emerging in other directions.

4.1 Introduction

Supernova explosions are traditionally described in terms of a spherical shock that expands into the homogeneous interstellar medium. This model ignores the kinematics of the actual supernova expansion where the boundary between the shock front and the interstellar medium is composed of material at varying temperatures and densities. An important consequence of this more complex model of supernova physics is the evolution of shock fronts that encounter an interstellar cloud of size comparable to the supernova remnant (SNR). Internal flows of the SNR are expected to become highly turbulent where the 67 shock front interacts with a cloud (Jun and Jones, 1999). Any spectral component should be velocity-broadened as a result of increased turbulence from a shock front/cloud interaction. Magneto-hydrodynamic models indicate that the reverse shock behind the surface of interaction between the two media should be a site for particle acceleration and magnetic field amplification (Jun and Jones, 1999). Consequently, a strong synchrotron component may be present between supernovae forward and reverse shocks. The SNR Cassiopeia A (Cas A) is a nearby (3.4 kpc) and young (∼ 350 yr) example of a Type IIb supernova (Krause et al., 2008) well-suited to observing these effects. The proximity of the SNR and its inherent brightness make it ideal for studying the details of supernova physics including nucleosynthesis, light echoes, dust physics, and cosmic rays. There is an extensive history of molecular spectroscopy of gas in the general direction of Cas A as well as the study of interactions between its ejecta and surrounding material (e.g., Weinreb et al., 1963; Barrett, 1964; Bieging and Crutcher, 1986; Heiles and Stevens, 1986; Keohane et al., 1996; Krause et al., 2004). Among this work, observations of carbon monoxide lines indicate that some of the gas toward the SNR is relatively warm (TK ∼ 20 K) (Wilson et al., 1993). However, whether there is an on-going interaction between the Cas A shock front and nearby molecular gas is ambiguous. Although CO line widths to the east and west may demonstrate some broadening (Liszt and Lucas, 1999), the 1720 MHz OH maser emission that is often seen in shocks is missing (Frail and Mitchell, 1998). Perhaps the broadened CO results from an interaction that occurred recently and has not yet had time to develop more complex structure. Given the youth of the SNR, the effects of interactions around Cas A in terms of warm molecular gas and broadened lines may be subtle. Nonetheless, there are some enticing hints at an interaction. The relatively warm temperature of the gas surrounding the SNR suggests an interaction between the shock front and molecular cloud along the western edge of Cas A. In addition, Arendt et al.(2014) find that the dust at the outer edge of the remnant has a composition consistent with that of interstellar dust. Fesen et al.(2011) identify variable knots of emission at a similar distance from the SNR center that they ascribe to interactions with ambient material. Cor- relation between radio synchrotron emission and X-ray emission, especially toward the 68 western edge of the remnant, suggest that the ejecta are moving into a denser medium while steep radio indices suggest relativistic particle acceleration (Anderson and Rudnick, 1996; Keohane et al., 1996). These phenomena are consistent with an interaction between the fast-moving ejecta and a dense molecular cloud. There is also a molecular cloud at the north shock front, with bright radio emission between this cloud and the Cas A reverse shock (Figure 4.1; Hines et al., 2004). The absence of any X-ray silicon emission or optical line emission at this position indicates that the radio emission does not result from an enhancement in the molecular gas density or from any obvious characteristic of the SNR. Therefore, we infer that this radio emission may be a tracer for particle acceleration behind the point of interaction of the shock front and a molecular cloud (Jun and Jones, 1999). Near-infrared Ks emission has also been demonstrated to be dominated by a synchrotron continuum component (Rho et al., 2003). The

Ks emission feature along the northern shock front (Rho et al., 2003), roughly coinciding with the enhanced radio emission, supports the possibility that we are seeing a SNR/MC interface. Enhancements in Ks emission to the west and south indicate some other regions of possible interaction. To search systematically for such interactions, we present radio and infrared spectroscopy of gas within several arcminutes of Cas A. In particular, we looked for line broadening in the J = 2 − 1 transitions of 12CO and 13CO. We also searched for signs of an interaction in spectra taken from the Spitzer Infrared Spectrograph (IRS) along two tracks in declination (Figure 4.1) across the northern shock front of Cas A. These observations were performed using the high-resolution module aboard IRS, and given the integration time, are the deepest mid-infrared spectra of this SNR available.

4.2 Observations

We mapped the vicinity of the Cas A SNR with the Heinrich Hertz Submillimeter Tele- scope (SMT) using the ALMA-type sideband separating 1.3 mm receiver on 2011 January 30. We observed the J = 2 − 1 transitions of 12CO and 13CO at 230 GHz and 220 GHz, simultaneously, conﬁguring the receiver with 12CO J = 2 − 1 in the upper sideband and 13CO J = 2 − 1 in the lower sideband and using an IF of 5000 MHz. The two sidebands 69

Figure 4.1: VLA 20 cm map of Cas A (Anderson et al., 1991) with positions of scan 1 (east) and scan 2 (west) overlaid as rectangular boxes (pink). Each of these boxes represents 16 pointings of the Spitzer IRS, with the southernmost pointing in each scan (i.e. S1, P1 and S2, P1) corresponding to the first pointing in each scan. Each successive pointing in each scan is spaced 200 to the north of the previous pointing (i.e. S1, P2 is 200 due north of S1, P1). The large circle (pink) represents the position of a molecular cloud identified in previous literature (Hines et al., 2004). The ellipse at the south end of the two shock tracks outlines the region of bright radio continuum emission that suggests a shock front/cloud interaction. 70 were configured with independent mixers receiving both the horizontal and vertical po- larizations. We used two spectrometers in parallel, providing 1 MHz and 250 kHz filters. The quality of the sky was good, with opacity around a few percent for the duration of the mapping. The system noise temperature varied with source elevation over the course of the observation, with values of ∼ 200 K over the majority of the map, but values as large as 269 K over the last several rows. The map was obtained on-the-fly (OTF) and covered a 100 × 100 region centered on the SNR. The rows were spaced 1000 in declination and the area was scanned boustrophe- donically, with a line-free reference position observed after every other row. The final map has 61 × 65 pixels, each 1000 × 1000 with a beam width of 3300 FWHM. Our reference position strategy provided adequately flat spectral baselines, although there was residual baseline ripple of ∼ 0.1 K peak-to-peak in some areas. The SMT data were processed using CLASS.1 A spectrum was extracted from each pixel and linear baseline subtraction was performed followed by sky subtraction using a spectrum derived from our reference. The spectra were then gridded into a datacube and converted into MIRIAD format (Sault et al., 1995) for analysis. The infrared data are high-resolution 10 − 40 µm spectra obtained in 2008 September (AORID 50322) using the Spitzer IRS. Two scans were constructed across the northern edge of Cas A consisting of 16 individual spectra in the short-high module integrating for 120 s per pointing. Long-high spectra were obtained at the same pointings integrating for 60 s per spectrum. The same parameters were used for each map, with pointing positions spaced approximately 200 apart (Figure 4.1). All IRS data were reduced using SMART (Higdon et al., 2004). The two spectral scans were broken down into 16 individual spectra each for 32 total spectra at different pointings. Background subtraction was performed from sky images taken during the observation. Finally, the spectra were extracted using the full-aperture algorithm for extended sources.

1http://www.iram.fr/IRAMFR/GILDAS 71

2.0

1.5

1.0 (K) * r T 0.5

0.0

-0.5 -60 -50 -40 -30 -20 -1 VLSR (km s )

Figure 4.2: 12CO J = 2 − 1, 250 kHz ﬁlter bank spectrum averaged across the 100 × 100 mapped region.

4.3 CO Line Emission Around Cas A

4.3.1 Kinematics of CO Line Emission as Observed in J = 2 − 1 Spectra

A spectrum of the 12CO J = 2 − 1 transition observed using the 250 kHz ﬁlters and averaged over the entire 100 ×100 mapped region is shown in Figure 4.2. The high-velocity gas is differentiated into three main components as shown in the CO J = 1−0 maps of Liszt and Lucas(1999). The most highly blue-shifted gas ( −50 to −45 km s−1) is concentrated in ﬁlamentary structure to the south and southeast of the remnant, the intermediate velocity gas (−45 to −39 km s−1) appears in the west and north edges, and the most red-shifted gas (−39 to −34 km s−1) is located in the southeast and west. These features can be observed in the channel maps in Figure 4.3 where the 12CO emission at 250 kHz resolution is displayed from −33 to −50 km s−1. Between −33 and −41 km s−1, the brightest emission appears west and south of the SNR, with a separate cloud that extends to regions on the order of 1 − 20 away in projected radius along the line of sight of the remnant. 72

Figure 4.3: Channel maps of 12CO J = 2 − 1 in the 250 kHz ﬁlter with VLA 20 cm emission contour overlaid. The velocity of each map is given in the upper left hand corner. The circle depicts the location of an over-density of molecular gas located along the northern edge of the remnant. Note that the over-densities of molecular gas to the west, south and east of the SNR correspond to peaks in the contours of 20 cm emission. 73

Figure 4.4: The integrated intensity (K km s−1) of 12CO J=2-1, 250 kHz ﬁlter bank spectrum masked for regions where the FWHM of 12CO J=2-1 lines exceeds 6 km s−1, which is a ﬁducial value we use to denote “broad-line regions.” The VLA 20 cm map of Cas A is overlaid to demonstrate the locations of various broad-line regions relative to the SNR. The only four regions in the entire map dominated by broad lines (A, B, C, D) are indicated by shaded pixels within thick-lined boxes, while two empty, thin-lined boxes (E, F) are randomly selected comparison regions dominated by narrow lines. The averaged spectra from each of these regions are plotted in Figure 4.5. 74

4.3.2 Possible Cloud/SNR Interactions

The frames at −40.86 and −41.51 km s−1 in Figure 4.3 show the molecular cloud at the north edge of Cas A described previously (Hines et al., 2004; Liszt and Lucas, 1999). The cloud is circled in white and lies near the two tracks of our Spitzer IRS observations outlined in Figure 4.1. The cloud has a mass of 650 M (Hines et al., 2004) and extends over a region 1.50 × 0.80 on the sky. Figure 4.1 also shows the enhanced zone of radio emission at the interface between the SNR and this cloud. This over-density and the radio feature so close to the edge of the remnant suggest an interaction between the two media. In general, the brightest CO-emitting regions along the line of sight to Cas A correspond well with the brightest emission in the VLA 20 cm map of Anderson et al.(1991) possibly indicating additional interactions between the shock front and molecular clouds.

4.3.3 CO Line Broadening

At points of interaction where shocks from the supernova ejecta have propagated through the molecular gas, we would expect to see line broadening due to enhanced turbulence. To evaluate the extent of that broadening in molecular clouds toward Cas A, we developed an algorithm to search the entire mapped area and velocity range in our 250 kHz filter maps for pixels with large line widths. The algorithm was optimized not to confuse multiple narrow lines with a single broad one. To this end, we identified pixels in our datacube with CO emission exceeding the noise baseline by at least 5-σ and fit a Gaussian profile with a linear baseline to the line profile. We iteratively increased the velocity range of interest around the velocity centroid of our fit until the fit either diverged (i.e., the reduced χ2 of the fit increased by a factor of two from the previous fit) or the fit parameters varied less than 10% from the previous one. In this way, we identified pixels that exhibit CO lines with widths exceeding 6 km s−1, our fiducial value for “broad-line regions.” We tested the sensitivity of the pixels selected on the threshold velocity width and found that the pixels identified as “broad-line regions” were not a strong function of the fiducial velocity-width above ∼ 3 km s−1. That is, these pixels stand out as qualitatively distinct components of the cloud, with significantly higher line widths than the molecular emission in adjacent 75

Table 4.1. Gaussian ﬁt values for CO lines

−1 −1 Line Peak (K) vc (km s ) FWHM (km s ) rms (K)

Ablue 3.12 −46.9 6.48 0.440

Ared 0.73 −37.2 3.42 0.083

Bblue 2.09 −46.5 6.45 0.158

Bred 0.81 −36.4 2.74 0.061

Cblue 1.40 −46.5 3.19 0.051

Cred 3.15 −37.7 6.19 0.116

Dblue 1.40 −46.6 3.01 0.051

Dred 3.55 −37.7 6.66 0.097

Eblue 3.99 −46.8 2.71 -0.009

Ered 2.51 −36.5 2.78 0.022

Fblue 2.40 −48.5 2.50 0.005

Fred 0.42 −37.3 2.69 0.027

pixels. Figure 4.4 locates the regions where the integrated CO intensity has strong wide components. For comparison, we overlaid the VLA 20 cm continuum map to demonstrate where the broad-line regions lie in relation to Cas A. In order to further examine the association between Cas A and these clouds, we have highlighted six regions in Figure 4.4 using boxes (A, B, C, D, E, F). The first four of these regions contain pixels that our algorithm identified as having “broad-line regions” while the last two contain only narrow lines in a similar velocity range for comparison. The average spectra of the pixels in each of the six boxes are plotted in Figure 4.5. The blue lines in panels A and B and the red lines in panels C and D do indeed appear broad, with FWHM velocities exceeding 6 km s−1 (see Table 1 for values from our Gaussian fits to these lines). Furthermore, these lines appear to demonstrate the “broad-wing” characteristic that is a signature of shocked molecular emission toward SNRs (e.g., Denoyer, 1979b,a; Burton et al., 1988). By comparison, the molecular emission lines in panels E and F appear narrower, which we approximate as Gaussian profiles with FWHMs less than 3 km s−1. We estimated the masses of the clouds located west, south, and southeast of the SNR, and the momentum necessary to generate the broad lines in these regions. We set a threshold of 6 km s−1 in the second moment of 12CO emission at each pixel to define the tur- 76 bulent regions. For clouds with observed sizes on the order of a few parsecs, this value is slightly more than twice as large as 12CO line widths typically observed in quiescent galactic molecular clouds (Heyer and Brunt, 2004). Using the ratio of 12CO to 13CO emission, we calculated the optical depth toward each pixel in three of the molecular clouds and used this value to determine the column density of molecular hydrogen at each position. From these values, we found cloud masses of approximately 660, 530, and 340 M for the turbulent regions of the western, southern, and southeastern clouds, respectively. The amount of momentum necessary to generate these line widths is then approximately, −1 −1 McloudσCO or 4000, 3200, 2000 M km s assuming a minimum line width of 6 km s for each cloud. Analysis of infrared transitions of CO toward Cas A suggests that astrochemical processes form molecular gas in the ejecta-rich reverse shock of the remnant (Rho et al., 2009a, 2012). Indeed, Herschel detections of broad, high-J rotational CO lines in the far infrared toward the Cas A remnant suggest that molecular gas has reformed recently in bright knots(Wallstrom¨ et al., 2013), reminiscent of optical fast moving knots (Fesen, 2001). However, the masses we detect in the western, southern, and southeastern clouds greatly exceed the mass of the progenitor star, so we infer that the molecular gas cannot originate from the remnant itself. Some of the broad-line regions (A, B, C) we have highlighted lie not merely toward the edge of the Cas A remnant, but toward radio continuum enhancements along the western and southern edges of the remnant, possibly due to an interaction between molecular gas and the supernova ejecta. The western position matches the filamentary structure observed in previous work (Liszt and Lucas, 1999) and lies where the molecular cloud morphology is consistent with an interaction as previously suggested by Keohane et al. (1996). However, one of the broad line regions (D) appears more than 20 from the forward shock, as demarcated by the edge of the remnant, which suggests an interaction with very fast-moving ejecta. We will explore this behavior in Section 4.5.1. 77

5 A B 4

3 (K) * r 2 T

0 5 C D 4

3 (K) * r 2 T

0 5 E F 4

3 (K) * r 2 T

-60 -55 -50 -45 -40 -35 -60 -55 -50 -45 -40 -35 -30 Velocity (km s-1) Velocity (km s-1)

Figure 4.5: 12CO J=2-1, 250 kHz filter bank spectrum averaged across the six roughly 5000 × 5000 regions denoted in Figure 4.4. Note that each spectrum contains multiple CO line components, although we have fitted Gaussian curves to the two brightest components in each panel. In each panel, we call the feature with velocity more negative than −43 km s−1 the “blue” component and the feature with velocity more positive than −40 km s−1 the “red” component, denoted by the blue and red dotted curves. The blue components in panels A and B and the red components in panels C and D are broadened, that is, they have a FWHM exceeding 6 km s−1 (line fits given in Table 1). We compare these lines to the narrower CO lines in panels E and F, which are more typical of CO in the ISM. 78 [S IV] [Ne II] [Ne V] [Ne III] [S III] [O IV]

Figure 4.6: Example of high-resolution shock front spectrum from the Spitzer IRS high- resolution data. This pointing represents the southernmost pointing along scan 2 (west in ﬁg. 1). The spectrum is separated into the short-wavelength (SH - upper panel) and long-wavelength (LH - lower panel) modules. All spectral features observed in each scan are identiﬁed. 79

4.4 Analysis of Cas A Infrared Spectroscopy

4.4.1 Spectral Features

A typical infrared spectrum with spectral line identifications, taken from one of our pointings toward the northern Cas A shock front, is shown in Figure 4.6. Both the emission lines and the dust continuum emission become less bright as the spectra are taken from pointings that extend further north of the remnant. In our example spectrum, the continuum is relatively faint at shorter wavelengths but peaks strongly around 21 µm and slowly decreases in brightness at longer wavelengths. This 21 µm peak observed by Lagage et al. (1996); Arendt et al.(1999); and Rho et al.(2008) has been associated with freshly formed silicate and iron dust in the ejecta. In scan 1, emission features are apparent at 10.5, 12.8 (doublet), 14.3, 18.7, and 25.9 µm. Similarly, scan 2 presents several strong features in the pointings furthest south. These were identified at 10.5, 12.8 (doublet), 14.3, 15.6, 18.7, and 25.9 µm. Virtually all the features identified in scans 1 and 2 correspond with previously identified fine-structure lines based on low-resolution IRS spectra (Rho et al., 2008; Smith et al., 2009; DeLaney et al.,

2010). In particular, [S IV] (10.51 µm), [Ne II] (12.81 µm), [Ne V] (14.32 µm), [Ne III]

(15.56 µm), [S III] (18.71 µm), and [O IV]+[Fe II] (25.94 µm) are commonly observed features in the ejecta toward this remnant. Apart from these lines, there is also a second emission line near 12.7 µm in both scans 1 and 2. Given the high-velocity components in other spectral features, we consider in the appendix whether this line can originate from neon in high-velocity gas.

The bright features (e.g., [O IV], [S III], [S IV]) also show considerable velocity structure. We have plotted these features in velocity-space for a spectrum composed of the two southernmost pointings along scan 2 averaged together (Figure 4.7). These averaged spectra emphasize the high-velocity features, which we have marked with vertical lines.

The remainder of our analysis concentrates on the [O IV] line; the other lines are discussed in the appendix. 80

1.8 2σ 1.6 [O IV] 1.4 1.2 1.0 0.8

Flux Density (Jy) 0.6 0.22 2σ 0.20 [S III] 0.18 0.16 0.14 0.12 0.10 Flux Density (Jy) 0.08

0.06 2σ [S IV] 0.04

0.02 Flux Density (Jy) 0.030 0.025 [Ne II] σ 0.020 2 0.015 0.010

Flux Density (Jy) 0.005 0.035 0.030 [Ne V] σ 0.025 2 0.020 0.015

Flux Density (Jy) 0.010 0.040 0.035 2σ 0.030 [Ne III] 0.025 Flux Density (Jy) −10000 −5000 0 5000 10000 Radial Velocity (km s−1)

Figure 4.7: Spectra from the six ([O IV], [S III], [S IV], [Ne II], [Ne V], [Ne III]) fine structure lines identified in Figure 4.6. These spectra represent the two southernmost pointings of scan 2 averaged to emphasize the presence of high-velocity features. A sample 2-σ error bar is shown in the upper-left hand corner to compare with the strength of the identified features. Five dotted, vertical lines mark some recurrent high-velocity features (at −2300, −200, 1100, 1700, 3200 km s−1). 81

2.0 2σ 1.8 S2,P1 1.6 1.4 1.2 1.0 0.8 Flux Density (Jy) 0.6 σ 1.4 2 S2,P2 1.2 1.0 0.8

Flux Density (Jy) 0.6 1.2 2σ 1.1 S2,P3 1.0 0.9 0.8 0.7 Flux Density (Jy)

1.0 2σ S2,P4

0.9

0.8 Flux Density (Jy) 0.9 σ 0.8 2 S1,P1 0.7 0.6 0.5

Flux Density (Jy) 0.4 0.9

0.8 2σ S1,P2 0.7 0.6 0.5 Flux Density (Jy) −10000 −5000 0 5000 10000 Radial Velocity (km s−1)

Figure 4.8: Identiﬁed [O IV] (25.94 µm) lines in scans 1 and 2 (S1 and S2) at the southernmost pointings in each scan (P1-2 in S1, P1-4 in S2). These lines have been plotted versus relative velocity to demonstrate the presence of high-velocity components in the spectra. A sample error bar is shown in the upper-left hand corner to compare with the strength of the identiﬁed features. Five dotted, vertical lines mark some recurrent high-velocity features (at −2300, −200, 1100, 1700, 3200 km s−1). 82

4.4.2 [O IV] 25.94 µm line

Observations

The [O IV]+[Fe II] feature around 25.94 µm is consistently the brightest fine-structure line detected toward the Cas A remnant. It is brightest around the edge of the remnant with significant emission toward the center relative to other lines (Rho et al., 2008). From the co-added spectra in both scans, it is possible to separate high-velocity features of the supernova outflow in the resolved [O IV] line profiles (around 25.94 µm). At pointing 1, there are five features centered at 25.63, 25.74, 25.86, 26.14, and 26.23 µm (Figure 4.8).

Assuming most of the ﬂux is due to [O IV] emission, these wavelengths correspond to radial velocities of −3580, −2300, −920, 2310, and 3350 km s−1, in agreement with the range of previous velocities observed within the remnant (−4000 to 6000 km s−1 (DeLaney et al., 2010)). Along scan 2, there appear to be three distinct features in the stacked spectrum at wavelengths 25.76, 25.92, 26.14 µm. For [O IV] emission, these wavelengths correspond to radial velocities of −2080, −230, and 2310 km s−1. At the high signal to noise on the central features in the scan 1 and 2 spectra, the wavelength difference between them (1.5 full IRS line widths) is signiﬁcant. Given the modest spectral resolution (R ∼ 600) and signal to noise, there is no convincing case that any of the features coincide between scans 1 and 2 (see Figure 4.8). In any case, the two spectra demonstrate that we observe ejecta over a velocity range of at least 5000 km s−1 in a relatively small region of the northern shock front (see Fig- ure 4.1). To demonstrate the behavior of [O IV] over the spatial region for which we have high-resolution spectra, we have plotted spectra for all 32 pointings (see description in

Section 4.2) centered around the [O IV]+[Fe II] line in Figure 4.9. A linear baseline was subtracted and ﬂux was added to each spectrum such that the line with the brightest baseline corresponds to the pointing furthest south (i.e. P1) and each subsequent spectrum stepping down in ﬂux represents the next pointing spaced 200 further north (i.e. P2, P3, etc.). We note in particular that the bright [O IV] feature disappears over an angular distance of about 600 for S1 and 1000 for S2, which is comparable to the IRS slit width of 83

1100 (Houck et al., 2004), and the relative emission from each high-velocity component remains the same at each pointing within the signal-to-noise. That is, the 5000 km s−1 velocity range persists all the way to a sharp cutoff in the line emission with increasing radius.

Is the velocity spread due to turbulence in the supernova ejecta?

The velocity spread in the infrared ﬁne structure lines is similar to that observed over most of the face of the SNR, where it is attributed to two causes: 1) turbulence in the explosion ejecta; and 2) Doppler shifts from systematic motions associated with radial expansion of the remnant. DeLaney et al.(2010) show that interpreting the extreme red and blue line components as due to Doppler shifts associated with line emission from opposite sides of an expanding sphere can yield a three-dimensional picture of the remnant. In this picture, the absolute Doppler shifts should tend to a minimum as one approaches the extreme edge of the remnant as projected onto the sky, since all of the systematic motions associated with the spherical expansion will lie on the plane of the sky. We will analyze this behavior quantitatively in the following section. Does the broad velocity spread in our observations, obtained at the extreme edge of the remnant, have the same causes? Attributing it to a combination of turbulence in the ejecta and the systematic expansion of the remnant has a number of implications. It might indicate that the lines we observe do not come from the extreme edge of the remnant, although we show in the next section that this explanation is not likely. It might indicate that the spherical geometry is not strictly correct at the edge of the remnant. Broadened lines around the extreme edge of the remnant might arise if the remnant were like a cylinder which we see end-on, but such a model seems physically untenable. As an alternative, in the following discussion we explore whether the turbulence is established in the area of our observations, which by Newton’s second law requires that there be mass for the ejecta to react against to modify their outward expansion. 84

3.0 3.0

S1 S2

2.5 2.5

2.0 2.0

P1 1.5 P2 1.5 P1 P2 P3 P3 P4 P4 P5 P5

Flux Density (Jy) P6 Flux Density (Jy) P6 1.0 P7 1.0 P7 P8 P8 P9 P9 P10 P10 P11 P11 0.5 P12 0.5 P12 P13 P13 P14 P14 P15 P15 P16 P16 0.0 0.0

25.6 25.8 26.0 26.2 26.4 26.6 25.6 25.8 26.0 26.2 26.4 26.6 Wavelength (µm) Wavelength (µm)

Figure 4.9: [O IV] 25.9 µm spectra depicted across the shock front along two scans, given as S1 (left) and S2 (right), to the north of Cassiopeia A. The spectra are depicted with the southernmost pointing (closest to Cas A center) at the highest flux density with spectra at pointings successively further north at lower flux density. These pointings start at a position approximately 30 north of the center of Cas A and step north in increments of 200. Note that the relative flux of high-velocity features surrounding the [O IV] line (dotted lines at −3580, −2300, −110, 1140, and 1640 km s−1 in S1 and −1850, −360, and 1140 for S2) remain roughly the same at each pointing and the feature disappears completely after roughly 3 pointings, or 6”, for S1 and 5 pointings, or 10”, for S2, which is comparable to the slit width of Spitzer IRS (1100). This feature suggests that there is a sharp cut off where the [O IV] emission, and thus the supernova ejecta, rapidly becomes less dense, consistent with a contact discontinuity between the ejecta and interstellar gas. 85

Does Radial Expansion Account for the Velocity Spread?

We now test whether the large spread in radial velocities discussed above can be explained in the context of pure radial expansion of the SNR. Assuming that the observed emission originates from ejecta in a thin shell toward the Bright Ring of Cas A, the observed Doppler

velocity (vD) of ejecta in that shell has been modeled using [Ar II] (6.99 µm) in DeLaney et al.(2010) using a three-parameter model,

q 2 2 w = |vD − vc| = (vc − vm) − (rP /S) (4.1)

where rP is the projected radius in arcseconds, vc represents the center of the distribution

of Doppler velocities, vm represents the lowest Doppler velocity for ejecta in the shell, S is a scale factor that relates Doppler velocity to projected radius, and w is the spread in

velocities around vc. We note that as rP increases, vD approaches vc and the line width tends to zero. This behavior is further expressed in the derivative of equation (1),

2r dr dw = P P (4.2) 2q 2 2 S (vc − vm) − (rP /S)

For the parameters that DeLaney et al.(2010) derive toward the Bright Ring of [Ar II] −1 −1 00 00 −1 (vc = 859 ± 100 km s , vm = −4077 ± 200 km s , S = 0.022 ± 0.002 per km s ), 00 for projected radii close to the furthest extent of the Bright Ring of [Ar II](rP ∼ 100 ) and 00 assuming incremental steps away from this position of drP ∼ 2 , we ﬁnd that dvD ∼ 200 km s−1. That is, for each 200 step in projected radius away from the center of expansion, w should decrease by about 400 km s−1, or about 2000 km s−1 over all ﬁve pointings in which we see the lines. In each spectral scan, the line components remain at the same relative strength until the ejecta are no longer present in our spectra. The observed spread −1 in Doppler velocities of the [O IV] line in Figure 4.9, roughly a few 1000 km s , cannot be reconciled with the 12900 projected radius of the spectral pointings from the observed center of expansion. We infer that there must be some additional source of energy that drives the observed velocity spread.

Finally, any explanation of the high-velocity Doppler components of the [O IV] emission must account for a range in velocities that is significantly in excess of the velocity 86 spread typically observed in Cas A ejecta at a single position (Figure 9 in DeLaney et al., 2010). This aspect of ejecta in the Cas A remnant has been confirmed through observation of optical ejecta knots at projected radii 2 − 30 and with Doppler velocities extending on the order of a few 100 km s−1 per knot (Fesen, 2001). Isensee et al.(2012) find that, to

ﬁrst order, the FWHM of the [O IV] line averaged over the entire remnant in low-resolution −1 Spitzer IRS spectra is ∼ 1000 km s . We contrast these observations with the [O IV] emission at one position, at 13000 projected radius from the expansion center of Cas A where high-velocity emission appears to extend 5000 km s−1. Any explanation of the ejecta at this position must account for such an anomalously large spread in Doppler velocity.

Fast Moving Knots and the Velocity Spread

We now consider the possibility that the emission results from a superposition of multiple fast-moving knots. Optical studies of regions beyond the bright, nebular shell surrounding Cas A have revealed dozens of emission-line knots at projected radii ∼ 130 − 210 arcseconds from the estimated center of expansion (Fesen et al., 2001; Hammell and Fesen, 2008). These observations cover the radius at which we observe highly Doppler shifted h m s ◦ 0 00 00 [O IV]+[Fe II] emission at approximately 23 23 26 , 58 50 57 , or about 130 north of the center of expansion as determined in Thorstensen et al.(2001) at a position angle of roughly 354◦. However, catalogs of fast-moving knots do not show any in the vicinity of our observations (Hammell and Fesen, 2008). Fast-moving knots tend to have a limited spread in velocity and exhibit structure on relatively small scales, with typical knot sizes between 0.2” and 0.4” (Fesen et al., 2001; DeLaney et al., 2010). We observe multiple high-velocity Doppler components in each of our pointings, which would indicate that several fast-moving knots are required to account for the line structure. However, each pointing (through, roughly, P5 in Figure 4.9) exhibits emission at the same velocities. This observation suggests that we observe a single line-emitting region with a broad range of velocities that sharply cuts off radially within approximately one slit width. The presence of multiple FMKs that could mimic a single emitting region within spatial scales of 1000 is unlikely in this part of the Cas A remnant. The majority of known 87

FMKs concentrate in the region along the northeast to southwest axis of the remnant where the “jet” is observed. Along the northern rim, ejecta knots are significantly less common. Hammell and Fesen(2008, see Figure 10, therein) find that there are 36 knots with projected radii between 13000 and 14000, or approximately 15 per square arcminute. In the on-sky area covered by the IRS slit (∼ 250 square arcseconds in Long-High mode), we would expect to observe approximately one FMK, or two FMKs in two separate pointings. For a Poisson random distribution of FMKs on-sky, we find that the probability of observing 8 FMKs in our two pointings (to match the eight velocity components in [OIV]) is roughly 0.1%. Since no knots are observed at the slit positions, this low probability is confirmed empirically.

Interaction-Induced Turbulence and the Velocity Spread

We ﬁnd that a scenario involving turbulent acceleration of the ejecta due to an interaction with interstellar material matches the observed spectra much more fully. The large spread in Doppler velocities over this small region, the lack of any known ejecta knots in this region, and the placement of this region beyond the observed Bright Ring of optical emission support the interpretation of ejecta located toward the forward shock interacting with interstellar gas. However, for a more in-depth analysis of the [O IV]+[Fe II] emission, we need to extend the study beyond the limited region covered by our high-resolution spectra (see Figure 4.1).

4.4.3 Analysis of Low-Resolution [O IV] Emission

We use archival low-resolution mid-infrared spectra that cover the remnant in its entirety to observe other regions along the shock front. Based on the range of velocities in [O IV] in our Spitzer IRS measurements, we expect that low-resolution data will exhibit a broadened spectral line indicating the spatial extent of Doppler shifted [O IV] emission in the SNR. We therefore obtained low-resolution Spitzer IRS spectra of Cas A (Rho et al., 2008; Smith et al., 2009; DeLaney et al., 2010) from the Spitzer Heritage Archive (AORID 3310) to compare the line width of [O IV] across the SNR shock front. 88

2.5

2.0 High-resolution

1.5

1.0 Flux Density (Jy) 0.5 10

8 Low-resolution PN 6 4 2 Flux Density (Jy) 0 1.4 Convolved High−resolution 1.2 1.0 0.8

Flux Density (Jy) 0.6 25.2 25.4 25.6 25.8 26.0 26.2 26.4 26.6 Wavelength (µm)

Figure 4.10: (Top) Sample spectrum from scan 2 of high-resolution [O IV] (Middle) Sam- ple low-resolution spectra taken from a Spitzer IRS image of PNG 002.7-04.8 (Bottom) The top two spectra convolved.

As a test, we convolved our own high-resolution data with a spectral response function for the IRS. We derived this function by ﬁtting a Gaussian curve to a narrow [O IV] emission feature observed in a planetary nebula (PNG 002.7-04.8) using the low-resolution module on Spitzer IRS and obtained from the Spitzer Heritage Archive (Figure 4.10). The convolution of our data to a lower spectral resolution effectively smooths and broadens the separate [O IV] features to a single line component. The IRS low-resolution Cas A spectra will resemble this convolved spectrum but exhibit larger widths where shock front/cloud interactions contribute to high-velocity [O IV] components.

From the convolved high-resolution spectrum, we evaluated the width of the [O IV] 89

Table 4.2. Gaussian ﬁt values for convolved high-resolution [O IV] spectra

Position C0 C2 Integrated width Velocity Width (Jy) (µm) (µm) (km s−1)

shock1-pos01 0.19 0.20 0.21 2400 shock1-pos02 0.14 0.24 0.25 2900 shock2-pos01 0.37 0.17 0.18 2100 shock2-pos02 0.25 0.18 0.19 2200 shock2-pos03 0.15 0.19 0.21 2400 shock2-pos04 0.085 0.24 0.26 3000

features using a ﬁve-parameter Gaussian,

2 λ − C1 flux = C0 exp(−0.5z ) + C3λ + C4, z = (4.3) C2 This profile proved insufficient as there was significant broadening in the wings of the

[O IV] feature and some of the lines appeared distinctly non-Gaussian. Instead, we used this fit to subtract a linear baseline from the data and integrated the intensity over the wavelength range of the line profile. We multiplied this value by 1.17441 , which for a πC0 purely Gaussian function would yield the full-width at half-maximum. In our data set, this value is comparable to the line width although less biased by non-Gaussian line profiles.

The values of C0, C2, the calculated line width, and the reduced chi-square of the Gaussian ﬁt are provided in Table 2 for each pointing in our high-resolution data where we produced a convolved low-resolution ﬁt. We infer from this analysis that the average line width in the region of our high-resolution spectroscopy is on the order of 2400 km s−1.

From the low-resolution [O IV] spectra of Cas A, we constructed a data cube of the remnant using IRS CUBISM (Smith et al., 2007a). The line width of [O IV] around the remnant can then be computed from the second moment of the spectrum at each position.

In many cases, the signal to noise of the [O IV] line was too low to compute a reliable

line width, so we masked every position in our data cube whose integrated [O IV] intensity 90

Figure 4.11: Low-resolution [O IV] (25.9 µm) (Smith et al., 2009) line widths given in 103 km s−1 with VLA 20 cm contours overlaid. The image has been masked for regions where the [O IV] integrated intensity was at least 3−σ the noise level, or about 10% of the −6 −2 −1 maximum value of 8 × 10 W m sr . Here we show regions with the largest [O IV] line widths, exceeding values inferred from our convolved high-resolution spectra (∼ 2400 km s−1). The black circle represents the approximate location of the reverse shock at 9500 from the expansion center (Thorstensen et al., 2001; DeLaney et al., 2010) while the green circle represents the forward shock. Note that in some places, the broad [O IV] emission extends well beyond the reverse shock, thus it is unlikely to be associated with the Bright Ring. 91 was less than the 3 − σ the level of the noise, or less than 10% of the maximum value of −6 −2 −1 3 −1 8 × 10 W m sr . In Figure 4.11, we show the [O IV] line width (in 10 km s ) in regions where it exceeds our average line width of 2400 km s−1.

There is some correlation between positions of broadening in [O IV](Figure 4.11) and CO broadening (Figure 4.4) as well as regions of bright radio emission. Highly broadened

[O IV] emission to the north, west, and southeast of Cas A support the hypothesis that ejecta in these regions are interacting with interstellar gas. The huge velocity spread of the

[O IV] line presumably arises at the outward edge of the shock where gas is being stirred strongly by the interaction, whereas the CO lines must correspond to greater depths in the molecular cloud where the effects of the shock are just being felt. In some regions, broad

[O IV] line proﬁles correspond to the location of the Bright Ring, where a large velocity spread can be attributed to the large quantity of ejecta brightened by the reverse shock. We can distinguish these two regions by their position relative to the positions of the reverse shock (Figure 4.11, black circle) and forward shock (green circle). Some of the broadened ejecta to the north, west, and southeast of Cas A correspond to positions well beyond the reverse shock, and in some cases, at projected radii consistent with the forward shock. We infer that the broad [O IV] line ejecta at these positions are associated with turbulence due to interaction with surrounding gas.

4.5 Discussion

Our study indicates that any signs of interaction between the shock from Cas A and the surrounding ISM are subtle. Broadened CO lines extending beyond the shock front offer the most direct evidence in support of an interaction between high-velocity ejecta (i.e. ejecta with velocities exceeding the spherical expansion velocity of Cas A) and molecular gas. We discuss this behavior ﬁrst, in Section 4.5.1. In Section 4.5.2 we discuss the mid- infrared lines and the evidence for turbulence they provide. We describe aspects of the SNR/ISM interaction in other supernovae and compare the physics of these interactions to Cas A in Section 4.5.3. 92

4.5.1 Possible Mechanisms for CO Line Broadening in Regions Beyond the Cas A Shock Front

Galactic molecular clouds typically have linewidths for 12CO of 1 − 3 km s−1 (Heyer and Brunt, 2004). We have observed linewidths as large as 8−10 km s−1 in regions that extend 1 − 20 past the southwest shock front of Cas A. We infer that, given the small projected distance between the gas and Cas A and LSR velocities of features in these clouds similar to those observed in previous studies of gas around Cas A (Liszt and Lucas, 1999), there must be some mechanism associated with the SNR other than direct interaction with the shock front that accounts for these linewidths. In the following discussion, we consider two possible causes of this behavior. Two remaining possibilities that we consider less likely are discussed in Appendix B. We conclude that the most likely cause of line broadening in the observed molecular clouds is interaction between the clouds and knots of material ejected at higher velocities than that of the shock front.

4.5.2 Broadening by Stellar Winds

It has been proposed that the Cas A progenitor was a star on the order of 15 − 40 M (Young et al., 2006). For stars in this range, episodes of signiﬁcant stellar winds can be described by the prescriptions of OB mass loss, red giant/supergiant mass loss, and

Wolf-Rayet (WR) mass loss leading to a total mass loss of 5 − 10 M . We consider a molecular cloud that intercepts ∼ 10% of the momentum from the wind (i.e. for wind emitted isotropically, a cloud with radius of 1.9 pc at 3 pc distance from the supernova progenitor would intercept approximately 10% of the wind). We then estimate that this wind would need to emerge with a velocity in excess of 2000 km s−1 to provide the momenta we calculated in each of our molecular clouds. This wind velocity is at the upper limit for the terminal velocity of winds emitted via OB mass loss (Kudritzki and Puls, 2000) and WR mass loss (Crowther, 2007), let alone red giant/supergiant mass loss which involves signiﬁcantly slower terminal velocities. Only in the most extreme mass loss scenario does the assumed velocity outﬂow match theoretical models. For example, Young et al.(2006) predict that for a 40 M progenitor, 93 mass loss in the form of Luminous Blue Variable (LBV) eruptions and subsequent WR mass loss would leave a pre-supernova collapse mass of 7.8 M . Ejecta mass on the −1 order of 30 M combined with a wind velocity on the order of 2000 km s for LBV winds approximately matches the momentum in turbulent molecular line widths observed in our molecular clouds, however the extreme mass-loss in this progenitor might uncover its C/O core before collapse. Observation of nitrogen-rich knots in the ejecta of Cas A preclude such a massive progenitor (Fesen, 2001). For example, recent models obtained from diffuse, thermal X-ray emission suggest that Cas A mass loss occurred via a slow RSG wind (∼ 15 km s−1) with only a moderate total mass stripped off the progenitor star

(∼ 11 M ) and even less in the wind itself (∼ 6 M )(Lee et al., 2013). Furthermore, while there is evidence that the Cas A progenitor exploded into a circumstellar bubble created by stellar winds (Reed et al., 1995; Borkowski et al., 1996), models developed from observations of the interaction between the remnant and this swept up material do not provide enough velocity in the wind (Perez-Rend´ on´ et al., 2009) to account for the CO line widths observed. We infer that while there is evidence for a circumstellar wind being swept up by Cas A, the wind could not account for the anomalous line proﬁles observed in molecular clouds toward the SNR.

4.5.3 Broadening by Interaction with Outﬂows of Fast-Moving Knots

There is evidence in optical and X-ray observations of Cas A that some fast-moving knots are moving along bipolar outflows toward the southwest and northeast of the SNR (Fesen et al., 2001, 2006; DeLaney et al., 2010). In addition, there are “ejecta pistons” (DeLaney et al., 2010) of fast-moving knots that extend in other directions besides along the outflows. These ejecta are characterized especially by iron, silicon, and oxygen emission at velocities on the order of 4000−10000 km s−1, appreciably in excess of the Cas A forward shock velocity. The bright optical and X-ray emission extends in many places 1 − 20 in projected radius beyond the forward shock. Along the observed outflow, points of bright emission form a cone with a projected opening angle calculated at 27◦ (DeLaney et al., 2010). Our observations are consistent with shock-broadened molecular lines caused by the 94

outflowing fast-moving knots. These knots have been observed at radii consistent with the projected radius of the observed molecular clouds. The energy in the fast-moving knots could be sufficient to create the broad lines. Less than one percent of the 2×1051 ergs in the Cas A supernova event (Young et al., 2006) would be sufficient to account for the ∼ 1049 4 −1 ergs (i.e. a 10 M cloud with 10 km s linewidths) of turbulent energy in the molecular cloud. For example, DeLaney et al.(2010) calculate that the fast-moving outflow carries 3 × 1050 ergs in kinetic energy, roughly 3% of which needs to be intercepted by one of the observed clouds to account for the turbulent energy we infer. Assuming the outflow is projected along a cone with opening angle 27◦ out to a distance of 8 pc (i.e. 2 pc beyond the shock front of Cas A), a cloud with a radius of 1.2 pc would intercept roughly 9% of the incident ejecta.

4.5.4 Turbulence in Mid-Infrared Line Emission

We have found that emission lines observed toward the Cassiopeia A shock front exhibit multiple, distinct Doppler components. Multiple optical and infrared studies of Cas A have revealed knots of ejecta at velocities ranging from −2300 to 3200 km s−1 (our Fig- ure 4.7; Lawrence et al., 1995; Reed et al., 1995; DeLaney et al., 2010). These knots reveal a complex, three-dimensional picture of the remnant, although one in which most of the expanding ejecta exhibit roughly spherical geometry with broad line regions located toward the optical shell. This characteristic of Cas A spectra supports the interpretation of extreme turbulence in bright, high-velocity [O IV] emission along the northern shock front. Here we will explore the possible mechanisms that may account for the conversion of the kinetic energy into turbulence. The physical scenario represented by turbulent line proﬁles in the forward shock of a SNR is at odds with the canonical model of a freely-expanding remnant. Any initial turbulence in the ejecta provided by the supernova event will rapidly dissipate unless some additional interaction can provide energy to the ejecta in order to excite turbulent motion on smaller scales. The most plausible candidates for this interaction are MHD driven turbulence due to the magnetic ﬁeld inherent in the shock itself or a direct interaction between the forward shock and dense gas in the interstellar medium. 95

To evaluate the first possibility, the magnetic field in the Cas A shock front can be indi- rectly probed through X-ray synchrotron emission driven by first-order Fermi acceleration (Vink and Laming, 2003b). Turbulence generated by, for example, the interaction between diffusive shock driven cosmic rays and the supernova ejecta may account for the observed magnetic field ‘frozen-in’ to plasma in the ejecta (Volk and McKenzie, 1981; Bell and Lucek, 2001; Bell, 2004). In this way, magnetic fields provide a direct probe of the energy necessary to drive turbulence in the shock front. The observed magnetic field toward the Cas A shock front has a strength measured between 0.08 mG (Vink and Laming, 2003b) and 0.5 mG (Saha et al., 2014). We approximate the average turbulent line velocity as,

s B2 vturb = (4.4) µmHneff If we take 0.5 mG as a rough upper limit for the magnetic field strength at the shock −3 front and assume a gas density on the order of neff = 30 cm (Saha et al., 2014), we can calculate an energy density in the magnetic field and assume this value is comparable to the energy density in turbulence. The calculated turbulent velocities are no more than 200 km s−1 which is significantly slower than the observed turbulence. We infer that any turbulence related to interactions with the ‘frozen-in’ magnetic field in the ejecta plasma is insufficient to explain the observed high-velocity features. The second mechanism that may generate turbulence in supernova ejecta is interaction with large scale density fluctuations in the medium surrounding Cas A. As the Cas A shock front freely expands outwards, interaction between the ejecta and interstellar material will provide turbulent energy such that high-velocity Doppler components and large line widths persist in the fine structure lines toward the shock front. A direct interaction between the SNR and surrounding gas explains both the sharp cut off in the spectral lines and the lack of change in the relative strength of the high-velocity features at different positions. Young SNRs exhibit a contact discontinuity in their radial density profile between ejecta and circumstellar or interstellar gas behind which turbulence associated with Rayleigh-Taylor instabilities generate magnetic fields (Gull, 1975). The dominant mechanism for radio emission in Cas A is believed to be synchrotron radiation 96 generated by an interaction between relativistic electrons in the ejecta and the magnetic field behind this contact discontinuity (Rosenberg, 1970; Keohane et al., 1998). Thus, a sharp contact discontinuity between Cas A and the interstellar gas would exhibit the cut off in line strength we observe over a relatively small (∼ 1000) transverse distance. Furthermore, the consistently high turbulent velocities over this distance indicate that the interaction between Cas A and the surrounding gas has stirred all the observed layers of ejecta recently such that the power in turbulent motion in the ejecta has not had time to dissipate from this position in the shock front. We can point to other signatures as an indication of interaction between ejecta and swept up interstellar or circumstellar material. The bright radio emission behind the northern, western, and southern shock front of Cas A suggests that either stronger magnetic fields or more efficient particle acceleration in this region must be invoked to explain the strong synchrotron component (DeLaney et al., 2002). We infer that an interaction between the ejecta and swept up material both causes the turbulent line velocities and forms a shock interface where efficient Fermi acceleration occurs. This explanation would also agree with a scenario where the forward shock is beginning to form an interface with a molecular cloud located toward this region of the SNR. Other regions where there exist both broadened molecular lines and bright radio emission could be sites of a contact discontinuity between the Cas A forward shock and molecular gas in the interstellar medium. There does not appear to be significant CO line broadening toward the cloud observed north of Cas A, despite the broad range of velocities in the IR fine-structure lines. It appears that the shock has not reached the bulk of the mass of the molecular cloud, represented by the CO. To understand the situation, we estimate the relative masses involved in the two types of emission. Hines et al.(2004) found that the CO-emitting molecular cloud is about 650 M . In the appendix, we show that the column density of ejecta along the northern shock front is N = 1.1×1020 cm−2. Assuming a distance of 3.4 kpc, we find that approximately 0.05 M of material within the IRS slit would be sufficient to produce the infrared lines at position S2, P1. That is, allowing for ten IRS slit areas of emission (about 10 by 100 arcsec), some three orders of magnitude less material is required for the infrared lines as is needed to produce the CO emission. Therefore, the observations are consistent 97 either with the infrared lines arising through interaction with a much smaller body of gas isolated from the main CO-detected cloud, or with a situation where the shock has just reached the cloud and has only penetrated slightly into it. The latter possibility would be consistent with Figure 4.1 and Figure 4.4, which show the enhanced radio synchrotron emission indicative of the interaction region only lying along one edge of the cloud.

4.5.5 Comparison to Other Studies of SNR/MC Interactions

Previous studies of SNR/MC interactions focus on several signatures of shocked molecular gas. The presence of the broadened molecular lines, the OH 1720 MHz maser, and near-IR molecular hydrogen lines are common tracers that distinguish regions of shocked gas around an expanding SNR. We describe the well-observed SNRs IC443, W44, and W28 and the tracers used to identify interaction with molecular clouds in their vicinity for comparison with Cas A. While these objects are at approximately the same distance as Cas A they are consid- erably older. IC 443 is 1.5 kpc and 4000 yr (Rosado et al., 2007; Lozinskaya, 1981; Troja et al., 2008), W44 is 3 kpc and 20,000 yr (Wolszczan et al., 1991), and W28 is 2.8 kpc and 58,000 yr (Kaspi et al., 1993). Age is an important factor in evaluating the characteristics of SNR/MC interactions for two reasons. First, Cas A is smaller than the more evolved examples of SNR/MC interactions. The entire SNR measures approximately 50 in diameter or about 5 pc at 3.4 kpc distance. Shocks caused by the blast wave from Cas A are less likely to have propagated through the molecular cloud, so any observed variations in the turbulent linewidths in these clouds are more likely to be confined to regions close to the Cas A shock front. Second, Cas A is still expanding at approximately the free expansion velocity (3000−4000 km s−1). For hydrodynamic models with shock speeds > 25 km s−1, molecular clouds exhibit significant collisional dissociation of H2 molecules (Kwan et al., 1977). Therefore, any interaction with molecular gas is likely to dissociate gas at the point of interaction between the two media. As the shock wave from an interaction propagates into a molecular cloud, emission lines from any surviving molecules will be broadened relative to the un-shocked material. This phenomenon is observed for a number of other “broad molecular line” sources, including 3C391 (Reach and Rho, 1999) and HB21 (Koo 98 et al., 2001). The SNR/MC interaction for IC 443 is thoroughly studied and is useful to compare with the younger interaction near Cas A. The interaction for IC 443 was first identified by Cornett et al.(1977) and has been confirmed by detection of molecular species toward the associated molecular clouds (White et al., 1987; Xu et al., 2011), shocked OH, CO, and H2 emission (Denoyer, 1979b,a; Burton et al., 1988), and shock excitation of the 1720 MHz OH maser (Claussen et al., 1997; Hewitt et al., 2006). The physical conditions in this SNR appear to be well-suited to detect tracers of SNR/MC interactions. The interacting molecular cloud is in front of the remnant along the line of sight and possesses more mass 4 (> 10 M ) than the clouds detected near Cas A. The cloud exhibits linewidths of > 30 km s−1, up to 90 km s−1 in regions with the most turbulent molecular gas (White et al., 1987). This result has been confirmed in several molecules, probing molecular gas of densities from 105 cm−3 to 3 × 106 cm−3 (van Dishoeck et al., 1993). The same result has been observed in W44 and W28, where CS, CO, and HCO+ line widths are 20 − 30 km s−1 (Reach et al., 2005). The resulting turbulence driven fragmentation may suppress star formation in regions of giant molecular clouds where SNR/MC interactions occur; smaller clouds have also been identified around IC 443, down to mass scales of 0.1 − 0.3 M . Shock excitation in the SNR/MC interactions located toward IC 443 cause this remnant to be one of the most luminous molecular hydrogen sources in the galaxy (Burton et al., 1988). Since it is difficult to excite bright emission in this species without dissociating the molecule, near-IR molecular hydrogen emission is predominantly useful in more evolved SNR as a tracer for SNR/MC interactions. In IC 443, the correlation between line emission + from H2 v = 1 − 0 S(1) and the more easily detected CO, HCO , and HCN emission at longer wavelengths provides a basis for mapping the region where the SNR is interacting with dense molecular gas. The OH 1720 MHz maser is one of the most commonly used tools to identify SNR/MC interactions. Spectral line emission from this phenomenon is suited to more evolved SNRs where C-type (nondissociative) shock interfaces are observed between the supernova and a molecular cloud. Masers around W28 and W44 are observed to have sizes of 150 - 1000 AU (Hoffman et al., 2005) with cores 60 AU in size. In these regions where the number 99

n ∼ 104 −3 density of molecular hydrogen is on the order of H2 cm , the shock velocity cannot exceed 10 km s−1. The physical conditions that give rise to the OH 1720 MHz maser are so specialized that they occur only in a small fraction of the remnant volume and these conditions may not appear at all in small and young remnants. Cas A is a very young SNR (∼ 350 yr) with a high shock velocity (3000 − 4000 km s−1). Any interaction between the SNR and nearby molecular clouds will be at an early stage of development considering that the diameter of Cas A is comparable to the 4 size of a typical 10 M molecular cloud. Our observations do show subtle indications of an interaction between the SNR and molecular clouds. These indicators include the extended region of broadened CO lines, and the multiple high-velocity components of the mid-infrared ﬁne structure lines around the rim of the SNR. However, the lack of more dramatic tracers, such as even broader CO lines, bright emission from molecular hydrogen, and OH maser emission can all plausibly be explained by the youth and small size of the Cas A remnant or, alternatively, a small amount of mass in the fastest ejecta knots.

4.6 Conclusion

From our OTF maps in 12CO and 13CO J = 2 − 1 of the region around Cas A, and high- resolution spectroscopy from 10−40 µm across two tracks in declination across the northern shock front of Cas A, supplemented by a VLA 20 cm map of Cas A (Anderson et al.,

1991) and low-resolution spectroscopy of [O IV] (25.94 µm) (Smith et al., 2009), we ﬁnd the following:

1. Bright radio continuum emission along the northern edge of Cas A toward a previously identiﬁed molecular cloud suggests acceleration of particles at the shock front/cloud interface.

2. There is broadening of 12CO J= 2 − 1 in CO spectral features across the western and southern shock fronts of Cas A. The presence of broadened molecular gas that extends well beyond the boundaries of the SNR in these directions implies an interaction between Cas A and the molecular gas. We infer that a shock interaction with 100

ejecta in the fast-moving knots, which emerge at velocities substantially greater than that of the shock front, has caused the observed CO line broadening.

3. There are a number of mid-infrared fine-structure lines across the northern shock front of Cas A. Some of these lines exhibit significant high-velocity splitting that is not accounted for by typical supernova outflows. These lines suggest an interaction with interstellar material, with the resulting turbulence yielding a distribution of large radial velocities.

4. Broadening in the low-resolution spectra of [O IV] conﬁrms that high velocities occur along the northern shock front of Cas A as well as along the western and southern shock fronts in the same regions where we observe CO broadening.

There is subtle evidence in our observations for an interaction between Cas A and the interstellar medium along the northern shock front as well as toward molecular clouds located to the west, south, and southeast. Our interpretation is consistent with previous submillimeter and mid-infrared studies of Cas A as well as optical and X-ray observations of fast-moving knots toward the remnant. 101

CHAPTER 5

A Systematic Survey for Broadened CO Lines Toward Galactic Supernova Remnants

We present molecular spectroscopy toward 50 Galactic supernova remnants (SNRs) taken at millimeter wavelengths in 12CO and 13CO J = 2 − 1 with the Heinrich Hertz Submil- limeter Telescope as part of a systematic survey for broad molecular line (BML) regions indicative of interactions with molecular clouds (MCs). These observations reveal BML regions toward nineteen SNRs, including nine newly identified BML regions associated with SNRs (G08.3−0.0, G09.9−0.8, G11.2−0.3, G12.2+0.3, G18.6−0.2, G23.6+0.3, 4C−04.71, G29.6+0.1, G32.4+0.1). The remaining ten SNRs with BML regions confirm previous evidence for MC interaction in most cases (G16.7+0.1, Kes 75, 3C 391, Kes 79, 3C 396, 3C 397, W49B, Cas A, IC 443), although we confirm that the BML region toward

HB 3 is associated with the W3(OH) H II region, not the SNR. Based on the systemic velocity of each MC, molecular line diagnostics, and cloud morphology, we test whether these detections represent SNR-MC interactions. One of the targets (G54.1+0.3) had previous indications of a BML region, but we did not detect broadened emission toward it. Although broadened 12CO J = 2 − 1 line emission should be detectable toward virtually all SNR-MC interactions we find relatively few examples; therefore, the number of interactions is low. This result favors mechanisms other than SN feedback as the basic trigger for star formation. In addition, we find no significant association between TeV gamma- ray sources and MC interactions, contrary to predictions that SNR-MC interfaces are the primary venues for cosmic ray acceleration.

5.1 Introduction

Supernovae (SNe) inject kinetic energy, momentum, and enriched ejecta into the interstellar medium (ISM) and are thought to be one of the main sources of interstellar turbulence and galactic outflows. An important factor in understanding these processes is the evolu- 102 tion of supernova remnants (SNRs) as their ejecta encounter various phases of the ISM. For example, a large fraction of core-collapse SNe from high-mass stars may explode promptly and interact with the molecular cloud (MC) from which they formed. This hypothesis suggests that the direct interaction between supernova (SN) ejecta and MCs, which we refer to as SNR-MC interactions, may occur at a rate significantly higher than expected from the filling fraction of molecular gas in the ISM (Elmegreen and Lada, 1977). While it has been estimated that as many as half of all SNRs have ejecta in physical contact with MCs (Reynoso and Mangum, 2001), this estimate is highly uncertain and relies on a number of assumptions about massive stars, SNe, and the relation between SNe and MCs. Detection of SNR-MC interactions has traditionally relied on a number of methods for observing shocked molecular gas directly. Strong, narrow emission from the OH 2 1720 MHz maser ( Π3/2,J = 3/2,F = 2 → 1) is often detected in the vicinity of SNRs, especially those with other signatures of SNR-MC interactions (Goss and Robinson, 1968; Haynes and Caswell, 1977; Denoyer, 1979b). More recent surveys have established that targeted searches for OH maser emission may be a reliable method of searching for SNR- MC interactions (e.g., Frail et al., 1996; Green et al., 1997). To date, however, OH masers have only been detected toward ∼ 10% of the 294 known SNRs in our galaxy (Green, 2014), although a lack of detailed spectroscopic information and localization makes it impossible to associate all of these detections with SNRs. This low incidence relative to the expected rate of SNR-MC interactions may be due to the unique physical conditions required for maser emission. OH 1720 MHz masers are collisionally excited behind slow (< 45 km s−1) C-type shocks in regions with moderate temperatures and densities T ∼ 50 − 125 n ∼ 105 −3 ( K, H2 cm )(Lockett et al., 1999). Studies of the warm, diffuse

H2 regions around interacting SNRs indicate that physical conditions are often outside this range (Hewitt et al., 2009). Regions with higher shock velocities could also depopulate the 2 hydroxyl Π3/2 state or dissociate the molecule entirely. Convincing evidence of SNR-MC interactions must come from other sources in these cases. Internal ﬂows within SNRs are expected to become highly turbulent where the SNR shock front interacts with the dense ISM (Jun and Jones, 1999). Molecular lines toward SNRs, such as the rotational transitions of 12CO, should be velocity-broadened as a result 103

of increased turbulence from a SNR-MC interaction. Shocked MCs are also known to exhibit asymmetric line profiles where these interactions occur (Reach and Rho, 1999; Koo et al., 2001; Jiang et al., 2010). As opposed to the presence of OH 1720 MHz maser emission, the detection of velocity-broadened, broad molecular line (BML) 12CO relies on more detailed spectroscopic information, especially in faint line “wings” where turbulent broadening is usually most apparent. At the same time, this type of observation offers insight into SNR evolution and the physical properties of both SN ejecta and MCs. Even if the studies of OH maser emission, BML emission, and other shock indicators are all taken together, only about 60 Galactic SNR-MC interactions have been identified. However, to date there has been no large-scale systematic survey for shock-broadened molecular lines. Instead, the interactions have not been identified in a homogeneous manner, but are generally a result of investigations of individual SNRs, notably IC 443, W28, W44, CTB 109, Kes 69, Kes 79, W51C, and Cassiopeia A (Cas A) (Wootten, 1977; De- noyer, 1979a; Tatematsu et al., 1990; Green and Dewdney, 1992; Koo and Moon, 1997; Arikawa et al., 1999; Reach and Rho, 1999; Zhou et al., 2009; Kilpatrick et al., 2014). These sources span a wide range of ages and morphologies, from young, shell-like SNRs such as Cas A, through middle-aged and older (> 104 yr) remnants, and including mixed- morphology remnants and plerions. A systematic and unbiased survey of broadened CO emission toward Galactic SNRs may yield a larger incidence of such interactions. The work described in this paper uses CO emission for the first large-scale, systematic search for BML regions toward SNRs. To interpret these observations, we developed an algorithm to identify regions of broadened 12CO J = 2 − 1 emission where they exist toward each SNR-MC system. We justify these identifications using the approximate systemic velocity at the distance inferred to each SNR, an estimate of the “broadened” CO emission, and the radio continuum morphology of the remnant itself. In this way, we have identified BML regions toward nineteen SNRs in our sample, including nine newly identified candidate SNR-MC interactions.

5.2 Source Selection and Distances 104

Table 5.1: Supernova Remnants in Survey

Cat. No. Name α (J2000) δ (J2000) Region RMS Distance1 BML? Type Pulsar γ-ray (σ) Ref. (h m s) (d m) (K) (kpc)

G4.5+6.8 Kepler 17 30 42 −21 29 100 × 100 0.17 4.0 S H(0.68) G08.3−0.0 18 04 34 −21 49 100 × 100 0.08 16.3 D S H(14.6) G09.9−0.8 18 10 41 −20 43 130 × 100 0.11 6.4 D S H(-0.1) G10.5−0.0 18 09 08 −19 47 100 × 100 0.15 14.4 S H(4.6) G11.2−0.3 18 11 27 −19 25 100 × 100 0.09 4.4 D C Y H(5.5) 1 G11.8−0.2 18 12 25 −18 44 100 × 100 0.09 19.4 S H(1.7) G12.0−0.1 18 12 11 −18 37 70 × 70 0.12 10.9 S H(0.7) G12.2+0.3 18 11 17 −18 10 100 × 100 0.09 15.6 D S H(2.4) G12.5+0.2 18 12 14 −17 55 100 × 100 0.12 15 C H(4.7) G12.8−0.0 18 13 37 −17 49 100 × 100 0.11 4.8 C Y H(24.3) 2 G13.5+0.2 18 14 14 −17 12 100 × 100 0.12 13.3 S H(3.3) G14.3+0.1 18 15 58 −16 27 100 × 100 0.11 18.7 S H(1.8) G15.9+0.2 18 18 52 −15 02 80 × 70 0.11 8.5 S H(1.4) G16.7+0.1 18 20 56 −14 20 100 × 100 0.15 10 DP C H(4.7) 3,4,5 G17.0−0.0 18 21 57 −14 08 100 × 100 0.18 18.1 S H(6.6) G18.1−0.1 18 24 34 −13 11 80 × 80 0.13 5.6 S H(6.9) G18.6−0.2 18 25 55 −12 50 60 × 60 0.13 13.2 D S H(2.8) G21.5−0.1 18 30 50 −10 09 100 × 100 0.09 18.9 S H(1.2) G21.5−0.9 18 33 33 −10 35 100 × 100 0.07 4.8 C Y H(6.4) 7 G23.6+0.3 18 33 03 −08 13 80 × 80 0.12 6.9 D ? H(1.5) G27.4+0.0 4C−04.71 18 41 19 −04 56 100 × 100 0.10 8.5 D S Y H(6.0) 3,8 G29.6+0.1 18 44 52 −02 57 80 × 80 0.16 10 D S Y H(4.5) 9 G29.7−0.3 Kes 75 18 46 25 −02 59 100 × 100 0.13 6.0 DP S Y H(10.1) 3,10,11 G30.7−2.0 18 54 25 −02 54 170 × 160 0.14 4.7 ? H(0.9) G31.5−0.6 18 51 10 −01 31 190 × 180 0.11 6.3 S H(1.9) G31.9+0.0 3C 391 18 49 25 −00 55 100 × 100 0.10 7.2 DP S H(0.0) 5,12,13,14 G32.4+0.1 18 50 05 −00 25 100 × 90 0.10 17 D S H(2.2) G33.2−0.6 18 53 50 −00 02 190 × 180 0.11 5.7 S H(2.5) G33.6+0.1 Kes 79 18 52 48 +00 41 100 × 100 0.13 7.1 DP S Y H(2.7) 3,15,16 G36.6+2.6 18 48 49 +04 26 190 × 130 0.09 8.7 S H(-1.4) G39.2−0.3 3C 396 19 04 08 +05 28 100 × 100 0.11 6.2 DP C Y? H(0.5) 3,17,18,19 G41.1−0.3 3C 397 19 07 34 +07 08 100 × 100 0.10 10.3 DP S H(3.0) 3,20 G43.3−0.2 W49B 19 11 08 +09 06 100 × 100 0.09 2.5 DP S Y? H(1.0) 3,21,22 G54.1+0.3 19 30 31 +18 52 100 × 100 0.09 6.2 P C Y H(3.8) 3,17,23,24 G57.2+0.8 4C-21.53 19 34 59 +21 57 130 × 100 0.19 8.2 S H(0.8) G59.8+1.2 19 38 55 +24 19 240 × 100 0.11 7.3 F H(1.6) G63.7+1.1 19 47 52 +27 45 100 × 100 0.12 3.8 F G69.0+2.7 CTB 80 19 53 20 +32 55 100 × 100 0.07 1.5 F Y M(0.6) 25 G69.7+1.0 20 02 40 +32 43 160 × 190 0.14 7.1 S G74.9+1.2 CTB 87 20 16 02 +37 12 80 × 60 0.08 6.1 F Y? V(6.2) 17,21 105

G76.9+1.0 20 22 20 +38 43 100 × 100 0.10 10 C Y 17 G83.0−0.3 20 46 55 +42 52 100 × 100 0.15 11.9 S G84.2−0.8 20 53 20 +43 27 220 × 160 0.11 6.0 S G111.7−2.1 Cas A 23 23 26 +58 48 100 × 100 0.05 3.4 DP S Y M(5.2) 26,27,28,29 G116.9+0.2 CTB 1 23 59 10 +62 26 see Section 5.3 0.09 3.1 S Y 30 G120.1+1.4 Tycho 00 25 18 +64 09 100 × 100 0.10 2.5 S V(5.8) G130.7+3.1 3C 58 02 05 41 +64 49 100 × 100 0.08 2.0 F Y M(5.7) 31 G132.7+1.3 HB 3 02 17 40 +62 45 see Section 5.3 0.11 2.0 DP S Y 32,33 G184.6−5.8 Crab 05 34 31 +22 01 100 × 100 0.10 2.0 F Y H(129) 34 G189.1+3.0 IC 443 06 17 00 +22 34 90 × 80 0.09 1.5 DP C Y? V(8.3) 35,36

The SNRs in our survey were selected from those accessible with the Hein- rich Hertz Submillimeter Telescope (SMT) on Mt. Graham, Arizona, (λ, φ) = (−109.8920, +32.7013). In general, all SNRs with declination greater than −22◦ were considered, yielding a sample between Galactic longitude +4.5 and +190 degrees. There

1Distances are those we adopt for this paper, as described in Section 5.2.1. For BML Detection, D: Detected broadened CO feature in this paper, P: Previous SNR-MC interactions as reported in Jiang et al. (2010). For Type, we report the SNR morphology from Green(2014) where S: shell, C: composite, and F: filled except for G12.0−0.1 which is described as shell-like in Kassim(1992), G59.8 +1.2 which is described as a plerion/filled shell SNR in Sun et al.(2011), and CTB 80 which is described as an extended plerion/filled shell SNR in Angerhofer et al.(1981). SNRs whose types have not yet been identified are indicated with a ?. For pulsar, Y: SNR has a confirmed (i.e., timed) pulsar, Y?: SNR has a coincident, compact X-ray or radio source. We include TeV gamma-ray detections by experiment and significance of detection in parentheses for HESS (H; Aharonian et al., 2006a, 2008a; Bochow, 2011), MAGIC (M; Albert et al., 2007a; Aleksic´ et al., 2014), and VERITAS (V; Acciari et al., 2009, 2011; Aliu, 2011). For references related to detections of pulsars/compact objects as well as previous evidence for SNR-MC interactions: (1) (Torii et al., 1997); (2) (Gotthelf and Halpern, 2009); (3) (Green et al., 1997); (4) (Reynoso and Mangum, 2000); (5) (Hewitt et al., 2008); (7) (Gupta et al., 2005); (8) (Vasisht and Gotthelf, 1997); (9) (Vasisht et al., 2000); (10) (Gotthelf et al., 2000); (11) (Su et al., 2009); (12) (Reynolds and Moffett, 1993); (13) (Frail et al., 1996); (14) (Reach and Rho, 1999); (15) (Green and Dewdney, 1992); (16) (Gotthelf et al., 2005); (17) (Biggs and Lyne, 1996); (18) (Hewitt et al., 2009); (19) (Su et al., 2011); (20) (Jiang et al., 2010); (21) (Gorham et al., 1996); (22) (Keohane et al., 2007); (23) (Camilo et al., 2002); (24) (Lee et al., 2012); (25) (Kulkarni et al., 1988); (26) (Tananbaum, 1999); (27) (Liszt and Lucas, 1999); (28) (Hines et al., 2004); (29) (Kilpatrick et al., 2014); (30) (Hailey and Craig, 1995); (31) (Murray et al., 2002); (32) (Lorimer et al., 1998); (33) (Routledge et al., 1991); (34) (Lovelace et al., 1968); (35) (Olbert et al., 2001); (36) (Cornett et al., 1977) 106 are currently 160 known SNRs in this volume (Green, 2014). From this larger sample, we selected SNRs with an apparent size less than 200 in diameter along their largest axis. The targets we observed are listed in Table 5.1. Some targets listed in Table 5.1 were selected for comparison to our larger sample. These SNRs are well-studied and large (diameter > 200) and in these cases we mapped a smaller subregion of the entire remnant. These targets include CTB 80, CTB 1, HB 3, and IC 443. IC 443 is a well-known SNR-MC system, while HB 3 is suspected to have SNR-MC interactions (Routledge et al., 1991).

5.2.1 Distances

Reliable distances are critical when characterizing SNRs. In order to draw an association between broad-line MCs and the SNRs themselves, a variety of distance indicators exist such as proper motion estimates, radial velocity, H I absorption, and OB associations, each with varying degrees of uncertainty. Where the properties of SNRs and their environments differ, distance indicators can yield estimates with signiﬁcant disagreement. For example, the distance to G4.5+6.8 (Kepler/SN 1604) has historically been controversial, with estimates ranging from around 4 − 6 kpc (Leibowitz and Danziger, 1983; White and Long, 1983; Reynoso and Goss, 1999; Vink, 2008b; Chiotellis et al., 2012) to 10 − 12 kpc (Milne, 1970; Ilovaisky and Lequeux, 1972; Borkowski et al., 1992), although proper motion estimates now favor the closer distance. This lack of agreement between various methods underscores the systematic uncertainties present in all distance indicators, from discrepant optical absorption (Danziger and Goss, 1980) to disagreement in the maximum optical brightnesses of historical SNe (van den Bergh et al., 1973). For these reasons, we carefully chose distance indicators from among the literature, favoring those determined from proper motion estimates, kinematic estimates from H I absorption and CO associations, and line-of-sight X-ray absorption. The values we chose are given in Table 5.1 and described in Appendix C. However, for nineteen of the remnants in our survey, we were unable to ﬁnd reliable distance indicators in the literature. For these objects, we resort to the surface brightness-size or Σ − D relation (Shklovskii, 1960; Clark and Caswell, 1976). This empirical relationship describes SNR evolution with 107

time, whereby the radio surface brightness (Σ) decreases monotonically while the physical diameter (D) increases monotonically. Thus, a power law relationship can be ﬁt such that

Σ = ADβ (5.1)

A and β are calibrated using SNRs whose distances are known from other methods. Mea- suring the radio surface brightness and angular size of the SNR then yields an angular diameter distance. Pavlovic et al.(2014) fit this relationship and derive distances to 225 Galactic SNRs. The authors find the best-fit slope to be β ≈ −5.2. This method involves significant systematic uncertainties and should be used with caution. Distance measurements for some of the SNRs in our sample are taken from this study where we adopt the authors’ most probable distance (i.e., the orthogonally fit value from Table 4 in Pavlovic et al.(2014)). These include sixteen out of the 50 SNRs we observed. We provide our own estimates for the remaining three SNRs, which we describe in Appendix C.

5.3 Observations

We mapped the vicinity of 50 SNRs (Table 5.1, Figure 5.1) with the 10m SMT using the ALMA-type sideband separating 1.3 mm receiver between 2013 April 21 and 2015 May 8. Apart from the size and average RMS noise of each map as given in Table 5.1, the observations and reduction for each remnant were identical to those described in Kilpatrick et al.(2014). As discussed in Section 5.2, there were several targets for which we did not observe the entire remnant. Rather, we targeted either a region we knew to have broadened lines or an X-ray (i.e., 0.5 − 10 keV) bright region toward any known molecular emission. These objects were used either for comparison to other targets (HB 3, IC 443) or as a targeted search for new BML regions toward large remnants (CTB 80, CTB 1). In particular, we targeted two regions of X-ray bright emission from CTB 1, centered at 23h59m05s, + 62◦3000000 with dimensions (α×δ) 100 ×300 and at 0h03m05s, +62◦4500000 with dimensions 100 × 200. Similarly, for HB 3 we targeted two regions centered at 2h16m15s, + 63◦4000000 with dimensions 100 ×300 and at 2h19m15s, +63◦5000000 with dimensions 100 ×100. These 108

Figure 5.1: 12CO J = 2 − 1 emission toward all 50 SNRs in our sample averaged over each OTF map from −150 to +150 km s−1. The spectra are scaled such that the maximum ∗ antenna temperature in each spectrum has Ta = 1. In the upper left of each panel, we give ∗ ∗ the name of the remnant and the scale factor where Ta(K) = Ta(scaled) × scale factor. regions were centered on local X-ray enhancements as reported by Lazendic and Slane (2006). For the other large targets, we used single ﬁelds as reported in Table 5.1. We analyzed these four remnants and we indicate whether BML regions were detected in each case. However, given the fact that we did not map the entirety of each remnant and some were selected on the basis of previous evidence for SNR-MC interactions, we discount all four in our discussion of SNR-MC interactions (except in Section 5.6.5) as a possible source of bias. 109

5.4 Algorithm for Selecting BML Regions from Observations

5.4.1 Data

In our analysis, we used the 1 MHz resolution data, which has 512 channels for a velocity range of 667 km s−1 at 230 GHz. For a receiver tuned to the 12CO J = 2 − 1 and centered at a systemic velocity of 0, we can be reasonably certain that all Galactic molecular emission along the line-of-sight is observed. The expected velocity range of the CO spectrum at each distance is determined by assuming a 50% uncertainty on the distances we assume to each SNR, which agrees with typical fractional errors for distances calculated using the Σ − D relation (Pavlovic et al., 2014). Comparison with the Clemens rotation curve (Clemens, 1985) allows us to limit the velocity range over which MCs might exist that are interacting with a SNR. We extend this range by 15 km s−1 at either end of the range in order to account for systematic errors in ﬁnding BML regions associated with SNRs (see, e.g., Section 5.4.5). Figure 5.2 demonstrates this method for the known BML region toward 3C 391. The distance to this SNR is taken to be 7.2 ± 3.6 kpc, which at Galactic coordinates l = 31.9, b = 0.0 corresponds to a systemic velocity range of +59 to +102 km s−1 or +44 to +117 km s−1 with the added velocity range.

5.4.2 Summary of BML Region Algorithm

Within the range of permitted velocities for each SNR, we imposed criteria on the line width and proﬁle of observed molecular emission which we describe and justify in Sec- tion 5.4.3. In summary, these criteria are:

• We required CO emission to be detected at 3 Trms (the root mean square of emission in the line-free spectrum at each pixel as in Table 5.1) in every spectral channel we used to calculate the line proﬁle.

• We used only ﬁve spectral channels for the width determination.

• We required that the line proﬁle not include secondary maxima. 110

Figure 5.2: (Left) Integrated 12CO J = 2 − 1 emission toward 3C 391 from +102 to +108 km s−1. Orange contours are 20 cm radio continuum emission of the remnant from Reynolds and Moffett(1993). Pixels with broadened molecular emission are shown as white squares. We denote a comparison region of unbroadened, bright molecular emission with a larger red square. (Right) The average spectrum from these pixels at the velocity of the BML identiﬁcation (solid black) along with a spectrum of the unbroadened, bright comparison region averaged over the pixels denoted by the red square (dashed red). Shaded parts of the spectrum correspond to velocities ruled out by the inferred distance to the SNR. 111

• We required the FWHM of observed 12CO lines > 6 km s−1.

The purpose of these criteria was to identify regions of molecular emission that correspond with disturbed molecular gas where a SNR-MC interaction has occurred. At the same time, we attempted to minimize false positives, e.g. due to molecular emission overlapping in velocity-space, and false negatives, e.g. where interactions are subtle and exhibit small velocity-widths, the CO line is weak, or strongly asymmetric line proﬁles are present.

5.4.3 Criteria for Possible Interactions

For clouds with sizes on the order of a few parsecs, quiescent 12CO line widths are typically ∼ 2 − 3 km s−1 (Heyer and Brunt, 2004). Therefore, we adopted a threshold of a 6 km s−1 FWHM to indicate BML regions. We evaluate and discuss the validity of this threshold in Section 5.4.5. Furthermore, BML regions can have asymmetric line profiles where the molecular gas has been shocked and accelerated. In addition to large line widths (i.e., > 10 km s−1) at negative or positive velocities relative to the peak, the line profile is characterized by faint wings where turbulent acceleration of molecular gas is apparent. This line profile can be confused with multiple MCs along the same line-of-sight (e.g., toward the Galactic Center) resulting in superimposed emission in the spectrum of a small region or even a single pixel. Thus, identifying BML regions ideally relies on as small a region as possible with as few spectral channels as possible. To account for line widths, asymmetry, and molecular emission from multiple sources along a line-of-sight, we developed an algorithm for identifying BML regions in 12CO J = 2 − 1. Initially, we constructed a list of individual spectral channels by examining the spectrum at each pixel. Where the 12CO emission in individual spectral channels exceeded

10 Trms and represented a local maximum, we flagged that channel for analysis. These flagged spectral features span the entire data cube and represent the line peaks of individual emission features in our data. At 230 GHz and for a spectrum with 1 MHz resolution, the 6 km s−1 velocity-width threshold we adopted corresponds to a feature whose FWHM is 4.6 MHz or roughly five 112

spectral channels. Thus, to measure the line width of a broad feature, we required at least ﬁve channels or four channels in addition to the ﬂagged local maximum of a spectral feature. Confusion where a single spectral channel contains emission from multiple MCs can be problematic. In part, we minimized confusion by limiting the number of channels we used to measure line widths to the absolute minimum. However, this method does not preclude the scenario in which two molecular features with narrow widths (i.e., > 3 km s−1) are separated by less than their half widths in velocity-space. Therefore, in addition to the requirements described above, our algorithm stipulates that the emission on either side of the line peak must decrease monotonically in each channel and that the

emission in each channel exceeds 3 Trms. Finally, using the five spectral channels around each flagged line peak, we measured the line widths using a method that accounts for possible asymmetries in the line profile.

We took the spectrum at a particular pixel (x0, y0) and channel v0 corresponding to one of the ﬂagged local maxima in our data cube. Here, we use the antenna temperature at a ∗ ∗ particular pixel, Ta(v) = Ta(x0, y0, v). Next, we measured the velocity centroid vc of the

line peak around v0 by calculating the ﬁrst moment of the spectrum

2 P ∗ viTa(vi) v = i=−2 c 2 (5.2) P ∗ Ta(vi) i=−2 We then evaluated the second moment of the spectrum in this range,

v u 2 u P ∗ 2 u Ta(vi)(vi − vc) u i=−2 σ = u 2 (5.3) u P ∗ t Ta(vi) i=−2 This statistic is the velocity dispersion of the spectrum around the line peak. For a Gaussian line profile, the FWHM = 2.355σ. For comparison to other line indicators, we used the velocity-width statistic ∆v = 2.355σ. If this value exceeds 6 km s−1, the peak of the corresponding line profile is flagged in our data cube as a BML. For SNRs with newly 113

Table 5.2. Newly Detected BML Regions Toward SNRs

Cat. No. # pix. (total) v¯ ∆v d (d ) c vc Tab. 1 (km s−1) (km s−1) (kpc)

G08.3−0.0 7 (10) +2.6 7.3 16.4 (16.3) G09.9−0.8 8 (10) +31 7.1 4.0 (6.4) G11.2−0.3 15 (21) +32 8.1 7.2 (4.4) G12.2+0.3 2 (3) +50 6.9 11.8 (15.6) G18.6−0.2 12 (15) +42 8.5 12.7 (13.2) G23.6+0.3 7 (11) +91 6.7 5.6 (6.9) G27.4+0.0 12 (20) +100 9.0 9.0 (12.8) G29.6+0.1 5 (5) +94 10.0 9.0 (10) G32.4+0.1 3 (4) +43 7.5 11.8 (18.5)

Note. — The number of pixels corresponds to the total number of pixels associated with BML regions (Section 5.4.4, Section 5.5.3). In parentheses, we report the total number of BML detections (Section 5.4.3), including those ruled out as potential BML regions as described in Section 5.4.4.

The velocity estimate v¯c is a signal-weighted averaged of the velocity centroid from all BML detections toward each SNR (Section 5.4.4). The velocity-width ∆v is an average of each BML detected toward the SNR. The distances d and d are the distance to each SNR as deter- vc Table 5.1 mined from the rotation curve at v¯c and the distance reported in Table 5.1, respectively. When calculating the former value, we prefer distances that are closer to the latter in order to resolve the distance ambiguity. identiﬁed BML regions, we give the total number of pixels with BML identiﬁcations in parentheses in Table 5.2.

5.4.4 BML Region Identiﬁcation

As described in Section 5.4.3, local maxima in the spectrum at each pixel are flagged if the line profile calculated at that point matches the outlined criteria. In this way, we can construct a separate “flagged” data cube with dimensions equal to the original data cube, ∗ that is, the antenna temperature Ta(x, y, v). This “flagged” data cube has F (x, y, v) = 1 if the spectrum at position (x, y) and in the velocity channel v meets our criteria, and F (x, y, v) = 0 otherwise. 114

We can further break down the “flagged” data cube into regions by using a group- finding algorithm simultaneously in spatial coordinates and velocity-space. For the purposes of this analysis, the definition of BML regions will ultimately depend on the physical significance of these groups. As described in Section 5.3, our 12CO J = 2 − 1 observations were performed with a 10m diameter telescope at 230 GHz with a FWHM beam of 3300. Each data cube was constructed by taking observations spaced 1000 apart. In our analysis, detection of pixels with broad lines whose centroids lie within each others’ FWHM and within a radius of 2 pixels represents a single BML region. For BML regions that appear extended on the sky, we can define groups that span several pixels, provided the broad-line detection is approximately at the same velocity (i.e., within a window spanning ∆v measured for each BML). A velocity centroid is measured for each BML detection representing a local maximum in the 12CO spectrum. However, each of these points is biased by the way in which we ∗ construct the original data cube Ta , and the determination of a group of BML detections must be physically significant. In the same way, the velocity centroid of each local maximum must be tied in some way to the actual velocity centroid for the entire BML region with which it is associated. Therefore, we utilized a signal-weighted average to measure this value, for which the BML region velocity centroid (v¯c) is defined by the number of ∗ local maxima it includes (N), the Gaussian amplitude of each local maximum (Ta,i), and the velocity centroid calculated (vc,i as described in Section 5.4.3)

N P ∗ Ta,ivc,i v¯ = i=1 (5.4) c N P ∗ Ta,i i=1 For newly discovered BML regions toward SNRs, we report in Table 5.2 the number of BML pixels we detect toward each remnant, the signal-weighted velocity centroid as described above, and the average velocity-width of all BML detections toward the SNR. Finally, we recalculated the distance to each remnant using the signal-weighted velocity centroid as a prior for the Clemens rotation curve and we report this distance in Table 5.2. 115

In order to resolve the distance ambiguity, we chose the distance closest to our original estimate from Table 5.1. The agreement between the distances reported in Table 5.1 is variable, although nearly always within the 50% uncertainty on the original estimate given the prior we used to ﬁnd associated CO emission (Section 5.4.1). In some cases, these values are very close to the original distances, especially for G08.3−0.0, G18.6−0.2, 4C−04.71, and G29.6+0.1, where the distances were estimated from the Σ − D relation for the ﬁrst three objects and

X-ray H I column density for the latter.

5.4.5 BML Region Examples: 3C 391 and Cas A

We applied our algorithm to two test cases involving known SNR-MC systems: 3C 391 (Section 5.4.5) and Cas A (Section 5.4.5). In this way, we demonstrate the types of BML regions selected by our algorithm and the accuracy of the measurements. These two systems present unique problems that must be addressed in identifying BML regions in a large sample of 12CO J = 2 − 1 data. Here, we discuss previously detected signatures of BML regions toward each remnant and the underlying SNR-MC interactions they may trace. From these examples, we examine the extent to which our algorithm can be optimized.

3C 391

Several observations of 12CO targeted on 3C 391 have revealed the presence of broadened molecular features from +90 to +110 km s−1 (Sanders et al., 1986; Wilner et al., 1998; Reach and Rho, 1999). The emission exhibits significant broadening relative to the ambient gas, with 12CO J = 2 − 1 line widths exceeding 20 km s−1 in some regions (Figure 5.2). This emission has consistently been interpreted as a signal of postshock gas. The molecular line profiles also exhibit a bright, narrow component (∆v ∼ 2 km s−1) at the same velocities and peaking in the same region. This observation suggests an association between molecules of significantly different temperatures and densities and presents a challenge when identifying regions of shocked molecular gas. One hypothesis for the presence of the narrow component is that it arises from cooler, less turbulent gas reforming behind a dissociative shock (Reach and Rho, 1999). That is, while the narrow component 116

may be associated with the SNR shock, in some regions it could be superimposed upon and dominate the line proﬁle from the broadened emission. Under this hypothesis, we might only detect BML regions where the SNR-MC interaction is ongoing or recent. Thus, detection of recently shocked gas in 12CO may be extremely sensitive to the criteria for identifying individual line components (i.e., local maxima) in the spectrum. In our analysis,

we selected all the local maxima in velocity-space whose signal exceeds 10 Trms. In this way, our criteria are unbiased by the coincidence of narrow and bright line components with broad and faint components as in 3C 391. One caveat is that false detections can occur where the local maxima are close together in velocity-space; for example, we might report no detection if the narrow component was within 2 spectral channels of the broad component. In this case, additional information on the line profile would be necessary to calculate an accurate line width. While this scenario does not occur in 3C 391, one could imagine dissociated molecular gas reforming behind a shock whose narrow emission peaks at roughly the same systemic velocity as the undissociated, turbulently broadened component. As a follow-up to our initial analysis, we modified our requirement for the five channels to include scenarios where the line profile does not decrease monotonically for two channels on either side of the local maxima. While we still required five velocity channels, we extended the channel window on either side of the line profile to avoid contamination from other emitting regions in velocity- space. As we demonstrate in Figure 5.2 (left), the well-studied, shocked MC toward 3C 391 is located to the south of the remnant along a region of bright radio continuum. The 12CO spectrum in this region is significantly broadened in a feature centered around a systemic velocity of +112 km s−1 as seen in Figure 5.2 (right). Our data is of sufficient spatial and spectral resolution to resolve and confirm the narrow and broad spectral features toward this cloud. Another important point that arises when examining similar BML regions is that the +112 km s−1 line appears to lie outside the velocity range inferred for an object at l = +31.9 and d = 7.2±3.6 km s−1, and the broad profile itself is at least 15 km s−1 from the maximum velocity allowed within this range. The remnant appears to be near the tangent 117

point of the Galactic rotation curve at this velocity (+102 km s−1), but the MC is slightly redshifted with respect to this velocity. Given that the shock interaction with the MC has been conﬁrmed in multiple wavelengths and is generally assumed to be associated with the remnant, one hypothesis for the velocity discrepancy is that remnants may accelerate MCs in an ongoing interaction. This point underscores the need for care in treating SNR- MC interactions whose distances are near the tangent point of the rotation curve along a particular line-of-sight. As we describe in Section 5.4.1, we extend the velocity window used as a prior in our CO spectra by 15 km s−1 at either end to account for this potential problem toward other SNRs.

Cas A

Cas A exhibits subtle indications of a SNR-MC interaction in several ways. Warm CO gas (T ∼ 20 K) directly to the west of the SNR implies some source of energy to heat the cloud (Wilson et al., 1993). There is also a correlation between the brightest portions of the radio shell at 1 GHz, nonthermal X-ray emission, and the radio spectral indices in the same region suggesting that the ejecta are encountering a denser medium (Anderson and

Rudnick, 1996; Keohane et al., 1996). In the near-infrared Ks band, strong nonthermal emission occurs in the same regions (Rho et al., 2003). These observations are consistent with an interaction scenario, although they do not conclusively point to the presence of a SNR-MC interaction. Recent results indicate that hard X-ray emission above 15 keV toward Cas A is dominated by the western half of the remnant (Grefenstette et al., 2015). Knots of bright emission in nonthermal emission are consistent with steep spectral indices in the X-ray and some change in the physical environment. These knots are coincident with bright radio emission and the large MC to the west of Cas A. Further evidence of an ongoing interaction scenario may come from changes in this and other nonthermal emission on short timescales.

In the MC itself, H2CO absorption studies suggest that certain species are broadened to as much as 5.5 km s−1 (Reynoso and Goss, 2002). Similarly, slight broadening in 12CO J = 1−0 up to 4.5 km s−1 is observed (Liszt and Lucas, 1999), while line widths in 12CO 118

J = 2 − 1 range from 6.2 − 6.7 km s−1 (Kilpatrick et al., 2014). Thus, evidence for a SNR-MC interaction toward Cas A is strongly correlated in the X-ray, radio, near-infrared, and sub-millimeter, but the overall result in 12CO J = 2 − 1 is subtle. In order to detect BML regions toward similar SNR-MC interactions, we must rely on a precise constraint in the velocity-width. Optimizing this constraint might be achieved by comparison with narrow molecular emission in the same field. Line widths of narrow components in the Cas A field range from 2.5 − 3.4 km s−1, or about half of the “broadened” line widths. Therefore, one cri- terion for “broadened” emission might be: can we divide these emission lines into distinct “narrow” and “broad” categories? For example, how sensitive is the identification of BML regions to the velocity-width constraint we applied in Section 5.4.3? We illustrate this point in Figure 5.3 (right) where we compare broadened features to unbroadened emission as determined by our analysis. The regions with the broadest 12CO line widths are qualitatively different from regions with narrow molecular lines. Furthermore, we report the total number of pixels selected by the BML algorithm for newly identified BML regions in Table 5.2, which includes pixels that were ruled out as potential BML regions by the criteria described in Section 5.4.4. For newly identified BML regions, the algorithm selects these pixels with little contamination from other positions. That is, the total number of BML pixels identified using the criteria described in Section 5.4.3 is generally 75 − 85% the number of BML pixels we determine are associated with the SNR as described in Section 5.4.4. We interpret this result to mean that a 6 km s−1 velocity-width threshold may be close to optimal for identifying BML regions toward SNRs and therefore minimizes false positives from spurious detections of, for example, overlapping molecular lines.

5.4.6 Comparison to Other BML Region Indicators

Jiang et al.(2010) compare molecular line studies to other traditional detection methods for SNR-MC interactions. These include detection of the OH 1720 MHz maser, morphological correspondence between a SNR and a MC, molecular line emission with a large high-to-low line ratio such as 12CO J = 2 − 1/J = 1 − 0, detection of near-infrared emis- 119

Figure 5.3: (Left) Integrated 12CO J = 2 − 1 emission toward Cas A from −43 to −35 km s−1. Orange contours are 20 cm radio continuum emission of the remnant from Anderson et al.(1991). Pixels with broadened molecular emission are shown as white squares. We denote a comparison region of unbroadened, bright molecular emission with a larger red square. (Right) The average spectrum from these pixels at the velocity of the BML identiﬁcation (solid black) along with a spectrum of the unbroadened, bright comparison region averaged over the pixels denoted by the red square (dashed red). Shaded parts of the spectrum correspond to velocities ruled out by the inferred distance to the SNR. 120

sion from shock excitation such as [Fe II] or H2 emission, and infrared colors consistent with molecular shocks. It is valuable to utilize as many of these indicators as possible, and in our analysis we rely on morphological correspondence between the SNR in radio emission and the MC as well as reports of OH 1720 MHz maser emission. However, the detection of broadened 12CO line emission is a unique tool for tracing SNR-MC interactions in our galaxy. Spectroscopic information is vital for identifying the systemic velocity of molecular gas and therefore associating the kinematic distance to molecular emission with distances to SNRs. This information reduces the probability of confusion by any chance alignment of a shock tracer with a SNR. Carbon monoxide is also the second most abundant molecule and possesses one of the highest bond-dissociation energies of any diatomic molecule. Therefore, we would expect CO to be more easily visible and less susceptible to dissociation behind a SNR shock than OH, H2 or any other molecular shock tracer. Finally, broad, low-J CO emission is a direct tracer of disturbed gas in a MC, as opposed to line emission from atomic species or near-infrared colors which do not directly or unambiguously trace the molecular gas. Probing the MC for molecular shocks and then drawing an association with the SNR is well-suited for ﬁnding SNR-MC interactions, even in a large survey.

5.5 Results

We applied the BML region algorithm to the 12CO J = 2 − 1 spectra (Figure 5.1) toward all 50 SNRs in our sample to determine the location of broadened CO features. The algorithm revealed broad-line detections toward nineteen SNRs, including nine previously unidentiﬁed SNR-MC systems. For the other ten systems with coincident BML regions, we discuss our detections toward known SNR-MC systems in the context of previous studies in Section 5.5.1. One of the remaining 31 SNRs, G54.1+0.3, had previous indications of a BML region. We argue in Section 5.5.2 that the associated MC may not be velocity- broadened. Finally, in Section 5.5.3, we discuss the results of our algorithm for newly identiﬁed SNR-MC systems and compare these detections to the observed properties of the SNRs themselves. For reference, we include a discussion of previous observations for some remnants in Appendix D, where we discuss details of the radio continuum, OH 121

maser detections, and pulsar observations.

5.5.1 SNRs with Previous Detections Supporting SNR-MC Interactions

BML regions were detected toward ten SNRs with known SNR-MC systems. Here we compare the characteristics of the detected BML regions to results from previous studies. G16.7+0.1 was observed to have a SNR-MC interaction by Reynoso and Mangum (2000). The authors report detection of a small MC in 12CO J = 1 − 0 at +25 km s−1 and coincident with OH maser detections to the south of the remnant. The CO emission −1 3 from this cloud is seen from +25.1 to +25.9 km s . The cloud is small (2.1 × 10 M , = 260 −3 1.5 −1 nH2 cm ) and has a CO line width of only km s , consistent with the line width of an ambient MC approximately 3 pc in size (Heyer and Brunt, 2004). Reynoso and Mangum(2000) report detection of two other MCs toward G16.7 +0.1, which they call the “central” and “northwestern” clouds, at systemic velocities between +25 and +27 km s−1. The central cloud is located directly to the east of the brightest 5 GHz enhancement in the remnant, implying a possible connection between nonthermal emission and the MC. The reported morphological correspondence, small size of the cen- 3 −1 tral cloud (1.8 × 10 M ), and moderately broadened CO profile (∆v ∼ 4.4 km s ) are all characteristic of a SNR-MC interaction. The northwestern cloud appears along the sharp boundary of the SNR shock front. While this evidence may indicate physical contact between the SNR ejecta and molecular gas, there is little indication in 12CO J = 1−0 emission to suggest an interaction scenario. The cloud extends to the northwest of the remnant as part of a larger complex, its observed line width (∆v ∼ 2.7 km s−1) is only slightly larger than the surrounding gas, and no other signs of SNR-MC interaction appear in this region. This potential interaction region may be supported by the presence of enhanced TeV gamma-ray emission, although the detection is just below the 5σ limit traditionally used to identify TeV sources (Table 5.1). In 12CO J = 2 − 1, we detect the same molecular emission as seen in 12CO J = 1 − 0 emission toward G16.7+0.1 at a systemic velocity of +25 km s−1. Toward the emission that Reynoso and Mangum(2000) identify as the “southern cloud,” some pixels appear 122 velocity-broadened in 12CO J = 2 − 1 to a width of around ∆v = 9.6 km s−1. This value contrasts with the smaller velocity-width in J = 1 − 0 of 1.5 km s−1 and suggests a recent or fast shock may be responsible for the large observed line widths. While this width is broad compared to emission observed in 12CO J = 1−0, it is also consistent with other shock indicators. The southern cloud is the site of OH maser emission as well as a significant enhancement in the radio continuum. We do not detect any broadened emission at the positions of the other clouds in this region. Kes 75 was first observed to have a BML region in 12CO J = 1 − 0 (Su et al., 2009). A large MC is coincident with the northern half of the remnant and at a systemic velocity between +45 and +58 km s−1. The evidence for interaction is in the broadened blue wing of this feature (+45 to +51 km s−1). This asymmetric and broadened feature is spectroscopically similar to the SNR-MC interactions seen toward IC 443 and 3C 391. Kes 75 is also the site of a strong TeV gamma-ray source (Djannati-Ata¨ı et al., 2008), although this emission is thought to be associated with a pulsar wind nebula in the remnant itself. We confirm the BML detection in 12CO J = 2 − 1 around +53 km s−1. Broadened molecular emission appears extended along the north and west of the remnant and is coincident with the brightest regions of 20 cm continuum emission. This observation is also consistent with detection of broadened emission in 12CO J = 1−0 where line widths were observed to be on the order of 3.7 km s−1. In our analysis, we find broader line widths in 12CO J = 2 − 1 up to 9.2 km s−1. 3C 391 is discussed in Section 5.4.5. Kes 79 was observed to have a SNR-MC interaction in bright 12CO J = 1−0 emission to the east and south of the radio shell of the remnant. This morphological correspondence between the SNR and MC strongly implies an association between these objects (Stanimirovic´ et al., 2003). The CO line profile at these positions is centered around +103 km s−1 with an integrated intensity of 75.9 K km s−1 and a peak brightness temperature of ∼ 8 K. Thus, the line has an approximate equivalent width of ∼ 9 km s−1, consistent with a shock-broadened cloud. However, the observation of 12CO J = 1−0 was performed with the NRAO 12 m telescope and a beamwidth of 4500. The integrated intensity measurements reported in Stanimirovic´ et al.(2003) are ambiguous and may include 123

line emission from multiple, ambient MCs and not a single, shocked cloud. We confirm detection of velocity-broadened 12CO emission in J = 2 − 1 from +99 to +109 km s−1 toward Kes 79. Broad line detections extend across the MC from southwest toward the northeast of the remnant with velocity-widths up to 8.8 km s−1. Agreement between our 12CO measurements and previous observations, as well as the OH measurements at similar velocities (Appendix D), argues for the presence of a SNR-MC interaction toward this remnant. 3C 396 is observed to have a possible SNR-MC interaction in 12CO J = 2 − 1 and J = 1 − 0 emission, where a bright cloud appears between +67 and +72 km s−1 (Su et al., 2011). At these velocities, the coincident MC extends around the remnant, with emission on the blue end of the line profile concentrated to the north and east and extending clockwise around the remnant toward redshifted emission along the western edge. This detection represents strong morphological agreement between the structure of the MC and the SNR. The spectral lines themselves exhibit signs of broadening, with an asymmetric profile and broad wings toward redder velocities in the eastern emission and blueshifted broad wings in the southwestern emission. In some regions, these line profiles have widths > 10 km s−1. The presence of broadened molecular emission along the western edge of 3C 396 as well as the OH maser at the same velocity argues for the presence of a SNR-MC interaction. We detect velocity-braodened 12CO J = 2 − 1 emission centered around +69 km s−1 to the north of the remnant and extending to the west where the broadened emission is centered around +77 km s−1. The detected broad line widths are generally between 6.6 and 7.6 km s−1. These broad-line detections are not bright or significantly enhanced compared to surrounding material. While our analysis supports the presence of a SNR- MC interaction, detailed analysis of this region may reveal shocked material that is not significantly distinct from the surrounding gas. Evidence for a SNR-MC interaction toward 3C 397 comes from 12CO J = 1 − 0 observations performed with the 13.7 m telescope at the Purple Mountain Observatory at Delingha (Jiang et al., 2010). These data reveal a MC at +32 km s−1 and to the north and west of the remnant. At these positions, the SNR shock appears to be embedded in 124 the MC and morphological correspondence between the 12CO and SNR radio continuum emission imply the objects are interacting. The CO emission in these regions indicates that the MC emits a line profile with broad blue wings (+28 to +31 km s−1), consistent with acceleration from a SNR-MC interaction. The authors report a FWHM line width of 2.5 km s−1 toward the MC at this velocity. We confirm detection of velocity-broadened emission along the western edge of the remnant and centered at a systemic velocity of +31 km s−1. The blue wing is not as bright in 12CO J = 2 − 1 emission as in the lower-excitation line, although the overall line profile appears to be enhanced to ∆v = 6.2 km s−1. Toward W49, evidence for a SNR-MC interaction can be seen in 13CO J = 1 − 0 emission (Simon et al., 2001). Broad lines have been detected at a systemic velocities from 0 to +20 km s−1 in multiple components and with a narrower line profile centered at +16 km s−1. Averaged over the entirety of W49B, the line profile has a width of ∼ 10 km s−1, although the spatially resolved emission is likely narrower. We detect broadened 12CO J = 2 − 1 emission toward W49B in the same velocity range and centered at a systemic velocity of +14 km s−1. Regions of broadened emission are seen to the west of W49B, and the majority of our detections come from a cloud located along the southwestern boundary of the remnant. In this region, we report BML detections with widths of ∆v = 10 km s−1. While this finding agrees with previous detections of broadened molecular emission, some interpretations argue against the presence of a SNR-MC interaction to the southwest of W49B. Lacey et al.(2001) found that low-frequency radio emission in this region is almost completely attenuated, and the intervening molecular material responsible for absorption may only be seen toward the remnant in projection. However, at higher frequencies W49B becomes significantly brighter. Indeed, at 90, 20, and 6 cm, the southwestern portions of the shell show significant local enhancement (Moffett and Reynolds, 1994). Added to this fact is the observation that spectral indices appear to flatten to the southwest of W49B. It may be that this region is not enhanced in nonthermal emission. Overall, the evidence for the presence of a SNR-MC interaction is ambiguous, though our observations support the presence of an interaction scenario. Detailed characterization of the 125 physical properties in the gas itself is needed to further verify the presence of a SNR-MC interaction. Cas A is discussed in Section 5.4.5. Evidence for a potential SNR-MC interaction toward HB 3 was reported by Routledge 12 et al.(1991), in which the authors perform H I line mapping and CO J = 1 − 0 observations. Toward this remnant, a bright 12CO “bar” is coincident with a region of enhanced radio emission near l = +134◦, b = +1◦ and extending east in Galactic coordinates. This morphological correspondence suggests the SNR is interacting with the cloud and a region of enhanced particle acceleration is seen in the SNR ejecta. Evidence from Shi et al. (2008) supports this hypothesis where radio continuum from 1.4 to 4.8 GHz is brighter toward the 12CO cloud. The spectral indices at these frequencies appear to steepen to α = −0.61. Emission from the MC itself extends to the south of the remnant, roughly in the direction of the OH maser detection, and the line profile peaks at a velocity of −43 km s−1. −1 However, the W3 H II region is embedded in a MC at −40.5 km s , and it is certain that the 12CO “bar” at −43 km s−1 is associated with this cloud. Bieging and Peters (2011) detect bright 12CO J = 2 − 1 and J = 3 − 2 emission centered on W3(OH) in 12 this same region. The correlation between the H II region, high line-ratio CO emission, and 1.1 mm dust continuum from the Bolocam Galactic Plane Survey (Rosolowsky et al., 2010) argues that these BML regions are not associated with HB 3. We detect 12CO J = 2 − 1 emission in this region and centered at a systemic velocity of −40 km s−1. Our spectral analysis confirms that this emission is broadened, up to ∆v = 12.7 km s−1, although generally to a velocity width of ∆v = 7.6 to 8 km s−1 in 12CO J = 2 − 1. IC 443 is a well-observed SNR with several indications of a SNR-MC interaction. This interaction was first identified in broadened CO emission by Cornett et al.(1977) and has been confirmed by detection of molecular species toward the associated MCs (White et al., 1987; Xu et al., 2011), OH 1720 MHz emission (Claussen et al., 1997; Hewitt et al., 2006), and H2 emission (Denoyer, 1979b,a; Burton et al., 1988). The physical conditions toward IC 443 are well-suited for every shock interaction tracer traditionally used to identify SNR- 126

MC interactions. As a result, this region is one of the most luminous molecular hydrogen sources in the galaxy. The interacting MC generally exhibits linewidths of > 30 km s−1, up to 90 km s−1 toward the most turbulent molecular gas (White et al., 1987). This result has been conﬁrmed in several molecules, probing molecular gas of densities up to 3 × 106 cm−3 (van Dishoeck et al., 1993). The SNR itself has a composite morphology and appears evolved, with an inferred age of ∼ 2 × 104 (Leahy, 2004; Lee et al., 2008). The overall radio morphology of the remnant was described by Braun and Strom(1986) in 327 MHz continuum emission. Two main subshells (shells A and B) make up the majority of the remnant in radio continuum emission. Shell A appears brighter and is coincident with the molecular shock indicators along its southern boundary and across the center of the remnant as a whole. The western edge of this shell appears to be associated with both the OH maser emission and TeV gamma-rays (Albert et al., 2007b). Given its approximate age and the assumed distance of ∼ 1.5 kpc, IC 443 has a large apparent size on the order of 1◦. We targeted a section of the remnant toward its center and along the southwestern edge of shell A. The shocked 12CO emission peaks around −17 km s−1 in this region and extends blueward in a broad, asymmetric line proﬁle. Our observations centered on a bright cloud referred to as shock clump B in Denoyer(1979b). We measured velocity-widths in 12CO J = 2−1 toward this region ranging from ∆v = 7.6 to 28 km s−1.

5.5.2 Non-Detection Toward G54.1+0.3

G54.1+0.3 was targeted with 12CO J = 2 − 1 observations in Lee et al.(2012). A large MC at +23 km s−1 extends along the north of the remnant, with a bright central core to the northeast of the remnant. The morphological correspondence between the outer blastwave of the SNR and the MC is not strong, with at least 3000 separation between the 20 cm contours of the remnant and the MC. Lee et al.(2012) report a FWHM of 7 km s−1 at +23 km s−1 across the entire cloud. Our observations indicate that the broad feature detected around +19 km s−1 with ∆v ∼ 11 km s−1 (Lee et al., 2012, see Figure 13, therein) is at least two clouds blended 127

Figure 5.4: (Left) Integrated 12CO J = 2 − 1 emission toward G54.1+0.3 from +16 to +23 km s−1. Orange contours are NVSS 1.4 GHz continuum emission of the remnant. Pixels with broadened molecular emission would be shown with white squares, but none passed our criteria. We denote a region of the +23 km s−1 cloud discussed in Lee et al. (2012) with a red square. (Right) The average spectrum across the pixels denoted by the red square. There appear to be two line components centered at +20 and +23 km s−1. Neither component exhibits broadening at the > 6 km s−1 level we used in our analysis. over the 3000 grid spacing applied in this study. While the broad and narrow components of emission at these velocities can be separated from +15 to +30 km s−1, the broad component is further separated into two features centered at +20 and +23 km s−1 (Fig- ure 5.4). These components are resolved in our spectra. However, these features may have been smoothed into a single, broad-line component in Lee et al.(2012) given the larger beam (4800) from the 6-m Seoul Radio Astronomy Observatory Telescope. In contrast to the Gaussian component analysis over a large region, our pixel-by-pixel analysis of the molecular emission reveals no indication of > 6 km s−1 velocity-broadening.

5.5.3 Newly Identiﬁed SNR-MC Systems

We detected broadened molecular features toward nine SNRs with no previous indication of SNR-MC interactions: G08.3−0.0, G09.9−0.8, G11.2−0.3, G12.2+0.3, G18.6−0.2, G23.6+0.3, 4C−04.71, G29.6+0.1, G32.4+0.1. In this section, we discuss the characteristics of the interaction as identiﬁed by our BML region algorithm along with previous 128 work on the associated SNRs. Our results from the detection of velocity-broadened regions toward each remnant are summarized in Table 5.2. For comparison with each remnant, we use 90 cm radio continuum contours from Brogan et al.(2006) (hereafter, BGGKL06) where l < +22◦ and NVSS 1.4 GHz contours (Condon et al., 1998) for the remaining objects.

G08.3−0.0

G08.3−0.0 was initially identified in 20 cm emission by Helfand et al.(2006) and subsequently confirmed by BGGKL06 at 90 cm as a small (50 × 40) SNR with shell-like morphology. The remnant has flux density 1.0 Jy at 20 cm, but its morphology suggests an enhancement in surface brightness toward the east. A subsequent search for associated OH maser emission at 1612, 1665, 1667, and 1720 MHz with the Green Bank Telescope (GBT) yielded a single-dish detection with no velocity information or localization (He- witt and Yusef-Zadeh, 2009). Of particular note is a TeV gamma-ray detection reported toward this remnant (Higashi et al., 2008). The gamma-ray emission was observed with a 70 offset from the center of G08.3−0.0 and detected with ∼ 140 resolution. Given the position of the remnant relative to the gamma-ray detection, there is reason to believe that the gamma-ray source originates from the northern half of this SNR. In projection, the radio contours of G08.3−0.0 appear to flatten toward the northwest (Figure 5.5) coincident with a MC between −0.7 and 4.6 km s−1. The spectrum of the molecular emission in this region appears velocity-broadened with 12CO Gaussian widths fit between 6.5 and 9.0 km s−1. The coincidence of this broadened 12CO along with the morphology of the remnant, the possible OH maser toward this position, and the gamma- ray source associated with this region support the presence of a SNR-MC interaction at this location at a systematic velocity near +2.6 km s−1. We note that the velocity centroid of the BML region and the kinematic distance inferred from this velocity, d = 16.4 kpc, is also in good agreement with the value given by the Σ − D relation (16.3 kpc). 129

Figure 5.5: (Left) Integrated 12CO J = 2 − 1 emission toward G08.3−0.0 from −2 to +6 km s−1. Orange contours are 90 cm radio continuum emission of the remnant from BGGKL06. Pixels with broadened molecular emission are shown as white squares. We denote a comparison region of unbroadened, bright molecular emission with a larger red square. (Right) The average spectrum from these pixels at the velocity of the BML identi- ﬁcation (solid black) along with a spectrum of the unbroadened, bright comparison region averaged over the pixels denoted by the red square (dashed red). Shaded parts of the spectrum correspond to velocities ruled out by the inferred distance to the SNR. 130

G09.9−0.8

G09.9−0.8 has been observed to have typical radio shell-like morphology (BGGKL06). Its spectral index between 90 and 20 cm is Crab-like (α = −0.4) implying nonthermal emission from the SNR. There may also be confusion with thermal emission from other sources, which suggests that the SNR has a steeper spectral proﬁle in agreement with enhanced particle acceleration. Hα emission toward the remnant is correlated with nonthermal emission at 1.4 GHz along the northern edge of the remnant (Stupar and Parker, 2011). This trend implies a connection between regions of strong SNR-ISM interaction and nonthermal emission from the SNR along that edge of the shell. We observe an association between a bright region of 90 cm radio emission and a MC toward the southwest of this remnant (Figure 5.6). The molecular emission appears at a systemic velocity of +27 to +33 km s−1 and extends in a small arc toward the center of the remnant. The region of velocity-broadened emission is concentrated toward a peak in the 90 cm contours, and the molecular emission to the east of this peak exhibits the most velocity-broadening. The 12CO Gaussian widths of the velocity-broadened emission in this region are ﬁt between 6.6 and 9.1 km s−1.

G11.2−0.3

G11.2−0.3 was identified as a SNR by Shaver and Goss(1970) at 408 and 5000 MHz. These observations revealed a small remnant (40) with a steep spectral index of α = −0.52. Follow-up with ASCA indicated there is a 65 ms pulsar toward the center of the remnant (Torii et al., 1997). Subsequent radio detection of the remnant in BGGKL06 revealed a probable shell-like SNR with a flux density at 20 cm of approximately 0.6 Jy. Emission from the remnant at 90 cm appears locally enhanced around the shell, and especially to the east and northwest. Detection of OH maser emission has been reported from GBT observations (Green et al., 1997), however, the lack of precise position and velocity information makes it difficult to associate the OH maser with the SNR. There appears to be a strong TeV gamma-ray source extended to the northwest of the remnant and HESS detects emission toward G11.2−0.3 at 5.5σ (Bochow, 2011). 131

Figure 5.6: (Left) Integrated 12CO J = 2 − 1 emission toward G09.9−0.8 from +26 to +34 km s−1. Orange contours are 90 cm radio continuum emission of the remnant from BGGKL06. Pixels with broadened molecular emission are shown as white squares. We denote a comparison region of unbroadened, bright molecular emission with a larger red square. (Right) The average spectrum from these pixels at the velocity of the BML identiﬁcation (solid black) along with a spectrum of the unbroadened, bright comparison region averaged over the pixels denoted by the red square (dashed red). Shaded parts of the spectrum correspond to velocities ruled out by the inferred distance to the SNR. 132

Figure 5.7: (Left) Integrated 12CO J = 2 − 1 emission toward G11.2−0.3 from +30 to +38 km s−1. Orange contours are 90 cm radio continuum emission of the remnant from BGGKL06. Pixels with broadened molecular emission are shown as white squares. We denote a comparison region of unbroadened, bright molecular emission with a larger red square. (Right) The average spectrum from these pixels at the velocity of the BML identiﬁcation (solid black) along with a spectrum of the unbroadened, bright comparison region averaged over the pixels denoted by the red square (dashed red). Shaded parts of the spectrum correspond to velocities ruled out by the inferred distance to the SNR.

We detect velocity-broadened molecular emission toward this remnant from +29 to +36 km s−1. At these velocities, emission extends from the southeast to the north of the remnant, roughly coincident with the bright shell of G11.2−0.3 as seen in 90 cm contours (Figure 5.7). Regions of velocity-broadened emission lie mostly toward the northern molecular emission, with some detections interior to the main radio shell. The Gaussian line width of velocity-broadened emission in this region is between 6.6 and 9.1 km s−1.

G12.2+0.3

G12.2+0.3 was discovered in BGGKL06 as a small, partial shell-like remnant. The radio spectral index from 90 to 20 cm is steep with α ∼ −0.7. The weak ﬂux density at 20 cm combined with the size suggest that the remnant has d > 10 kpc, and the Σ − D relationship implies a distance of 15.6 kpc. 1720 MHz OH maser emission has been detected toward this remnant in single-dish observations with no reported velocity information or 133

Figure 5.8: (Left) Integrated 12CO J = 2 − 1 emission toward G12.2+0.3 from +48 to +55 km s−1. Orange contours are 90 cm radio continuum emission of the remnant from BGGKL06. Pixels with broadened molecular emission are shown as white squares. We denote a comparison region of unbroadened, bright molecular emission with a larger red square. (Right) The average spectrum from these pixels at the velocity of the BML identification (solid black) along with a spectrum of the unbroadened, bright comparison region averaged over the pixels denoted by the red square (dashed red). Shaded parts of the spectrum correspond to velocities ruled out by the inferred distance to the SNR. localization (Hewitt and Yusef-Zadeh, 2009). We detect velocity-broadened 12CO emission coincident with this remnant at a systemic velocity of +50 km s−1 (Figure 5.8). The broadened molecular emission lies toward the eastern shock front of the remannt and to the north of the brightest region in 90 cm continuum emission. This emission is weakly detected in two pixels of our map, and the velocity-width between the two BML detections extends from 6.6 to 7.2 km s−1. However, the morphology of the coincident MC and SNR are suggestive of an interaction at this velocity. As seen in projection, the shape of the SNR appears flattened to the east and along the boundary of the SNR. This shape contrasts with the northwest edge of the SNR where the shell structure appears much fainter and more extended. We infer that the eastern shock front is deccelerated relative to the rest of the remnant, which would require a significant density enhancement in the ISM. 134

Figure 5.9: (Left) Integrated 12CO J = 2 − 1 emission toward G18.6−0.2 from +40 to +48 km s−1. Orange contours are 90 cm radio continuum emission of the remnant from BGGKL06. Pixels with broadened molecular emission are shown as white squares. We denote a comparison region of unbroadened, bright molecular emission with a larger red square. (Right) The average spectrum from these pixels at the velocity of the BML identiﬁcation (solid black) along with a spectrum of the unbroadened, bright comparison region averaged over the pixels denoted by the red square (dashed red).

G18.6−0.2

G18.6−0.2 was first confirmed as a SNR by BGGKL06 in 90 cm continuum emission, although Helfand et al.(2006) listed the source G18.6375-0.2917 as a “high-probability supernova remnant candidate” before this confirmation. Subsequent detection at 20 cm identified a partial shell with a spectral index of α = −0.3. The remnant is small (∼ 60) and faint in 90 cm emission (1.9 Jy), in agreement with a far distance as determined from the Σ − D relation (13.2 kpc). As shown in Figure 5.9, the region of brightest radio continuum emission toward this remnant extends in an arc to the east with fainter lobes of emission to the southwest and north. We find evidence for a SNR-MC interaction to the north of the remnant at a systemic velocity of +44 km s−1. The MC in this region covers most of the northern and western half of the SNR, although the disturbed region appears localized to the north. The velocity- widths of emission lines detected toward this cloud are broad, with most lines exceeding ∆v = 8.0 km s−1. 135

G23.6+0.3

It has been suggested based on its high infrared to radio luminosity ratio that G23.6+0.3 is

an H II region and not a SNR (Pinheiro Gonc¸alves et al., 2011). However, this object has a Crab-like spectral index (α = −0.34), which implies some nonthermal component to the emission from this source (Shaver and Goss, 1970). While the radio emission is elongated, it correlates well with the spatial extent of the infrared emission and suggests that the two

are associated. This observation is consistent with the possibility that G23.6+0.3 is an H II region and SNR projected onto each other. We detect velocity-broadened lines toward the remnant at a systemic velocity of +91 km s−1. The regions of broadened molecular emission lie along the brightest part of the remnant’s northern shell (Figure 5.10). Emission in this region appears moderately broadened, and the emission lines have velocity-widths around ∆v ∼ 6.5 km s−1 up to 7.5 km s−1. As seen in projection, the morphology of the SNR relative to the MC is highly suggestive. There appears to be a boundary along the northern edge of G23.6+0.3 where an ongoing interaction has broadened the coincident MC.

4C−04.71

4C−04.71 is often referred to as Kes 73 and was identified as a potential SNR in part of a larger radio source centered around Galactic coordinates l = +27.3, b = 0.0 by Milne (1969). The remnant is compact with a central pulsar identified at X-ray wavelengths and an age of ∼ 2000 yr (Vasisht and Gotthelf, 1997). Green et al.(1997) were able to localize OH maser emission toward this source 120 west of the SNR around +33 km s−1. 13 Subsequent observations revealed associated CO J = 1 − 0 and H I emission toward −1 this source and two nearby H II regions imply a connection with gas around +110 km s (Tian and Leahy, 2008b). HESS detected TeV gamma-ray emission toward this remnant at 6σ significance (Aharonian et al., 2008b). There appears to be a TeV source extending to the south of the remnant, although several other counterparts have been discussed for TeV detections in this region. We observe broadening from the same molecular emission toward 4C−04.71 as seen 136

Figure 5.10: (Left) Integrated 12CO J = 2 − 1 emission toward G23.6+0.3 from −43 to −35 km s−1. Orange contours are NVSS 1.4 GHz continuum emission of the remnant. Pixels with broadened molecular emission are shown as white squares. We denote a comparison region of unbroadened, bright molecular emission with a larger red square. (Right) The average spectrum from these pixels at the velocity of the BML identiﬁcation (solid black) along with a spectrum of the unbroadened, bright comparison region averaged over the pixels denoted by the red square (dashed red). Shaded parts of the spectrum correspond to velocities ruled out by the inferred distance to the SNR. 137

Figure 5.11: (Left) Integrated 12CO J = 2 − 1 emission toward 4C−04.71 from +94 to +102 km s−1. Orange contours are NVSS 1.4 GHz continuum emission of the remnant. Pixels with broadened molecular emission are shown as white squares. We denote a comparison region of unbroadened, bright molecular emission with a larger red square. (Right) The average spectrum from these pixels at the velocity of the BML identiﬁcation (solid black) along with a spectrum of the unbroadened, bright comparison region averaged over the pixels denoted by the red square (dashed red). Shaded parts of the spectrum correspond to velocities ruled out by the inferred distance to the SNR. in Tian and Leahy(2008b), although in 12CO J = 2 − 1 the line proﬁle of the broadened emission is centered around +100 km s−1 with average line widths of ∆v = 9.0 km s−1 (Figure 5.11). However, the broadened regions lie at large separations from 4C−04.71 (∼ 30). It is possible that these BML regions toward 4C−04.71 are associated with other

nearby radio sources (Sanbonmatsu and Helfand, 1992, see, e.g.,) often assumed to be H II regions. These sources would need to extend to within 30 in projection toward 4C−04.71 to its west and southwest, which is not seen in the radio continuum. A second possibility, which has been raised in the context of Cas A (Kilpatrick et al., 2014), is that faint, fast-moving ejecta from 4C−04.71 could interact with the molecular gas without being visible in radio emission. If ejected at high-velocity during early phases of SN evolution, this material could carry enough momentum to turbulently accelerate the MC to the velocity width we observe. This hypothesis is supported by the observation that 4C−04.71 is still relatively young and the turbulent clouds are nearby as seen in projection. Presumably some fraction of the ejecta has not yet been signiﬁcantly deccelerated by the 138

surrounding ISM and would not need to travel far to shock the observed MCs.

G29.6+0.1

G29.6+0.1 was discovered in VLA observations at 6 and 3.7 cm by Gaensler et al.(1999). The radio source is a faint, shell-like remnant, and several local enhancements were detected in the continuum. The spectral index from 6 to 3.7 cm is roughly α ∼ −0.5 across the remnant, although the brightest region of 6 cm continuum toward the eastern shell is steeper, with α = −0.77. The northwestern edge of the shell also appears relatively bright, although there is no evidence of steeper nonthermal emission. Subsequent observations toward G29.6+0.1 with ASCA revealed a compact X-ray source thought to be a pulsar (Vasisht et al., 2000). The X-ray source is relatively faint, and no timing or age estimate have been performed, although the coincidence of this source strongly implies that G29.6+0.1 is a distant and possibly young SNR from a core-collapse SN. We detected broadened 12CO emission toward G29.6+0.1 and roughly coincident with a region of bright 20 cm continuum emission, at a velocity of +94 km s−1 (Figure 5.12). This position is consistent with enhancements in the 6 cm continuum and spectral index, as well as a region of bright X-ray emission as mapped by Vasisht et al.(2000). The majority of the molecular emission at these velocities coincides with the remnant, implying a strong morphological association between the SNR and MCs. A weak red wing is also visible in the 12CO spectrum demonstrating the characteristic velocity proﬁle of shocked-broadened molecular emission. The two pixels where we ﬁnd evidence for velocity-broadened molecular emission are highly disturbed, with line widths ∆v ∼ 10.0 km s−1.

G32.4+0.1

G32.4+0.1 was discovered in the NRAO/VLA Sky Survey (NVSS) at 1.4 GHz (Condon et al., 1998). The remnant has clear shell-like structure visible in both radio and X-ray wavelengths (Yamaguchi et al., 2004). While the morphology of the remnant and its X-ray spectrum imply a strong synchrotron component, a radio spectral index has not yet been measured. Of particular note are two lobes seen in nonthermal radio and X-ray emission 139

Figure 5.12: (Left) Integrated 12CO J = 2 − 1 emission toward G29.6+0.1 from +34 to +42 km s−1. Orange contours are NVSS 1.4 GHz continuum emission of the remnant. Pixels with broadened molecular emission are shown as white squares. We denote a comparison region of unbroadened, bright molecular emission with a larger red square. (Right) The average spectrum from these pixels at the velocity of the BML identiﬁcation (solid black) along with a spectrum of the unbroadened, bright comparison region averaged over the pixels denoted by the red square (dashed red). Shaded parts of the spectrum correspond to velocities ruled out by the inferred distance to the SNR. 140

Figure 5.13: (Left) Integrated 12CO J = 2 − 1 emission toward G32.4+0.1 from +39 to +47 km s−1. Orange contours are NVSS 1.4 GHz continuum emission of the remnant. Pixels with broadened molecular emission are shown as white squares. We denote a comparison region of unbroadened, bright molecular emission with a larger red square. (Right) The average spectrum from these pixels at the velocity of the BML identiﬁcation (solid black) along with a spectrum of the unbroadened, bright comparison region averaged over the pixels denoted by the red square (dashed red). along the eastern edge of the remnant (Ueno et al., 2005). These features appear to be local enhancements in nonthermal emission (Figure 5.13). Moreover, they are adjacent to regions where the radio shell appears to break and little or no emission can be seen. This morphology contrasts with the western edge where the radio shell appears more contiguous albeit no brighter than the eastern lobes. We detect broadened molecular emission toward a cloud at a systemic velocity of +43 km s−1 and roughly coincident with the eastern radio lobe of G32.4+0.1. The associated MC extends along the southeastern and eastern part of the shell, although we detect velocity-broadening only along the northernmost portion of this cloud where the radio continuum of G32.4+0.1 is brightest. The Gaussian velocity-width of molecular emission in this region ranges from 7.1 to 8.3 km s−1. The association with the nonthermal enhancement toward G32.4+0.1 is highly suggestive, and the spatial correlation between the MC and the shell of the remnant imply an interaction scenario. However, the kinematic distance inferred from this association 141

is in tension with the values inferred from the Σ − D relation (18.5 kpc) and X-ray H I absorption (22 kpc) (Ueno et al., 2005). The Σ − D estimate may be affected by the low

ﬂux density of G32.4+0.1 (S1GHz ∼ 0.25 Jy from Green(2014)), which renders surface brightness and size estimates uncertain. Finally, it may be that the density of hydrogen −3 toward this source is signiﬁcantly larger than the nH = 1.0 cm assumed for the latter estimate.

5.6 Discussion

5.6.1 The SNR-MC Interaction Rate

Reynoso and Mangum(2001) argued that, because half of all SNRs are expected to result from Type II SNe and their progenitors are young, massive stars born in large MC complexes, it is expected that roughly half of SNRs should be in contact with MCs. New information on the incidence of core-collapse SNe places additional constraints on the expected rate of this type of SNR-MC interaction. Here, we discuss the total number of SNR-MC interactions in our galaxy as a fraction of all known Galactic SNRs, which we call the “SNR-MC interaction rate.” We examine this value in light of our ﬁndings and the factors that may contribute to an enhanced or reduced rate. Recent transient surveys have greatly improved our understanding of the intrinsic fraction of both thermonuclear and core-collapse SN types in volume-limited samples. For example, the Lick Observatory Supernova Search (LOSS) observed 175 SNe between 60 and 80 Mpc and determined that 57% of SNe are Type II, another 19% are Type Ib/c and the remaining 24% are Type Ia (Li et al., 2011), implying that roughly 3/4 of SNRs come from core-collapse SNe. We might expect the SNR-MC interaction rate to be close to this value. However, discounting the four objects in our sample mentioned in Section 5.3, our fraction of candidate interactions (17/46 or 37%) is much lower than the observed fraction of core-collapse SNe. If we rule out possible Type Ia SNe by looking only at SNRs with known pulsars or compact X-ray sources, which we infer resulted from core-collapse SNe, the total rate appears to change somewhat (7/15 or 47%). We show in Section 5.6.2 142

below that virtually all SNR-MC interactions identified by other means also show broad CO and will be identified by our approach. However, in Section 5.5.2 we illustrated the potential for false identifications. Therefore, 7/46 and 7/15 should be interpreted as upper limits. The difference between these two rates also supports the hypothesis that SNR-MC interactions tend to be associated with core-collapse SNe, but both values are low compared to the expected rate. There appears to be an underlying physical mechanism that suppresses the number of SNRs that interact with MCs compared to the known fraction of core-collapse SNe or the number of SNR-MC interactions detectable in CO lines. In Section 5.6.2, we discuss the possibility of ongoing or recent interactions between SNRs and MCs that do not excite broadened 12CO J = 2 − 1 emission. We examine the possibility that some underlying physical mechanism suppresses the SNR-MC interaction rate in Section 5.6.3 and in Section 5.6.4 we discuss the implications of a low SNR-MC interaction rate for sequential star formation and simulations of SN feedback. Finally in Section 5.6.5, we analyze the correlation between SNR-MC interactions and TeV gamma-ray emission from SNRs.

5.6.2 To What Extent Are SNR-MC Interactions Detectable in 12CO J = 2 − 1?

One of the most fundamental questions in our 12CO J = 2 − 1 analysis of BML regions toward SNRs is what fraction of SNR-MC interactions are detectable in molecular emission? Jiang et al.(2010) argue that OH maser detections are a reliable (if incomplete) signpost of SNR-MC interactions, so a test of the completeness of 12CO J = 2 − 1 line broadening detections for SNR-MC interactions might be “what fraction of SNRs with nearby OH maser emission have also been detected in broad-line molecular emission?” In Table 5.3, we list remnants with OH maser detections and whether molecular line studies, where they exist, support the presence of broadened molecular lines. While all of the studies cited in Table 5.3 detect molecular features at roughly the same systemic velocity as the coincident OH maser emission, the incidence of broad-line detections is 13/15 or 87%. Another study is inconclusive: Lazendic et al.(2002) conﬁrm the presence of shocked molecular hydrogen emission and broad-line absorption toward G359.1−0.5 but their results are ambiguous for the CO lines, so we have not included this 143

Table 5.3. SNRs with Detected OH Maser Emission

Cat. No. Transitions in Previous Studies Line Broadening? Reference

G0.0+0.0 CS (1−0, 5−4, 7−6) Y 1,2 G1.05−0.1 G1.4−0.1 G5.4−1.2 CO (1−0) N 3 G5.7−0.0 CO (1−0) N 3 G6.4−0.1 CO (1−0, 3−2) Y 4 G8.7−0.1 G9.7−0.0 G16.7+0.1 CO (1−0, 2−1) Y 5,6 G21.8−0.6 CO (1−0), 13CO (1−0), C18O (1−0), Y 7 HCO+ (1−0) G27.4+0.0 CO (2−1) Y 5 G31.9+0.0 CO (2−1), CS (2−1, 3−2, 5−4), Y 5,8 HCO+ (1−0) G32.8−0.1 CO (1−0, 2−1) 13CO (1−0) Y 9 G34.7−0.4 CO (2−1) Y 10 G49.2−0.7 CO (1−0, 2−1), 13CO (1−0), Y 11 HCO+ (1−0) G189.1+3.0 CO (1−0, 2−1) Y 5,12 G337.0−0.1 G337.8−0.1 G346.6−0.2 G348.5−0.0 G348.5+0.1 CO (1−0) Y 6 G349.7+0.2 CO (1−0) Y 13 G357.7−0.1 CO (1−0, 2−1, 4−3),13CO (1−0, 2−1), CS (2−1, 3−2), Y 14 + HCO (1−0),HCN (1−0),H2CO (3(2,2)−2(2,1), 3(0,3)−2(0,2)) G357.7+0.3 G359.1−0.5 CO (1−0, 2−1), 13CO (1−0, 2−1),C18O (1−0, 2−1), CS (2−1, 3−2),HCO+ (1−0), ? 15

HCN (1−0),H2CO (3(2,2)−2(2,1), 3(0,3)−2(0,2)),

H2CO (3(2,2)−2(2,1), 3(0,3)−2(0,2)),SiO (2−1, 5−4)

Note. — SNRs are selected from those with identiﬁed OH maser emission features in Frail et al.(1996); Green et al.(1997); Yusef-Zadeh et al.(1999); Hewitt et al.(2008); Hewitt and Yusef-Zadeh(2009). For remnants where studies have also been performed in molecular line emission, we give the transitions observed. We also indicate whether the study concluded that the observed emission exhibits line-broadening (LB) relative to the surrounding gas (Y) or there was no evidence of line broadening (N). For references that performed molecular line studies: (1) (Serabyn et al., 1992); (2) (Tsuboi and Miyazaki, 2012); (3) (Liszt, 2009); (4) (Arikawa et al., 1999); (5) this study; (6) (Reynoso and Mangum, 2000); (7) (Zhou et al., 2009); (8) (Reach and Rho, 1999); (9) (Zhou and Chen, 2011); (10) (Seta et al., 1998); (11) (Koo and Moon, 1997); (12) (Cornett et al., 1977); (13) (Reynoso and Mangum, 2001); (14) (Lazendic et al., 2004);a (15) (Lazendic et al., 2002)b

aOnly the observations mentioned in 12CO J = 4 − 3 toward G357.7-0.1 have broad-line detections, although these data have not been published.

bBroad-line absorption is seen in CO, CS, and HCO+ while broad-line emission appears in 13CO. See Section 5.6.2. 144

object. Molecular line broadening toward shocked MCs may not be detected for several reasons. Molecular gas from a single cloud can be optically thick and obscure the shocked regions. The faint line wings characteristic of broadened molecular lines may also be obscured by bright, narrow emission from other clouds. These effects are usually more severe in lower energy transitions where line emission from quiescent clouds is intrinsically broad. Indeed, both of the BML non-detections in Table 5.3 were performed in 12CO J = 1 − 0 and toward SNRs near the Galactic Center (G5.4-1.2, G5.7-0.0) (Liszt, 2009). Observations in higher energy transitions might show less emission from quiescent clouds along the line-of-sight while revealing broad-line emission. Every observation we report that was performed in transitions at higher energies than 12CO J = 1 − 0 has found signs of broadened molecular emission. We infer that broad molecular line emission is a highly complete (e.g., > 87%) tracer of SNR-MC interactions, and 12CO J = 2 − 1 may exhibit broadened molecular lines in almost all instances of a SNR-shocked MC. Detection of broadened 12CO J = 2 − 1 lines may be a necessary, but not sufﬁcient, condition for the presence of SNR-MC interactions. This hypothesis would imply the limiting factor in assessing the “true” number of SNR- MC interactions in 12CO J = 2 − 1 is not false negatives from optically thick or obscured emission, but false positives from multiple MCs in velocity-space that give the appearance of a broad line or from BML regions toward other sources such as H II regions. We report evidence of both effects, for example toward HB 3 where the coincident BML region

appears to be associated with the nearby H II region W3(OH), G54.1+0.3 where the line emission from multiple velocity components appears to have been mistaken for broad-line emission in previous work, and 4C−04.71 where BML regions are detected at multiple velocities implying contribution from other sources. Given the high completeness we infer for detections of SNR-MC interactions observed in 12CO J = 2−1, the ratio of BML regions we detected to observed SNRs (i.e., 37%) may only be an upper limit to the “true” SNR-MC interaction rate in the galaxy. A lower SNR- MC interaction rate only exacerbates the discrepancy with the fraction of core-collapse SNe. We infer that some additional mechanism is required in order to suppress the number 145

of SNR-MC interactions occurring in the galaxy, and we explore possibilities and their consequences below.

5.6.3 Suppressing the SNR-MC Interaction Rate

Suppression of the SNR-MC interaction rate relies on a SN event that is too distant from a MC for interaction to occur on a short timescale. For core-collapse events, we propose that there are two main mechanisms by which suppression can occur. One possibility is that massive stars are kicked or migrate away from their parent MC such that they are too distant to cause a SNR-MC interaction. The second is a delay between the SN event and SNR-MC interaction, which we can express as some characteristic timescale or, equiva- lently, a characteristic separation between massive stars and MCs. These mechanisms are not mutually exclusive, and indeed, the former might provide an explanation for the delay between SN explosions and SNR-MC interactions. The evolution of the SNR-MC interaction rate with age provides a way of disentangling each possibility and its underlying cause. For example, one explanation of the low SNR-MC interaction rate may be a large, distinct population of runaway massive stars. Tetzlaff et al.(2011) detected 2547 runaway star candidates in Hipparcos, including over 200 stars with peculiar velocities from 10 − 120 km s−1 toward OB associations. For a star migrating directly away from its parent cloud at 10 km s−1 over 106 yr, a subsequent core-collapse SN would explode roughly 10 pc from its starting position. Thus, the time between explosion and SNR-MC interaction would be delayed by approximately 6000 yr (i.e., the age at which a SNR has a radius of 10 pc). Age is a critical factor in the SNR-MC interaction rate. We hypothesize that the SNR- MC interaction rate for SNRs of varying ages may be solely a function of the “delay” between the explosion and interaction. This supposition is supported by the observation that the incidence of runaway stars is appreciable compared to the overall population of young, nearby stars. Tetzlaff et al.(2011) found that the runaway frequency of OB stars is 27% for their sample. This population of stars may explode as SNe and interact with MCs only as evolved SNRs or never exhibit signatures of SNR-MC interactions. 146

Assuming massive stars are “kicked” or otherwise migrate a large distance from their parent MC, an analysis of the SNR-MC interaction rate with SNR age would provide a constraint on the delay caused by this migration. A constant SNR-MC interaction rate with age would suggest that there are two distinct populations of massive stars; those that remain close to their parent MC and those that migrate a large distance such that all known Galactic SNRs are too young for the latter to have encountered a MC. The more likely scenario is that the SNR-MC interaction rate increases with SNR age, which suggests that massive stars migrate a short distance (i.e., < 30 pc) from their parent MC over their lifetimes. We would expect the SNR-MC interaction rate to level off at a certain SNR age. This analysis may provide a constraint on the distribution of massive star runaway velocities and the fraction of massive stars that are runaways. We note several SNRs with age determinations from free-expansion velocity or pulsar timing (Appendix D). However, the number of such SNRs is too low to draw any statistically meaningful conclusions. As more SNR-MC interactions are identiﬁed, accurate age determinatons may be obtained from kinematic distances and modeling of SNR size, morphology, and surface brightness. Vink(2012) provides an analytic overview of age determinations that accounts for SN parameters as well as momentum and radiative loss. Once ages have been determined for a sufﬁcient number of SNRs, an estimate of the dependence of the SNR-MC interaction rate with age should be possible.

5.6.4 Astrophysical Implications of a Low SNR-MC Interaction Rate

Another rich field of study is the direct impact SNRs may have on star formation. Models of sequential or triggered star formation rely on an understanding of the feedback between massive stars and star forming regions. While the relationship between OB associations and star-forming regions has long been recognized (Elmegreen and Lada, 1977; Herbst and Assousa, 1977), the contribution of shocks from SNRs is more ambiguous. Some studies have proposed that these shocks will compress molecular regions and accelerate triggered star formation (Tomisaka et al., 1981), although it is possible the shock velocities involved are dissociative and will destroy the molecular gas instead. There is also evidence that SNRs play little or no role in regulating star formation in individual star clusters due 147 to the time delay before SNe occur, but turbulent outflows powered by SNe may suppress star formation in giant MCs or even galaxy-wide star formation (Krumholz et al., 2014). Recent numerical simulations of SN feedback have examined the impact SN feedback has on galactic environments in order to evaluate the importance of SNRs in star formation (e.g., Hopkins et al., 2012; Agertz et al., 2013; Aumer et al., 2013; Stinson et al., 2013; Hopkins et al., 2014). Many of these simulations fail to produce the observed stellar masses due to excessive cooling or the lack of momentum transfer from SNe to giant MCs. This problem may be exacerbated by a large fraction of SNe that explode far from any molecular gas. Over short timescales, a large fraction of SNe may play no role in galaxy-wide feedback. In the absence of momentum transfer from SNRs, it is possible that stellar feedback mechanisms other than SNe are dominant. For example, H II regions are known to pho- toionize dense MCs and may suppress star formation where a sufficient level of background radiation is present. This hypothesis is supported by observations of giant H II regions (such as 30 Doradus in the LMC) indicating that radiation pressure from stars is the dominant mechanism in adding energy and momentum to the surrounding region (Lopez et al., 2011). The relative role of each mechanism must be assessed in the context of specific timescales and physical conditions. Our observations support arguments that SNRs may contribute less to stellar feedback than previously thought.

5.6.5 Gamma-Ray Emission from SNR-MC Interactions

Particle acceleration is one of the most important galactic processes in which SNRs play a central role. Cosmic rays are produced in SNRs where the SN explosion can deliver enough energy to rapidly accelerate nuclear particles (Berezhko et al., 1996), and the energy density of cosmic rays produced in the shock region increases with the density of the surrounding medium (Jun and Jones, 1999). Enhanced turbulence and magnetic ﬁelds can excite signiﬁcant hadronic particle acceleration which produces π0 mesons that decay immediately into gamma-rays. Consistent with this hypothesis, there are a number of instances where TeV gamma-rays are reported from regions of interaction between SNRs and MCs (e.g., Aharonian et al., 2006b; Albert et al., 2007a; Acciari et al., 2009). 148

A challenge to such studies is the complex and dense population of MCs and the high density of TeV gamma-ray sources in the inner Milky Way where most known TeV sources lie. There may also be a tendency to target regions of possible interaction in searching for gamma-ray sources, potentially introducing a bias. These issues can be mitigated in our survey. We use the positions of the specific SNRs targeted in the CO survey as priors to identify possible TeV counterparts, which we accept if their positions coincide with the SNRs within 100. This positional tolerance was adopted from a conservative estimate of the positional accuracy of 60 for the TeV measurements (e.g., Hinton and the HESS Collab- oration, 2004, Felix Aharonian, private communication), plus an allowance for the extent of the SNR. We can then compare the number of detections for the SNRs where we detect broad CO lines indicative of SNR-MC interactions with the number where we do not find broad lines. We adopt the conventional threshold of 5σ for a TeV detection. Hahn(2014) has evaluated the HESS Milky Way survey (the source of nearly all the TeV detections we use) and shown that all of the sources at this detection level have been identified, that is, the sample is complete to this level of significance. We correct the detection rates for chance coincidences by identifying TeV sources with the same positional tolerance relative to 120 positions selected arbitrarily in the same region. Where there were possible identifications, we report the separation between the approximate location of the SNR-MC interaction region or the center of the SNR for remnants without SNR-MC interactions and the nearest TeV gamma-ray source in Table 5.4. We have taken positions for the TeV sources from TeVCat (Wakely and Horan, 2008). There are nine cases within our positional criteria, four with indications of MC interactions and five without. Our test on arbitrary positions indicates that there are about two false identifications, one for the cases with interactions and one for those without. That is, there are three true cases coinciding with interactions and four not coinciding. We have estimated the number of SNRs with TeV measurements from Bochow(2011), as shown in Table 5.1. The two sample sizes are virtually the same (19 without interactions, 20 with), so the rates of detections are identical within statistical errors. The similar number of associations between SNRs and gamma-ray emission in these two samples does not support a strong correlation between SNR-MC interactions and TeV 149

Table 5.4. SNRs with Possible TeV Gamma-Ray Emission at 5σ Signiﬁance

Cat. No. α (J2000) δ (J2000) BML Detection γ-ray Separationa γ-ray Source Ref. (h m s) (d m s) arcmin

G08.3−0.0 18 04 28 −21 47 00 D 5.1 HESS J1804-216 Aharonian et al.(2006b) G11.2−0.3 18 11 25 −19 23 20 D 22 HESS J1809-193 Renaud et al.(2008) G12.8−0.0 18 13 37 −17 49 00 1.4 HESS J1813-178 Aharonian et al.(2006b) G17.0−0.0 18 21 57 −14 08 00 57 HESS J1825-137 Aharonian et al.(2006c) G18.1−0.1 18 24 34 −13 11 00 42 HESS J1825-137 Aharonian et al.(2006c) G21.5−0.9 18 30 50 −10 09 00 17 HESS J1831-098 Sheidaei(2011) G27.4+0.0 18 41 18 −04 54 20 D 39 HESS J1841-055 Aharonian et al.(2008b) G29.7−0.3 18 46 26 −02 57 30 D 1.5 HESS J1846-029 Djannati-Ata¨ı et al.(2008) G74.9+1.2 20 16 02 +37 12 00 0.1 VER J2016+371 Aliu(2011) G111.7−2.1 23 24 15 +58 48 20 D 7.9 TeV J2323+588 Becker et al.(1991) G120.1+1.4 00 25 18 +64 09 00 2.1 TeV J0025+641 Acciari et al.(2011) G130.7+3.1 02 05 41 +64 49 00 2.0 TeV J0209+648 Aleksic´ et al.(2014) G184.6−5.8 05 34 31 +22 01 00 0.2 TeV J0534+220 Aharonian et al.(2006a) G189.1+3.0 06 17 24 +22 26 20 D 8.4 TeV J0616+255 Acciari et al.(2009)

Note. — SNRs with nearby TeV gamma-ray detections at the 5σ significance level. We report the position of the detected BML region from our study for SNRs with BML detections or the location of the SNR itself for those without BML detections from Green(2014). The gamma-ray separation is determined from the difference between the reported position and the nearby TeV gamma-ray source. The gamma-ray source identifiers are also given with the reference for the position of each source indicated. aOnly cases with separation < 100 are accepted for our study. 150 gamma-ray emission. The observed gamma-ray detections may also be contaminated by particle acceleration in pulsars, which are known gamma-ray sources. For example, the Crab pulsar has long been associated with hadronic particle acceleration and TeV gamma- ray emission (Hayakawa, 1958; Fazio et al., 1968; Ostriker and Gunn, 1969). TeV gamma- ray surveys typically distinguish between “compact” and “extended”’ gamma-ray sources, which may aid in resolving the ambiguity between gamma-rays from pulsars and more extended regions along the SNR shell. However, this distinction may fail to separate pulsar point sources from SNR-MC interactions where the conditions for particle acceleration are met only in a small region. Moreover, H I clouds are observed near SNRs with a high enough density gradient to excite particle acceleration at shock fronts (Fukui et al., 2012; Xing et al., 2015). In the absence of a pulsar, it may not be necessary to have a SNR-MC interaction to create a gamma-ray source toward a SNR, although the contact discontinuity between SNR ejecta and the densest parts of the ISM would still be expected to produce the brightest and hardest gamma-ray sources. We can account for these uncertainties to an extent. Ten out of the fourteen SNRs we report with TeV gamma-ray emission at 5σ significance have pulsars or coincident, compact X-ray sources, including five of the six objects with BML regions and five of the eight without. Removing these sources leaves only G08.3−0.0 and Tycho as the two objects with TeV gamma-ray sources detected at close separations, and while we report a coincident BML region toward the former, we do not detect any strong 12CO emission toward the latter. G08.3−0.0 is a recently discovered object (Section 5.5.3) and it is possible that an undiscovered pulsar exists in the SNR and provides the particle acceleration necessary to generate TeV gamma-rays. Higashi et al.(2008) discuss several possible sources for the gamma-ray detection in this region, including the pulsar PSR B1800-21, although this source appears to be separated from the gamma-ray emission by at least 100. Barring these possibilities, G08.3−0.0 represents the best candidate in our sample for TeV gamma-ray emission from a SNR-MC interaction. In addition to G08.3−0.0, compelling evidence for TeV gamma-ray emission can be seen toward a number of SNR-MC interactions, notably W28 (Aharonian et al., 2008c) and IC 443 (Acciari et al., 2009). Although these cases show that SNR-MC interactions can 151 produce TeV sources in some cases, the lack of a strong correlation in our study shows that other processes must play an important, perhaps dominant, role. Larger statistical samples are needed to evaluate the association of TeV sources with SNR-MC interactions more definitively. Bochow(2011) lists all measurements of SNRs, detections and upper limits from HESS. Out of the nineteen SNRs we observed with evidence for SNR-MC interaction, twelve also have evidence for gamma-ray emission at the 2σ significance level including six of the nine SNRs with newly reported interactions. This fraction is appreciably higher than the overall incidence of gamma-ray emission toward SNRs, which is 24/50 in our sample. Previous systematic studies for gamma-ray emission toward SNRs have found an even lower ratio of SNRs with coincident TeV gamma-ray emission (1/3 in Bochow, 2011). At this level, our findings seem to support the hypothesis that TeV gamma-ray emission is correlated with SNR-MC interactions. However, confirmation of gamma-ray emission at this level is problematic. Detections at 2σ may be close to the confusion limit with background sources along the Galactic plane. In addition, the diffuse emission of the Galaxy can contribute to such detections without being associated with a specific source. Combined with the positional uncertainty in gamma-ray experiments, the correlation between SNR-MC interactions and TeV sources at 2σ is tantalizing but not necessarily reliable. We conclude that there is no convincing evidence with the existing data that SNR-MC interactions are a dominant source of TeV gamma-ray emission from SNRs. Other sources of particle acceleration in SNRs, especially pulsars and interaction with the H I gas in the cold neutral medium, may contribute significantly to the number of gamma-ray detections.

5.7 Conclusion

Our survey of 12CO J = 2 − 1 molecular spectroscopy toward 50 SNRs is the ﬁrst large, systematic search for this shock tracer toward SNRs. We demonstrate that our technique for analyzing 12CO J = 2 − 1 data cubes satisfactorily reproduces broad-line detections in systems with known BML features. In applying this technique to our data set, we have found: 152

1. BML features were detected toward nine remnants (G16.7+0.1, Kes 75, 3C 391, Kes 79, 3C 396, 3C 397, W49B, Cas A, HB 3, IC 443) with previous evidence of SNR-MC interactions and nine remnants (G08.3−0.0, G09.9−0.8, G11.2−0.3, G12.2+0.3, G18.6−0.2, G23.6+0.3, 4C−04.71, G29.6+0.1, G32.4+0.1) with no previous evidence of interaction. The BML region detected toward HB 3 is likely

associated with the H II region W3(OH). We are uncertain whether the BML regions associated with 3C 396 and W49B represent SNR-MC interactions, and detailed follow-up toward these clouds is required to draw any ﬁrm conclusions. For one of the remaining remnants, G54.1+0.3, we ﬁnd that previous indications of a SNR-MC interaction may have been due to the blending of multiple narrow line components.

2. From our analysis of SNRs with signs of molecular shocks from OH masers, we infer broadened 12CO J = 2 − 1 features are a necessary but not sufﬁcient condition to determine the presence of SNR-MC interactions. Conﬁrmation must come from other sources, such as 12CO line ratios, to demonstrate that the gas in BML regions is physically distinct from the surrounding material. The total number of SNR-MC interactions we observed as a fraction of observed SNRs may only be an upper limit to the “true” SNR-MC interaction rate in the galaxy.

3. The fraction of detected BML regions toward SNRs with known pulsars or compact X-ray sources (7/15) is slightly larger than the overall number (17/46), supporting the hypothesis that SNR-MC interactions are associated with core-collapse SNe from massive stars. However, the number of BML regions we detect toward SNRs is signiﬁcantly lower than the incidence of core-collapse SNe. We infer that age may be a critical factor in the overall rate of SNR-MC interactions. A low SNR- MC interaction rate may also have far-reaching consequences for models of stellar feedback and sequential star formation.

4. We found no convincing evidence for a correlation between TeV gamma-ray sources and SNR-MC interactions. The number of TeV detections within 100 of SNRs with BMLs is the same as the number for SNRs with no evidence for interactions. We conclude that SNR-MC interactions are not a dominant source of TeV gamma-rays 153 in the galaxy. As an alternative, pulsars or interaction with H I gas in the cold neutral medium may contribute signiﬁcantly to the detection rates for TeV gamma-rays toward SNRs. 154

CHAPTER 6

Conclusion

6.1 Results

In this dissertation, I sought to apply the physics of SN shocks described in Chapter 1 to observations of a range of physical systems involving SNe and SNRs. Through these observations, I addressed a few key science questions: when SN exhibit strong shocks with CSM, how can we decompose their optical spectra to glean information about the underlying SN?; what does nonthermal emission from the young SNR Cas A tell us about the validity of diffusive shock acceleration, especially as it pertains to magnetic ﬁeld ampli- ﬁcation?; how often do shock interactions between SNRs and MCs occur in our galaxy?; are these molecular shocks toward SNRs associated with TeV gamma-rays and sequential star formation as has been predicted in previous studies? Chapter 2 describes the young SN PS15si, which exhibits signatures of CSM interaction in the form of narrow emission lines of hydrogen. I demonstrated that this SN belongs to a class of objects with spectroscopic signatures of SNe Ia and that spectral decomposition of PS15si into SN Ia combined with blackbody emission from CSM interaction does not accurately reproduce the physical parameters of the combined SN Ia/IIn spectrum. While overluminous SN 1991T-like spectra provide the most consistent match to PS15si as has been observed toward other SNe Ia/IIn, the CSM temperatures implied by this spectral template are too low to account for the luminosity of the interaction-powered SN. The 1991T-like spectrum may in fact be an artifact of the shock interaction between the SN ejecta and CSM, which produces ultraviolet/X-ray radiation that may change the ionisation state of the visible ejecta layers. This interaction would also produce a quasicontinuum with distinct spectral features and explain why the SN spectrum is easily visible at a time when the CSM should outshine the SN by several magnitudes.

Observations of nonthermal Ks band emission from Cas A described in Chapter 3 155

reveals that there may be synchrotron-emitting knots entrained in the ejecta observed at infrared wavelengths. I use a novel technique to constrain the magnetic field and electron density of emission observed over multiple epochs. Using this technique, I find magnetic field strengths significantly larger than predicted downstream from the forward shock. This observation argues in favor of magnetic field amplification as predicted by Reynolds (1998). I examine Cas A in greater detail in Chapter 4 and address a major question from the previous chapter: what is unique about the southwest shock of Cas A that could generate bright synchrotron-emitting knots and enhance magnetic field amplification? This study uses mid-infrared spectroscopy to examine the northern shock where a molecular cloud appears coincident with the shock front. By using the dispersion of Doppler velocities in mid-infrared fine-structure lines, I argue that there are subtle indications of a shock interaction in this region. Examining the remnant as a whole in 12CO J = 2 − 1 emission reveals a large molecular cloud to the west and south of the remnant and the CO lines to the southwest are statistically broader than lines in other regions. I argue that this signature points toward a shock interaction between Cas A and the molecular cloud, which may also be responsible for radio and X-ray emission enhancements and steep spectral indices associated with nonthermal emission. Curiously, there are also indications of molecular line broadening at least 10 in projection away from the southwestern Cas A shock front, although along the northeast-southwest bipolar outflow from the remnant. I suggest that the energetics in this broadening agree with a scenario where the molecular gas has been shocked by fast-moving ejecta from the SNR. Finally, in Chapter 5 I extend the use of line broadening to a survey of 50 Galactic SNRs in which I examined 12CO J = 2 − 1 emission to look for signatures of SNR-MC interactions. It had been predicted in previous studies that the incidence of these interactions should be comparable to the overall fraction of core-collapse SNe (roughly 3/4). I found that 19 out of 50 remnants exhibited signatures of broad molecular lines, which is appreciably smaller than expected. I interpreted these results to indicate that SNRs play a much smaller role in stellar feedback and sequential star formation than previously predicted. Moreover, localizing regions where SNRs shock MCs gave me an opportunity to 156 test predictions related to physical processes that occur along SNR-MC interfaces, especially particle acceleration and gamma-ray emission, I examined the incidence of Galactic TeV gamma-ray emission and compared locations of these events to SNR-MC interactions. I found that TeV sources toward SNRs are not strongly correlated with SNR-MC interactions. This finding argues that SNR-MC interfaces are not the dominant source of TeV gamma-rays in our galaxy.

6.2 Future Direction

6.2.1 Late-time Radio Emission from Supernovae

Recent detailed studies of SNe IIn spectra have raised new questions about the latest phases of stellar evolution in massive stars. As discussed in Chapter 1 and Chapter 2, SNe IIn are thought to be caused by slow-moving circumstellar gas that is hit by the blast wave or illuminated by UV radiation from the SN. In most cases, the presence of these narrow lines agrees with a scenario in which a large amount of CSM is ejected by the SN progenitor star before core collapse. In some extreme cases, as much as 25 M has been ejected in the decade before core collapse (Smith et al., 2007b, 2008b, 2010; Woosley et al., 2007), resulting in the most luminous SNe yet recognized in the universe. SNe IIn events are about 8-9% of all core-collapse SNe (Smith et al., 2011), so violent precursor eruptions evidently precede some SNe, but not all. So far we have no explanation for the cause of these eruptions. A key question is whether these eruptions are actually timed to go off just before the SN - for example, due to instabilities in the final burning stages - or if instead SNe IIn stand out because of a selection effect. Because nearby material is illuminated when the SN is still young, if there was recently ejected material it will be observable during the prime observational time of transient searches, and the SN will be classified as a SN IIn. If, on the other hand, any major mass-loss episodes are less recent, it may take months or years for the SN to interact with the CSM. Most often this will go undetected as follow- up observations rarely extend past the first few months. This introduces a bias in our interpretation of pre-SN mass loss. Given this uncertainty, key science questions that must 157 be addressed are: Is strong pre-SN mass loss really limited to just the few years before core-collapse in SNe IIn? If so, what are the physical implications for models of massive star evolution? If pre-SN eruptions are limited to just a few years before core collapse, that would point to a violent instability in the latest O, Ne, or Si burning phases. If these eruptions also occur in the centuries before core-collapse, then perhaps massive stars undergo an instability in C- burning. If there is evidence for pre-SN eruptions several millennia before core-collapse, then other mechanisms must be explored. Do other types of core-collapse SNe (e.g., Type II-L SNe) show signs of strong CSM interaction when the ejecta reach the CSM at late times? This is still very unclear as most SNe are not followed with deep, panchromatic observations after they fade. A SN might not be categorized as Type IIn upon discovery but still have distant CSM that will interact with the SN ejecta at late times. Of particular interest, Type II-L SNe (SNe II-L) progenitors suffer weaker pre-SN mass loss than SNe IIn progenitors and there is spectroscopic evidence that a continuum exists between these two types (Smith et al., 2015b). For these SNe, there must still be significant pre-SN mass loss to remove most or all of their hydrogen envelopes despite a lack of evidence for dense CSM when the SN is young. Could weaker mass loss from SNe II-L progenitors immediately preceding the SN imply more significant mass loss during earlier burning phases? If only SNe IIn show this evidence for mass loss during earlier burning stages (i.e., when a more evolved SN encounters distant CSM ejected during these stages), then perhaps this phenomenon only occurs in very massive stars, which might then point to a near-Eddington (radiation pressure) instability. This evidence would also raise significant new questions about the continuity and strength of mass loss in SNe II-L progenitors. Over the past few years, several radio searches have looked for emission toward SNe IIn and II-L. These studies have generally focused on a limited number of targets or on emission within the first few years after the initial SN event. While the feasibility of late-time monitoring of SNe IIn and II-L has been assessed and radio emission has been detected toward these objects - notably for 1995N (Chandra et al., 2009) - so far no systematic studies have been published. I am currently undertaking a radio survey to evaluate 158

1.4 GHz target range 6.0 GHz 10. GHz target range

σ VLA 3− RMS limit

Figure 6.1: (Right) Multi-frequency (1.4, 6, 10 GHz) light curves of a model SN at 30 −6 −1 −4 −1 Mpc with constant 5 × 10 M yr mass-loss and a 2 × 10 M yr eruption timed 7000 yr before the SN explosion. (Left) The corresponding spectral index between 6 and 1.4 GHz. the presence of radio synchrotron emission in the late-time light curves of SNe IIn and II-L. In this way, I will constrain the histories of a large sample of SN progenitors and conclude whether instabilities in earlier burning phases contribute signiﬁcantly to pre-SN mass loss. Models of radio SN radio light curves must account for emission due to the synchrotron radiation from accelerated electrons in CSM interactions (Chevalier, 1982) as well as attenuation due to features of the local CSM, internal free-free and self-absorption, and any additional absorption along the line of sight unassociated with the SN. The former processes are affected by the expansion of the SN shock front and vary over time. For uniform CSM distribution, the emission and attenuation due to SN evolution is often simpliﬁed as, !7/4 M˙ L6 GHz, peak ∝ (6.1) vw !2 M˙ τ6 GHz, peak ∝ (6.2) vw ˙ where M is the pre-SN mass-loss rate and vw is the pre-SN wind velocity (Weiler et al., 1986). The radio emission is the product of the magnetic energy and electron density 159

in the expanding shell and evolves with the size of the shell (R) and time (t) roughly as 3 −7/2 −1/2 Lν ∝ R t ∝ t . Similarly, the free-free optical depth for a shell expanding into the CSM evolves roughly as τ ∝ t−3. Thus, the light curve can be described by the product of these two processes: decaying radio emission with CSM that is becoming transparent. Figure 6.1 demonstrates a prediction for the ﬂux density at 6 GHz toward an example SN at 30 Mpc over a range of mass-loss rates. The mass-loss rates represent a range that has been detected at early times for SNe IIn and II-L in optical studies.

6.2.2 Measuring the Physical Properties of Molecular Clouds Shocked by Supernova Remnants

To date, the limiting factor in systematically analyzing SNR-MC interactions has been the small sample size – only about 30 of the 300 known galactic SNRs have strong evidence of a molecular shock (Jiang et al., 2010). Given the known association between core-collapse SNe, massive stars, and giant molecular clouds, this fraction sharply contrasts with the incidence of core-collapse SNe - roughly 3/4 of all SNe in volume-limited surveys (Li et al., 2011). Is the incidence of SNR-MC interactions truly low compared to the fraction of SNRs that result from core-collapse SNe or are our statistics incomplete? In Chapter 5, 12CO line emission spectra were examined for broadened molecular lines as a signpost of molecular shocks from SNR-MC interactions. This survey was the first of its kind and covered a large, volume-limited sample. I demonstrated that over 90% of SNRs observed with known, broad molecular features (discovered using a variety of telescopes, molecular lines, and analytical techniques) are also found using the SMT, in 12CO J = 2 − 1, and using our detection methods. Extending this analysis to all of the SNRs in our sample, I found nine new systems with broad molecular features. Given the incidence of detected interactions (19/50), I concluded that the fraction of SNRs with SNR-MC interactions is intrinsically low compared to the fraction of SNRs that result from core-collapse SNe. This result favors mechanisms other than SN feedback as the basic trigger for star formation. I also found that there is no strong correlation between sites of TeV gamma-ray emission and SNRs with MC interactions, contrary to predictions that SNR-MC interfaces are one of the primary venues for cosmic ray acceleration. This result implies alternate sources 160 for TeV gamma-ray emission observed toward SNRs such as pulsars and particle acceleration toward SNR interfaces with the warm, neutral medium. However, unambiguous examples of TeV gamma-rays have been found toward some SNR-MC interactions (e.g., W28; Aharonian et al., 2008a). Comparison between the physical conditions toward SNR- MC interfaces with and without TeV emission will be a vital next step in understanding the processes that lead high-energy particle acceleration and gamma-rays. Shock heating in MCs associated with SNRs is known to raise the kinetic temperature significantly higher than the ∼ 10 K typical of quiescent MCs. For example, molecular line surveys toward IC 443, a well-studied SNR associated with molecular shocks, suggest that temperatures exceed 50 K in some regions while densities range from 104 cm−3 − 105 cm−3 (Ziurys et al., 1989; Turner et al., 1992; Snell et al., 2005). These values imply 12 A p m 4 −3 that CO J = 3 − 2 (Eupper = 33.19 K, ncrit ∼ σ kT = 4.6 × 10 cm at 50 K) may be ideally suited for detecting shocked regions of molecular gas toward similar SNRs. My own data from the SuperCam 345 GHz receiver suggest that 12CO J = 3 − 2 is also a better tracer of regions associated with line broadening than J = 2 − 1 as observed toward the known SNR-MC interaction associated with 3C 391 (Figure 6.2). Given the utility of line ratios in characterizing SNR-MC interaction regions - especially those with signatures of line broadening - I am currently expanding upon previous observations of 12CO J = 2−1 toward SNRs by obtaining 12CO J = 3−2 maps of SNRs that may be sites of SNR-MC interactions. While surveys of molecular line strengths around individual SNRs exist, my observations have been directed toward a large-scale, volume-limited survey of SNRs. In some cases, the emission from SNR-MC interactions can occur in clouds beyond the radius of the SN shock front where “ejecta pistons” whose velocity exceeds the remnant free-expansion velocity (DeLaney et al., 2010) interact directly with a MC (Chapter 4). My observations cover sufficiently large areas around the remnants to sample molecular gas outside the immediate vicinity of the shock and a larger variety of possible molecular shocks toward Galactic SNRs. 161

Figure 6.2: (Top) 20 cm radio continuum emission of 3C 391 with contours (3σ, 4σ, 5σ) of integrated intensity of 12CO J = 2 − 1 from +102 to +108 km s−1 overlaid. (Bottom) Same as on the left but with contours of 12CO J = 3−2. The highlighted molecular cloud is clearly identiﬁed 12CO J = 3 − 2 emission, indicating that it has been shocked and heated. 162

APPENDIX A

Mid-Infrared Fine Structure Lines

We discuss the velocity structure and other aspects of mid-infrared ﬁne structure lines

other than [O IV] in this appendix. These lines have been observed previously in Cas A (Ennis et al., 2006; Rho et al., 2008). In this appendix, we present the deepest mid-infrared spectroscopic analysis performed on this remnant.

A.1 [S III] 18.71 µm line

The second brightest ﬁne-structure line after [O IV] is the [S III] feature at 18.71 µm. In

scan 1, there are three unique spectral lines around the [S III] rest wavelength centered at 18.65, 18.77, 18.81 µm (Figure A.1). These wavelengths correspond to radial velocities of −1600, −960, and 960 km s−1. These values do not appear to correspond closely with

any of the velocities calculated for the [O IV] emission, however, the limits to our spectral

resolution and the poor detection of [S III] in these frames may account for the discrepancy.

In scan 2, [S III] was more strongly detected in each frame and across more point-

ings. The high-velocity components in [S III] correspond to three distinct spectral features centered around the rest wavelength. These features are centered at 18.78, 18.89, 19.03 µm corresponding to radial velocities of −5130, −2880, −1120 km s−1. This pointing yields the largest absolute radial velocities detected in any line, although the two smaller

velocities agree with measurements of the [O IV] line.

A.2 [S IV] 10.51 µm line

Fine structure emission was observed around 10.5 µm along both scans, indicating some

emission due to [S IV]. At scan 1, we observed a single Doppler component around [S IV] (Figure A.2) where the line is dominated by the brightest component at 10.54 µm, corresponding to radial velocity of 960 km s−1. 163

2σ 0.25 S2,P1

0.20

0.15

Flux Density (Jy) 0.10

σ 0.16 2 S2,P2 0.14 0.12 0.10

Flux Density (Jy) 0.08 0.12 S2,P3 0.11 2σ 0.10

0.09

Flux Density (Jy) 0.08 0.10

2σ S1,P1 0.08

0.06

Flux Density (Jy) 0.04

0.060 0.055 S1,P2 0.050 2σ 0.045 0.040

Flux Density (Jy) 0.035 0.030 −10000 −5000 0 5000 10000 Radial Velocity (km s−1)

Figure A.1: Identiﬁed [S III](18.71 µm) lines in scans 1 and 2 (S1 and S2) at the southernmost pointings in each shock (P1-2 in S1, P1-3 in S2). These lines have been plotted versus relative velocity to demonstrate the presence of high-velocity components in the spectra. A sample error bar is shown in the upper-left hand corner to compare with the strength of the identiﬁed features. Five dotted, vertical lines mark some recurrent high-velocity features (at −2300, −200, 1100, 1700, 3200 km s−1). 164

2σ 0.08 S2,P1 0.06 0.04 0.02 Flux Density (Jy)

0.05 σ 2 S2,P2 0.04 0.03 0.02 0.01 Flux Density (Jy) 0.00

0.015 S2,P3 2σ 0.010

0.005

Flux Density (Jy) 0.000 0.030 0.025 2σ S1,P1 0.020 0.015 0.010 0.005 Flux Density (Jy) 0.000 0.010 0.008 S1,P2 0.006 2σ 0.004 0.002

Flux Density (Jy) 0.000

−10000 −5000 0 5000 10000 Radial Velocity (km s−1)

Figure A.2: Identiﬁed [S IV](10.51µm) lines in scans 1 and 2 (S1 and S2) at the southernmost pointings in each shock (P1-2 in S1, P1-3 in S2). These lines have been plotted versus relative velocity to demonstrate the presence of high-velocity components in the spectra. A sample error bar is shown in the upper-left hand corner to compare with the strength of the identiﬁed features. Five dotted, vertical lines mark some recurrent high-velocity features (at −2300, −200, 1100, 1700, 3200 km s−1). 165

Similarly, in scan 2 there is very little high-velocity splitting, and stacking the spectra reveals only one main component due to [S IV] at 10.57µm corresponding to a radial velocity of 1700 km s−1. This result is anomalous considering that we detected a component of differing velocity at pointing 1 and the critical densities of [S III] and [S IV] are essentially

the same yet we have found only negative velocities for [S III]. We infer that there may be some shock-excitation of high-velocity features apart from the dominant components

observed in the [S IV] line.

A.3 [Ne II] 12.81 µm line

Figure A.3 shows that there are three distinct high-velocity features at both scans. In scan 1, we identified two blue-shifted features centered at 12.74 and 12.76 µm and a third feature at 12.80 µm. These wavelengths correspond to radial velocities of 320, 1260, and 1730 km s−1. However, this identification is anomalous in that the blue-shifted features have approximately the same flux as the low-velocity feature and do not correspond to any

feature observed in [S III] or [S IV].

There is some agreement with the velocities observed in scan 1 for [O IV] and [S III]; assuming an error of roughly 200 km s−1 per measurement (given the spectral resolution of −1 the IRS high-resolution modules), the [O IV] feature at 1740 km s and [Ne II] feature at −1 1730 km s likely correspond to the same high-velocity component, as well as the [O IV] −1 −1 feature at 340 km s and the [Ne II] feature at 320 km s . Furthermore, the [S III] feature −1 −1 at 960 km s and [Ne II] feature at 1260 km s may correspond to the same component. In scan 2, the high-velocity structure exhibits a similar pattern with a bright peak at low-velocity and two distinct features that are both blue-shifted. These features occur at wavelengths 12.74 and 12.78 µm for the blue-shifted peaks and 12.82 µm for the low- velocity peak. These wavelengths correspond to radial velocities of 1730, 800, and −140 −1 −1 −1 km s . We can associate the [O IV] 1510 km s feature with [Ne II] at 1730 km s −1 within a reasonable margin of error in scan 2, as well as the [S IV] −290 km s and −1 [Ne II] −140 km s features.

It is puzzling that [Ne II] would exhibit exclusively blue-shifted features at high- velocity where the other lines indicate that there are red-shifted shock components. Fur- 166

0.04 S2,P1 2σ 0.03

0.02 Flux Density (Jy) 0.01

0.020 S2,P2 2σ 0.015

0.010 Flux Density (Jy) 0.005 0.014

S2,P3 0.012

0.010 2σ 0.008

0.006 Flux Density (Jy) 0.004

0.015 S1,P1

σ 0.010 2

0.005 Flux Density (Jy)

−10000 −5000 0 5000 10000 Radial Velocity (km s−1)

Figure A.3: Identiﬁed [Ne II] (12.81 µm) lines in scans 1 and 2 (S1 and S2) at the southernmost pointings in each shock (P1 in S1, P1-3 in S2). These lines have been plotted versus relative velocity to demonstrate the presence of high-velocity components in the spectra. A sample error bar is shown in the upper-left hand corner to compare with the strength of the identiﬁed features. Five dotted, vertical lines mark some recurrent high-velocity features (at −2300, −200, 1100, 1700, 3200 km s−1). 167 thermore, [Ne II] exhibited multiple blue-shifted components in scan 2 where [S IV] and

[S III] have none and [O IV] appears at only the highest velocity feature. These observed spectral properties must be explained by any model for the shocked gas.

A.4 Remaining emission lines

The remaining two emission lines [Ne V] at 14.32 µm (Figure A.4) and [Ne III] at 15.56 µm (Figure A.5) do not show any high-velocity structure above the noise level. Both lines have distinct peaks between 2000 and 3000 km s−1.

A.5 Cloudy Models of Purely Photoionized Gas

We used the spectral synthesis code Cloudy version 13.03 (Ferland et al., 1998) to approximate the contribution of radiative excitation to line intensities in the photonionization zone around Cas A. In models of photoionized ejecta in SNRs, the intensity of the ionizing ﬁeld is determined by Rmax, the lateral extent of the shock front compared to its thickness (Rho et al., 2009b; Raymond, 1979). In our models, we set the lateral extent of the shock to 2.4 pc, in agreement with ejecta models in Laming and Hwang(2003) and found the best

fit to our mid-infrared spectra with Rmax = 0.81. The ionizing radiation field was set by an approximation of X-ray spectra toward Cas A presented in Hwang and Laming(2012). X-ray spectra toward Cas A peak sharply around 1.9 keV, and most ejecta knots can be fit with an electron temperature around this value. Therefore, taking the total X-ray luminosity of Cas A to be 3 × 1036 ergs s−1 (Stewart et al., 1983), we approximated the ionizing radiation field in our Cloudy model as a blackbody with a temperature kBT = 1.9 keV. Finally, the abundances were fixed to those calculated for the entire remnant in (Hwang and Laming, 2012). Using these approximations, we determined the ejecta column densities that best fit the line intensities in our mid-infrared spectra (Figure A.6). The continuum in this spectrum is dominated by freshly formed dust composed of SiO2, Mg protosilicates, and FeO grains (Rho et al., 2008). In order to approximate the dust continuum in the region of our fine structure lines, we fit two power-laws independently to the emission shortward and 168

S2,P1 0.04

2σ

0.03 Flux Density (Jy) 0.02

0.01

S2,P2 0.025

2σ 0.020

0.015 Flux Density (Jy)

0.010

−10000 −5000 0 5000 10000 Radial Velocity (km s−1)

Figure A.4: Identiﬁed [Ne V] (14.32 µm) lines in scan 2 (S2) at the southernmost pointings (P1-2). These lines have been plotted versus relative velocity to demonstrate the presence of high-velocity components in the spectra. A sample error bar is shown in the upper-left hand corner to compare with the strength of the identiﬁed features. Five dotted, vertical lines mark some recurrent high-velocity features (at −2300, −200, 1100, 1700, 3200 km s−1). 169

S2,P1

0.04 2σ

0.03 Flux Density (Jy)

0.02

S2,P2 0.035

σ 0.030 2

0.025

0.020 Flux Density (Jy)

0.015

0.010 −10000 −5000 0 5000 10000 Radial Velocity (km s−1)

Figure A.5: Identified [Ne III] (15.56 µm) lines in scan 2 (S2) at the southernmost pointings (P1-2). These lines have been plotted versus relative velocity to demonstrate the presence of high-velocity components in the spectra. A sample error bar is shown in the upper- left hand corner to compare with the strength of the identified features. Five dotted, vertical lines mark some recurrent high-velocity features (at −2300, −200, 1100, 1700, 3200 km s−1). 170 longward of the 21 micron peak. We find that the column density of ejecta toward our spectra is approximately N = 1.1 × 1020 cm−2 toward the southernmost spectrum (P1) in S2. Given the IRS slit size and assuming a 3.4 kpc distance to Cas A, we approximate the total mass of ejecta in our spectra to be 0.05 M . 171 [S IV] [Ne II] [Mg V] [Ne V] [Ne III] [S III] [Ne V] [O IV]

Figure A.6: Synthesized spectrum (black, solid line) of gas illuminated by a blackbody to the approximate Cas A gas temperature in X-ray wavelengths ﬁt to our observed spectrum at S2, P1 (red, dashed line). The line intensities were determined using the CLOUDY pho- toionization code and the dust continuum emission was ﬁt using two power-laws shortward and longward of the 21 micron dust continuum peak. 172

APPENDIX B

Alternative Methods for CO Line Broadening

B.1 Broadening by Radiative Acceleration

We can test whether the gas was accelerated by the burst of radiation that was emitted from Cas A when the Type IIb event occurred. From infrared light echoes, Dwek and Arendt 11 (2008) determined that this burst had a luminosity of approximately Lr = 1.5 × 10 L with the maximum intensity at a wavelength around 100 nm, and lasted for approximately tr = 40 days. Using this information, an assumed distance from the supernova of d = 3 pc, and an estimate of the cross-sectional cloud area, we can calculate the amount of radiative momentum that could be transferred to a molecular cloud surrounding Cas A and thus estimate the line width of CO produced by the affected gas. For the size of each molecular cloud, we estimate a surface area oriented toward the source of radiation by assuming that the molecular cloud is approximately a sphere with mass M, radius R, and column density σ¯ averaged across the region where we evaluated the cloud mass (Section 3.3). Thus, R = pM/(πσ¯). By this method, we obtain radii 1.6, 1.6 and 1.2 pc for the western, southwestern, and southeastern clouds, respectively. The amount of momentum transferred to each cloud by the incident radiation is approximately

Z ∞ 2 Lrtr 2 2h LrtrR pr = 2 πR dν ≈ 2 (B.1) 0 4πd hν c 2d c Here we assume the proﬁle of the spectral energy distribution is sharply peaked at 100 nm (Dwek and Arendt, 2008) such that the majority of the radiation from the Cas A event comes out at this wavelength. The factor is an efﬁciency with which momentum is transferred to the cloud. For the limiting case in this paper, we assume that = 1, although typical values for momentum transfer by scattering of UV photons off of interstellar dust grains suggest that ∼ 0.3 (Draine, 2003). The momenta we obtain by substituting the 173

−1 radius of each cloud into this formula are 47, 47, 27 M km s . These momenta are two orders of magnitude smaller than required to account for the widths we observe in our molecular clouds (Section 3.3). It is not possible for radiation from the supernova ﬂash to provide enough acceleration to explain the line widths in our clouds.

B.2 Broadening by Gravitational Acceleration

Acceleration in molecular clouds may be a sign of gravitational collapse. To evaluate the likelihood of this scenario, we have estimated the virial mass of each cloud as

5Rσ2 M = w . (B.2) vir 8ln(2)G Here we use 6 km s−1 as the velocity width and the radius is calculated from surface density estimates. We have assumed that the clouds are spherical and measured their radii along their shortest axis of integrated 12CO intensity. The virial masses are then

12000, 12000 and 9000M . These values are at least an order of magnitude larger than the estimated cloud masses in Section 3.3, thus the clouds are stable against collapse and gravitational acceleration has most likely had no role in producing the observed line widths. 174

APPENDIX C

Individual SNR Distance Determinations

Kepler (G4.5+6.8) is a young SNR whose proper motion is fast enough to be measured. From Chandra ACIS-S measurements of Kepler’s expansion and a shock velocity inferred from X-ray synchrotron emission, Vink(2008b) determined that Kepler is located at a distance of 4.0 kpc. G11.2−0.3 was observed in neutral hydrogen by Becker et al.(1985) who found that −1 weak H I absorption can be seen up to 55 km s and roughly consistent with a distance of 4.4 kpc. G12.8−0.0 has an association with stellar cluster Cl 1813178 whose distance is known to be 4.8 kpc (Messineo et al., 2011). G15.9+0.0 does not yet have a trustworthy distance estimate, although multiple Σ- D relationships have been calculated for this remnant. Caswell et al.(1982) calculated the distance to be 16.7 kpc (updated to 8.5 kpc for the correct galactocentric distance as in Reynolds et al., 2006). Pavlovic et al.(2014) use the Σ-D relationship and place the remnant slightly further at 10.4 kpc. We use 8.5 kpc in our analysis for consistency with previous studies.

G16.7+0.1 has an associated pulsar with high Galactic H I column density along its line-of-sight (Helfand et al., 2003), roughly consistent with a distance of 10 kpc.

It has been argued that the complex of H II regions near G18.1−0.1 implies an asso-

ciation with the SNR. Several studies have used H I absorption estimates and inferred a distance of ∼ 4.0 kpc to G18.1−0.1 (e.g., Downes et al., 1980; Kolpak et al., 2003; An-

derson and Bania, 2009; Paron et al., 2013). However, new H I absorption estimates to the remnant itself imply this assumption may be in error and the remnant is located beyond

the H II regions at 5.6 kpc (Leahy et al., 2014), which we use in our analysis.

G21.5−0.9 has been identiﬁed from H I absorption between systemic velocities +67 km s−1 and +69 km s−1 for a distance of 4.8 kpc (Tian and Leahy, 2008b). 175

Radio recombination lines at systemic velocities from +42 to +99 km s−1 provide kinematic constraints on the distance to G29.6+0.1 (Lockman et al., 1996). These values, along with the distance to the nearby pulsar AX J1845-0258 (Torii et al., 1998), led Gaensler et al.(1999) to estimate that G29.6 +0.1 is at a distance of 10 kpc.

Tian and Leahy(2008a) measured the distance to G27.4 +0.0 (4C-04.71) in H I absorption to be 8.5 kpc. Based on X-ray absorption measurements, Vasisht et al.(2000) determined that the distance to G29.6+0.1 is 5 − −15 kpc (i.e., 10 ± 5 kpc).

G29.7−0.3 (Kes 75) has several distance estimates derived from H I absorption, separate Σ − D relationships, and CO associations (e.g., Caswell et al., 1975; Milne, 1979; Becker and Helfand, 1984; Su et al., 2009). We adopt the distance 6.0 kpc, derived by

Leahy and Tian(2008) from H I absorption studies. Case and Bhattacharya(1998) derived a distance of 12.9 kpc to G31.5−0.6 based on Σ − D estimates, and this value has been used in previous studies of this remnant (e.g., Mavromatakis et al., 2001). We adopt this distance in our study. The distance to G32.4+0.1 was measured in X-ray absorption by Yamaguchi et al. (2004), who determined the remnant lies at d = 17 kpc. (Frail and Clifton, 1989; Green and Dewdney, 1992) both measured the distance to

G33.6+0.1 (Kes 79) kinematically from H I and OH absorption, respectively. Their measurements both agree with a distance 7.1 kpc. The distance to 3C 396 has been derived in several ways, sometimes with discrepant results. Caswell et al.(1975) use H I absorption to determine 3C 396 is > 7.7 kpc, while

X-ray H I column density and CO associations indicate 6.2 − 8 kpc (Olbert et al., 2003; Hewitt et al., 2009; Su et al., 2011). We adopt the latest kinematic estimate of 6.2 kpc as in Su et al.(2011).

Toward 3C 397, H I absorption indicates a distance of at least 7.5 kpc while kinematic estimates from CO give 10.3 kpc (Caswell et al., 1975; Jiang et al., 2010). We use this estimate in our analysis.

Enhanced H I absorption toward W49B suggests the distance is 12.5−−14 kpc (Lock- hart and Goss, 1978; Brogan and Troland, 2001), from which we adopt the lower estimate 176 of 12.5 kpc. (Leahy et al., 2008) determine the distance to G54.1+0.3 from morphological association with CO around 53 km s−1 to be 6.2 kpc.

Optical observations of G59.8+1.2 reveal line emission in Hα,Hβ, [N II], [S II], and

[O III](Boumis et al., 2005). These data suggest that the SNR has a large angular size 0 −1 (∼ 20 ) and slow shocks (< 70 km s ). G59.8+1.2 is also obscured by a H I column 22 −2 ◦ density of NH ∼ 2.5 − 2.8 × 10 cm . For a Galactic longitude of l = 59.8 and a galactocentric distance of the Sun of R0 = 8 kpc, the minimum galactocentric distance along a line-of-sight near the mid-plane is approximately 7 kpc. We can apply an estimate of the distribution of neutral hydrogen along the line-of-sight to G59.8+1.2 determined in −3 Kalberla and Dedes(2008) to infer a distance to this SNR. For n0 = 0.9 cm , Rn = 3.15 kpc and assuming the galactocentric distance is 7 kpc . r . 35 kpc, the average mid-plane volume density of neutral hydrogen is

−(r−R0)/Rn nH (r) = n0e (C.1)

The distance to which the neutral hydrogen density proﬁle integrates to the measured value of 2.6 × 1022 cm−2 toward G59.8+1.2 is d = 7.3 kpc, which we use as the distance to this SNR.

Associated H I and CO features toward G63.7+1.1 yield a distance of 3.8 kpc (Wallace et al., 1997).

The distance to G69.0+2.7 (CTB 80) has been determined from H I emission to be approximately 1.5 kpc (Leahy and Ranasinghe, 2012; Park et al., 2013).

H I absorption and a possible CO association toward G74.9+1.2 (CTB 87) yield a distance estimate of 6.1 kpc (Kothes et al., 2003). G76.9+1.0 has an associated pulsar whose coordinates and spectra best place the SNR in the Outer Arm at d ∼ 10 kpc (Arzoumanian et al., 2011). Leahy and Green(2012) determined the distance to G84.2 −0.8 from X-ray absorption to be 6 kpc. Proper motion measurements of Cas A along with spectroscopic measurements of the expansion velocity suggest the remnant is located at a distance of 3.4 kpc (Reed et al., 177

1995). Hailey and Craig(1994) determined the distance to CTB 1 kinematically from optical spectroscopy to be 3.1 kpc. The distance to Tycho has historically been controversial (e.g., Chevalier et al., 1980;

Schwarz et al., 1980). Measurements of H I absorption at the near and far side of Tycho suggest the remnant is located at d = 2.5–3.0 kpc (Tian and Leahy, 2011) while X-ray proper motion studies suggest 4 ± 1 kpc (Hayato et al., 2010). We adopt the near distance estimate 2.5 kpc in agreement with models of the shock wave (Smith et al., 1991; Ghavamian et al., 2001) and measurements of a reported companion star (Ruiz-Lapuente et al., 2004).

Several observations of G130.7+3.1 (3C 58) in H I absorption indicate a distance of 2 − 3 kpc (Green and Gull, 1982; Roberts et al., 1993; Wallace et al., 1994; Kothes, 2013), from which we adopt the latest estimate of 2 kpc. As we indicate in Section 5.5.1, HB 3 is located near W3(OH) in the W3 star-forming region. Hachisuka et al.(2006) measure the distance to W3(OH) kinematically from water masers and find it is located at d = 2 kpc. Proper motions of ejecta toward G184.6−5.8 (Crab/Crab Nebula) in optical and radio wavelengths yield a distance of 2 kpc (Fesen and Staker, 1993; Bietenholz et al., 2004; Rudie et al., 2008). Welsh et al.(2002) constrained the distance to IC 443 by examining stars whose spectra indicated Na I and Ca II absorption from the remnant. These stars constrained bracketed the distance at 1.5 kpc. The remaining three SNRs – G12.5+0.2, G23.6+0.3, and G30.7−2.0 – have no distance estimates and either have no known type or are classified as having composite morphology (Table 5.1). However, the targets have observed sizes and 1 GHz surface brightnesses, to which we can apply the Σ−D relation and obtain a distance approximation. We acknowledge that this relationship should not be applied to remnants that are not shell-like unless the shell structure is well-resolved or bright enough that all other emission can be ignored. Therefore, these distance values should not be interpreted as a measurement to each SNR but a starting point we calculate for the purposes of our analysis. Using the 178 relationship calculated in Pavlovic et al.(2014) and the angular size and 1 GHz flux in Green(2014), we find distances of 15 kpc to G12.5+0.2, 6.9 kpc to G23.6+0.3, and 8.7 kpc to G30.7−2.0. 179

APPENDIX D

Ancillary Data on Previously Detected BML Regions Toward SNRs

G16.7+0.1 is a composite remnant initially identified at 5 GHz in Helfand et al.(1989). The remnant has a steep spectral index between 1.4 GHz and 5 GHz of α ∼ −0.6. Mul- tiple OH 1720 MHz maser searches have revealed a source along the southern half of the remnant, consistent with one of the brightest regions of the SNR at 5 GHz (Green et al., 1997; Hewitt et al., 2008). This maser emission occurs at +20 km s−1, and a tentative detection of OH (1667 MHz) absorption has also been reported at +25 km s−1. Kes 75 was initially identified as a SNR (referred to as G29.7 − 0.2) between 408 and 5000 MHz by its spectral index (Shaver and Goss, 1970). The spectral index measured in that study demonstrated nonthermal emission with α = −0.60. Low-frequency observations of this SNR are dominated by the outer shell where the emission demonstrates has an even steeper spectral index, up to α = −0.8 (Kassim, 1992). While the overall morphology appears to be composite, the steep spectral indices, even at high frequencies, suggests that the shell-like structure may be the dominant source of nonthermal emission. An X- ray pulsar was identified toward Kes 75 by Gotthelf et al.(2000). Timing measurements suggest that the pulsar is relatively young at 723 yr. Kes 79 was first confirmed as a SNR at 408 and 5000 MHz by Caswell et al.(1975). Subsequent pulsar detection and X-ray timing observations were performed with XMM- Newton (Gotthelf et al., 2005; Halpern et al., 2007). These observations confirmed the existence of a 105 ms X-ray pulsar toward Kes 79; however, the characteristic age was ˙ determined to be τc = 2P/P > 8 Myr. This value is significantly larger than the age estimated in Sun et al.(2004), where the authors identified a bright inner ring of material at 0.5 − 3 keV as a wind-driven shell and, assuming this material is expanding outwards at 5000 − 10000 km s−1, inferred an age of ∼ 6000 yr. Thus, while Kes 79 is likely a shell-like SNR with an embedded pulsar, the timing measurements imply the pulsar was spun up at birth and the younger age is more accurate. 180

OH lines at 1666, 1667, and 1720 MHz have been seen in both emission and absorption toward Kes 79 (Green, 1989; Green et al., 1997; Koralesky et al., 1998; Stanimirovic´ et al., 2003). The OH 1720 MHz maser was observed at +90 km s−1 in a single-dish detection, although the source has yet to be localized. In absorption, OH has been observed at +12 and +55 km s−1 and in a broad feature from +95 to +115 km s−1. It is inferred that this last feature corresponds to a coincident MC, and the broadening in OH is due to a shock interaction with the SNR. 3C 396 was classified as a SNR by Shaver and Goss(1970). 20 cm emission maps reveal composite morphology with significant brightening along the western half of the remnant, and especially toward the southwest (Patnaik et al., 1990). The spectrum at this frequency appears nonthermal (α = −0.42), and measurements of the spectral index at 90 cm agree with this value (Kassim, 1992). Green et al.(1997) reported detection of OH 1720 MHz maser emission at more than one pointing in single-dish observations of the remnant, although no velocity information or localization was obtained from these observations. This detection was subsequently confirmed by Koralesky et al.(1998) in VLA observations, and the authors report the presence of a +70 km s−1 OH line.

3C 397 was resolved and distinguished from a nearby H II region to its west by Kundu et al.(1974). The remnant is bright and compact, with several enhancements in the nonthermal radio emission to its west. The spectral index is steep from 6 to 20 cm in this region (α = −0.66) relative to the rest of the shell (α = −0.59)(Anderson and Rudnick, 1993). Observations at 1720 MHz revealed no OH maser emission and imply that if there is a SNR-MC interaction the physical conditions may be too extreme to excite a maser (Green et al., 1997). W49B was initially identified as a strong radio source and probable SNR at 408 MHz by Shaver and Goss(1970). The remnant radio spectrum has a moderately nonthermal spectral index between 408 MHz and 15 GHz of α ∼ −0.48 (Green et al., 1975). Thus far, no OH maser emission has been detected toward this remnant in targeted searches (e.g., Green et al., 1997). G54.1+0.3 is a diffuse, low surface brightness radio source identified as a SNR in Reich et al.(1985). The source is significantly polarized ( 7%) at 4.75 GHz, but has a 181

ﬂat spectral index at this frequency of α ∼ −0.1. Subsequent observations in the X-ray with ASCA and at 1175 MHz with Arecibo revealed a 136 ms pulsar, which combined with the SNR morphology, suggests that G54.1+0.3 is a Crab-like remnant with an age of ∼ 2900 yr (Camilo et al., 2002). HB 3 is a large SNR with well-deﬁned shell structure at 2.7 GHz (Velusamy and Kundu, 1974). The remnant is embedded in a star-forming region that contains the nearby

H II region W3. Spectral indices from 38 MHz to 2.7 GHz (α ∼ −0.5) suggest HB 3 has a strong nonthermal component. There is also evidence for a SNR-MC interaction from several OH 1720 MHz masers that have been detected to the southeast of this SNR and near W3 (Koralesky et al., 1998). One of these detections, referred to as “W3 northwest” in Koralesky et al.(1998), occurs at −38.6 km s−1 and is directly to the southeast of the remnant. Toward HB 3, this systemic velocity corresponds to a distance of 3.7 kpc, slightly farther than inferred from this Σ − D relation, although HB 3 is extended (∼ 800) and presumably an old remnant at this distance (∼ 30, 000 yr in Lazendic and Slane(2006)). Thus, surface brightness and size estimates are less reliable than for younger remnants and may yield discrepant distances compared to other methods. 182

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