EXPERIMENTAL DEVELOPMENTS FOR THE STUDY OF EXPLOSIVE NUCLEOSYNTHESIS IN STARS

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EXPERIMENTAL DEVELOPMENTS FOR THE STUDY OF EXPLOSIVE

NUCLEOSYNTHESIS IN STARS

by

Luke Erikson

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A thesis submitted to the Faculty and the Board of Trustees of the Colorado School of Mines in partial fulfillment of the requirements for the degree of Doctor of Philosophy in Applied Physics.

Golden, Colorado Date

Signed:

Luke E. Erikson

Approved:

Uwe Greife Professor

Department of Physics

Golden, Colorado Date

Tom Furtak

Professor and Department Head Department of Physics

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ABSTRACT

For several years now, theν-SNS collaboration has been working to place a small neutrino detector at the Spallation Neutron Source at Oak Ridge National Lab. If successful, the experiment may produce the needed neutrino-nucleus cross sections on solid targets such as iron and aluminum. These reaction probabilites are of great interest for a number of reasons, including: neutrino astronomy, explosive nucleosynthesis, and nuclear structure.

However, success for this project requires a very efficient cosmic ray detector to exclude backgrounds. The system would need to be ∼99% efficient while remaining affordable in a difficult financial climate for basic science. The first half of this thesis addresses a prototype cosmic ray veto based on extruded scintillator with embedded wave-length-shifting fibers.

This approach has been successfully used before, and may provide the performance needed for this project. However, our results suggest some additional research and development would be required to meet the requirements for theν-SNS experiment.

The second half of this thesis relates to experimental work to study the resonance strength of the ²³Mg(p,γ)²⁴Al reaction. For this purpose a radioactive ion beam experiment has been conducted at TRIUMF using the DRAGON experiment. This reaction is thought to play an important role during explosive nucleosynthesis such as novae and X-ray bursts. If so, then accurate knowledge of this break-out reaction would help explain the isotopic abundances around that mass range in the universe.

Our results suggest the rate of this reaction at astrophysically relevant energies is lower than predicted and might further exclude explosive binary systems as the production site for such elements as²⁶Al.

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FOREWORD

This is it. That moment they told us about in high school when one day algebra would save our lives.

—Gallagher in Red Planet After spending 5 years of my life pursuing a thesis in nuclear astrophysics, I find it is easy to take myself much too seriously. Working in this field seems to give you license to tell your friends and family how you are doing such things as peering at the inner workings of the Sun or explaining how our universe came to be composed of exactly this mix of materials. To a large degree, I feel these grand and egotistical statements are a sort of mantra we repeat to prove that what we do is important to ourselves and others. Lets face it, we’re not saving lives.

The fact is, discovering such things as exactly how a star explodes are difficult to convert into something that contributes to the quality of life for the humans of today or tomorrow or maybe ever. Clearly one or more people at the U.S. Department of Energy (the primary financial sponsor of my work) are interested in explosive nucleosynthesis and how it occurs in the universe. For my part, I am humbled for someone to think I’m the person to find out such wonderful things. With all the other things happening in the world, to give me so much time and money to explore the subject either seems very farsighted or very foolish depending on the headlines I read.

I wish for this work to bring something good to this world. It can be small, such as satisfying the curiosity of how the minute details of these events occur. Hopefully it will help other scientists advance the research described in this work. Perhaps it will.

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This may sound odd, but I also wish for something from you the reader.

Please do not spend time contemplating ways to use this work to bring evil to this world. Besides, I’d like to think you’d be wasting your time. Thank you very much.

Finally, I wish to express my sincere appreciation to everyone who helped me accomplish this degree. Most especially I want to give special thanks to the following people: My wife, Rebecca, who’s love for me was a guid- ing light for the last 5 years. My mother, who’s love and many sacrifices brought to me everything I can imagine to succeed in life. My advisor, Uwe Greife, who is the exception to all those grim stories you hear about PhD advisors. And finally, Christof Vockenhuber, who was a friend, colleague and mentor at TRIUMF when I most needed it. Thank you for everything.

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CONTENTS

ABSTRACT . . . iii

FOREWORD . . . iv

LIST OF FIGURES . . . xi

LIST OF TABLES . . . xxiv

CHAPTER 1 SOURCES OF EXPLOSIVE NUCLEOSYNTHESIS . . 1

1.1 Explosive Nucleosynthesis in Binary Systems . . . 1

1.1.1 Novae . . . 4

1.1.2 X-Ray Bursts . . . 5

1.1.3 Type Ia Supernova . . . 6

1.2 Type II Supernovae . . . 7

CHAPTER 2 THEν-SNS COSMIC RAY VETO . . . 13

2.1 Scientific Goals . . . 14

2.1.1 Supernovae . . . 14

2.1.2 Nucleosynthesis . . . 16

2.1.3 Astronomy . . . 16

2.1.4 Other Examples . . . 20

2.2 Neutrino-nucleus Interactions . . . 20

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2.2.1 Charged Current Interactions . . . 21

2.2.2 Neutral Current Interactions . . . 23

2.3 The SNS Facility . . . 25

2.3.1 Neutrino Production at the SNS . . . 25

2.3.2 Neutrino Detection at the SNS . . . 29

2.4 Backgrounds at the SNS . . . 32

2.4.1 Cosmic Rays . . . 32

2.4.2 Neutron Backgrounds Caused by the SNS . . . 35

2.5 Design of the Veto . . . 41

2.6 Extruded Scintillator . . . 44

2.6.1 Scintillator Production . . . 44

2.6.2 Embedding Fiber in Scintillator . . . 48

2.6.3 Connecting Fibers . . . 55

2.6.4 Measuring Scintillator Quality . . . 60

2.6.5 Photon Calibration . . . 70

2.7 Cosmic Ray Telescope (CRT) . . . 74

2.7.1 Construction of the CRT . . . 74

2.7.2 Accidental Count Rate of the CRT . . . 76

2.7.3 Timing Resolution of the CRT . . . 77

2.8 Prototype Cosmic Ray Veto . . . 79

2.8.1 Construction of the Prototype . . . 79

2.8.2 Detection Efficiency of Extruded Planks . . . 81

2.8.3 Timing Resolution of Extruded Planks . . . 83 2.8.4 Detection Efficiency of a Multi-Layer Prototype Veto 85 2.8.5 Timing Resolution of the Multi-Layer Prototype Veto 88

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2.8.6 Effects of Shielding the Multi-Layer Prototype . . . . 90

2.9 Simulations . . . 93

2.9.1 Stage 1: Physical Simulation of Extruded Scintillator 95 2.9.2 Stage 2: Simulating the Prototype Veto . . . 99

2.10 Results and Discussion . . . 103

CHAPTER 3 ²³Mg(p,γ)²⁴Al WITH THE DRAGON EXPERIMENT . . 105

3.1 Scientific Goals . . . 105

3.2 The ISAC Facility at TRIUMF . . . 108

3.3 The DRAGON Experiment . . . 115

3.3.1 Gas Target . . . 117

3.3.2 Gamma Array . . . 120

3.3.3 Elastic Monitors . . . 121

3.3.4 Electromagnetic Mass Separator . . . 122

3.3.5 Ion Chamber . . . 125

3.3.6 Mass-Slit Radiation Monitors . . . 126

3.4 Upgrades to the DRAGON . . . 128

3.4.1 Ion Chamber . . . 128

3.4.2 Multi-Channel-Plate Local Time-Of-Flight System . . 130

3.4.3 DRAGON Attenuator . . . 133

3.4.4 Software Upgrades . . . 134

3.5 Acceptance of the DRAGON Separator . . . 136

3.6 Charge State Distributions . . . 137

3.7 Timeline of the²³Mg Beam Campaign . . . 140

3.8 ²³Mg Beam Production . . . 142

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3.9 Measuring²³Mg Beam Intensity . . . 147

3.9.1 Measuring the Beam Mixture . . . 147

3.9.2 Calibrating the Beta Monitor (S7) . . . 147

3.9.3 Calibration of the NaI Detectors . . . 153

3.9.4 Calibration of the HPGe Detector . . . 154

3.10 The Search for²⁴Al . . . 155

3.10.1 Recoil Identification with MCPs . . . 155

3.10.2 Particle Identification with the Ion Chamber . . . 156

3.10.3 Results of the Search . . . 157

3.11 Calculating the Expected Yield . . . 160

3.11.1 Definition of the Maximum Yield (ymax) . . . 160

3.11.2 Stopping Power and Energies Covered . . . 162

3.11.3 Calculating the Yield . . . 166

3.12 Calculating the Upper Limit . . . 168

3.12.1 Region 1 . . . 169

3.12.2 Region 2 . . . 169

3.13 Conclusions . . . 170

CHAPTER 4 SUMMARY . . . 171

4.1 Theν-SNS Cosmic Ray Veto . . . 171

4.2 ²³Mg(p,γ)²⁴Al with the DRAGON Experiment . . . 173

REFERENCES . . . 175

APPENDIX A RESULTS OF EXPERIMENTS WITH WLSF . . . 181

APPENDIX B EXPERIMENTS CONNECTING FIBERS . . . 187

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APPENDIX C COSMIC RAY TELESCOPE . . . 193 APPENDIX D MULTI-LAYER PROTOTYPE . . . 201 APPENDIX E OPTICAL SPECTRA OF RELEVANT MATERIALS USED

IN SIMULATIONS . . . 213 APPENDIX F CHARGE STATE DISTRIBUTIONS . . . 227 APPENDIX G TIMELINE OF THE²³Mg(p,γ)²⁴Al EXPERIMENT . . . 235 APPENDIX H E810 RUN SUMMARIES . . . 237

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LIST OF FIGURES

1.1 Hubble image of the binary system composed of Sirius A (the bright main-sequence star) and Sirius B (the white dwarf

in the lower left quadrant). . . 2

1.2 Artist’s rendition of a binary system showing the accumulation of matter from a main sequence partner. Here the degenerate partner is a black hole. . . 3

1.3 Optical view of supernova 1987A[11]. . . 10

1.4 Kamiokande II detecting SN1987’s neutrinos. . . 11

2.1 Example of the composition of the stellar core at one instant during collapse[24]. . . 22

2.2 The charged current neutrino-nucleus interaction; a = gen- eral case, b = low energy four fermion approximation(q² << M²_W)[48]. . . 23

2.3 The neutral current neutrino-nucleus interaction; a = gen- eral case, b = low energy four fermion approximation(q² << M²_Z)[48]. . . 24

2.4 Arial view of the SNS facility grounds. . . 26

2.5 Schematic of the SNS target hall. . . 27

2.6 The spallation process at the SNS. . . 28

2.7 Neutrino energy spectra for the SNS and supernova. . . 29

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2.8 Schematic of theν-SNS bunker. . . 30

2.9 Cosmic muon energy spectra at sea level. black triangles=µ⁺, red squares=µ⁻[52]. . . 33

2.10 Neutron sources from the SNS. . . 36

2.11 Neutron flux from the SNS[60]. . . 37

2.12 Time structure of the SNS. . . 38

2.13 Electronics configuration for n-γ experiments. . . 39

2.14 Example of n-gamma separation with BC-501A detector; blue gate=neutrons, red gate=gammas. . . 39

2.15 Neutron backgrounds detected at SNS; the peak shows neutrons consistently arriving after the proton pulse timing signal. A continuous neutron background is also present. . . . 40

2.16 Schematic of the cosmic ray veto surrounding the ν-SNS bunker. . . 43

2.17 The extrusion process at Itasca. . . 45

2.18 Various pictures taken during the Dec 2005 test extrusion. . 46

2.19 Cross section of the MINOS extrusion[65]. . . 47

2.20 Cross section of the MECO extrusion[60]. . . 48 2.21 Pictures showing the various stages of embedding fibers

in scintillator; a=cut and polished fibers, b=cut and painted fibers, c=fibers awaiting epoxy, d=epoxy filled groove with fiber, e=Al foil covering grooves, f=BC-620 painted sample. 51

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2.22 Setup for measuring fiber connections. . . 56 2.23 Electronics configuration for measuring transmission through

fiber connections. . . 57 2.24 Connection experiments on 19 June 2007. From left to right:

cyan=non-index matched OF with CSM hole connector, violet = non-index matched OF with W&M hole connector, blue

= index matched OF with CSM hole connector, red = index matched OF with W&M hole connector, black = no connection (baseline). . . 58 2.25 Connection experiments on 25 June 2007. From left to right:

blue = non-index matched OF with W&M hole connector (2 runs), red = index matched OF with W&M hole connector (2 runs), black = no connection (baseline, 4 runs). . . 59 2.26 End on views of 3 samples, each at a different stage of

preparation for light output measurements. . . 62 2.27 Direct light output setup at CSM. . . 63 2.28 Electronics setup for direct light output measurements. . . . 64 2.29 The relative light output for selected samples from the Dec

2005 test extrusion (the complete MECO and CSM production). . . 65 2.30 The relative light output for selected samples from the Dec

2005 test extrusion (CSM’s materials). . . 66 2.31 Sample extruded scintillator with wavelength-shifting fiber

embedded. . . 68 2.32 Schematic of the embedded fiber light output setup. . . 69

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2.33 The embedded fiber light output setup at CSM. . . 69 2.34 The electronics configuration for embedded fiber light output

setup. . . 71 2.35 Relative light output for a sample exposed to cosmic rays;

red = extruded, black = cast scintillator. . . 71 2.36 The electronics configuration for photoelectron calibration. . 73 2.37 An example single photon calibration, black = experimen-

tal data, red = model fit, dotted-blue = components of the photon calibration model. . . 73 2.38 Small cosmic ray detectors to be used for efficiency experi-

ments; the penny is shown for scale. . . 75 2.39 Electronics setup for measuring accidental coincidences. . 76 2.40 Measuring the accidental coincidence rate with the paddles

0 degrees of vertical. . . 77 2.41 The count rates of the CRT by angle (while properly aligned

in time). The black line is a fit to help guide the eye. . . 78 2.42 Electronics configuration for the measurement of CRT tim-

ing resolution. . . 78 2.43 Measurement of the time resolution of the CRT. . . 79 2.44 Plank numbers used in construction of MLP, viewed from

the end of the MLP furthest from fiber readout. . . 80 2.45 Experimental setup to measure the response of an extruded

plank to cosmic rays. . . 81

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2.46 Electronics configuration for measuring the detection efficiency for a single plank. . . 82 2.47 Detection efficiency along a plank. Black = CRD4 placed

above the center groove, red = left groove, blue = right groove as viewed from the fiber readout. The lines are linear fits to guide the eye. . . 83 2.48 Electronics setup to measure the timing resolution for a sin-

gle plank. . . 84 2.49 Time resolution of a single plank as measured on distance

from fiber readout. The line is a linear fit shown to guide the eye. . . 85 2.50 Experimental setup to sensitivity of multi-layer prototype cos-

mic rays. . . 85 2.51 Electronics configuration to measure detection efficiency of

multi-layer prototype. . . 86 2.52 Detection efficiency of MLP by offset from center (0 degrees

off vertical and requiring activation of at least 3 layers). The red line is the theoretical efficiency for a MLP with planks composed of 94.5% detection efficiency. . . 87 2.53 Detection efficiency of MLP by angle. Black = theoretical

efficiency of MLP composed of planks with 94.5% detection efficiency and requiring at least 2 layers activated, blue = experimental results requiring at least 2 layers, red = 3 layers, green = 4 layers. Lines are shown to help guide the eye. . . 88 2.54 Electronics setup for measuring time resolution for a single

layer (X) of the MLP. T0 indicates the first TDC channel and T1 is the second channel. . . 89

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2.55 Electronics configuration for detecting coincidences between layers in the MLP. . . 91 2.56 Effects of shielding with a 160milliCi²⁴⁴Cm¹³C source. Black

circles=doubles, Red triangles=triples. The Al shielding is at 215cm and steel shielding at 235cm. Lines are shown to help guide the eye. The blue and green lines indicate the expected response if there had been no shielding. . . 92 2.57 Effect of shielding with¹³⁷Cs and double coincidences. Black

= experimental data, red = background rate, blue = fit to datapoints excluding those with shielding. . . 93 2.58 Effects of shielding with¹³⁷Cs and triple coincidences. Blue

is measured count rates with the source, green indicates the background count rate. . . 94 2.59 Early results from simulating the extruded scintillator. . . 98 2.60 Simulating the extruded scintillator (including PMT); black =

experiment, red = simulation . . . 98 2.61 Simulation of energy deposited in a single plank by cosmic

muons. Black = 0 degree of vertical, red = 60 degrees off vertical. . . 100 2.62 Comparison of simulation (red) and experiment (black) for

energy signals from a single plank. . . 101 2.63 Comparison of simulation (red) and experiment (black) for

timing signals from a single plank. . . 101 3.1 Diagram of NeNa and MgAl cycles[68]. . . 106 3.2 Diagram of TRIUMF beam lines and experimental facilities. 109

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3.3 Diagram of ISAC beam lines and experimental facilities. Ar-

eas in grey are located below the main floor level. . . 111

3.4 Diagram of TRILIS in relation to the ISAC target-ion source. 112 3.5 Diagram of OLIS. . . 113

3.6 Picture of surface ionizing (left) and microwave (right) ion sources. . . 114

3.7 Diagram of the DRAGON from the TRIUMF website. . . 116

3.8 Diagram of a the gas target. . . 118

3.9 Diagram of gas recycling system [19]. . . 119

3.10 A diagram of BGO array around gas target [19]. . . 120

3.11 Diagram of the DRAGON electromagnetic separator elements. . . 123

3.12 Diagram of a DRAGON electric dipole [19]. . . 124

3.13 A diagram of a DRAGON ion chamber[18]. . . 125

3.14 A diagram of the beta monitoring detectors at the mass slit box. . . 127

3.15 Idealized behavior for an ion chamber with 2 particles with the same energy. The higher Z particle (red) initially de- posits more energy than the lower Z particle (black), later the roles are reversed. . . 129

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3.16 Particle identification using the 5 anode ion chamber during the Fall of 2007. This was done with a ²³Na(p,γ)²⁴Mg beam on the 647keV resonance. The horizontal axis presents the total energy measured across the full 25cm detector while only the second anode of length 5cm (starting from anode 0) is plotted on the vertical axis. Black=attenuated beam, blue=recoil singles, red=recoil coincidences. . . 130

3.17 Rendering of the DRAGON end-station elements. . . 131

3.18 Time calibration of the MCP-TOF. Each peak is a separate run with a different delay. The delays (from left to right) used were +4,+2,-2,-4,-8,-16ns. The axis is reversed in time because a delayed signal from the upstream detector stops the TOF measurement. . . 132

3.19 Diagram of apertures in the gas target [19]. . . 136

3.20 Example of rossum being used to measure a charge state fraction. The graphs show current readings from Faraday cups (in amps) with the X axis showing time (in seconds).

The order of readings, from left to right, are: black=FC4, red=FC1, black=FC4, red=FCCH, black=FC4. . . 139

3.21 Effects of pressure on the charge state distribution of²⁷Al at Elab = 720keV/u. Each line connects the datapoints for a single charge state: red=6+, green=7+, blue=8+, magenta=9+, cyan=10+. As the pressure increases, the particles reach equilibrium and the charge state fractions become more stable. . . 141

3.22 Shift summary of the ²³Mg beam campaign. Grey background shows time periods when beam was unavailable or too low in intensity. The red line indicates the minimum intensity required. . . 143

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3.23 Separation in ion chamber (3rd anode) with laser ionization (TRILIS). . . 148

3.24 Separation in ion chamber (3rd anode) without laser ionization. . . 148

3.25 Identification of 511keV coincidences between 2 NaI detectors. . . 153

3.26 Comparison of NaI 511keV coincidences to S7 scalers (the red line is a linear fit). . . 154

3.27 Recoil identification using the MCP TOF system with a 519 keV/u (lab)²³Na beam. Black counts are recoil singles and red counts are recoil coincidences (simultaneous detection of a gamma and a heavy ion). . . 156

3.28 Particle identification with a thick window and a 519keV/u (lab) ²³Na beam. Grey = attenuated ²³Na beam, black = recoil singles²⁴Mg, violet = recoil coincidences. . . 157

3.29 Particle identification with a thin window and a 497keV/u (lab) ²³Mg beam (runs #19236-19239). Grey = attenuated beam, black = recoil singles. . . 158

3.30 Particle identification using both MCP and IC with a 519keV/u (lab) ²³Na beam. Grey = attenuated beam, black = recoil singles, magenta = recoil coincidences. . . 159

3.31 Particle identification using both MCP and IC with a 497keV/u (lab) ²³Mg beam (run #19221). Grey = attenuated beam, black = recoil singles. . . 160

3.32 Measuring the stopping power of²⁴Mg . . . 164

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3.33 Energy ranges covered by E810 (in lab system). Light Green

= 8.0 torr in target, Dark Green = 6.5 torr in target, Blue = Range of Predicted Resonances. . . 166 3.34 The²⁴Al level scheme suggested by a recent Gammasphere

experiment[23]. . . 167 A.1 A comparison between fibers painted with BC-620 and cov-

ered with Al vapor. . . 181 A.2 A comparison between 3 fibers (1.5mm dia) and 2 fibers

(2mm dia) with 25cm leads. . . 182 A.3 A comparison between 3 fibers (1.5mm dia) and 2 fibers

(2mm dia) with 5m leads. . . 183 A.4 A comparison between 25cm and 5m leads (3 fibers with

1.5mm dia). . . 184 A.5 A comparison between 25cm and 5m leads (2 fibers with

2mm dia). . . 185 C.1 Plexiglass light guides (left) and cast scintillator used in cre-

ating the CRDs. . . 193 C.2 A plexiglass/scintillator paddle joined by BC-600 epoxy. . . 194 C.3 A paddle wrapped in aluminum foil and electrical tape. . . . 195 C.4 Creating a PMT holder for the CRD caddy. . . 196 C.5 Pieces of the caddy ready for assembly. . . 197 C.6 A top and bottom pair for a single caddy. The black pieces

hold the PMT. . . 198

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C.7 Adding a black clamping block to secure the paddle. . . 198 C.8 Joining the paddle to the PMT. . . 199 C.9 Two CRDs, one above and one below a scintillating plank

(enclosed in the light-tight black tube, functioning as a CRT. 200 D.1 A view of a plank awaiting a mixture of Epon 815c / Epicure

3234 epoxy in the grooves. The perspective is above the far end of the plank — opposite from readout. The end of the fibers glow because of light reflecting off the BC-620 painted tips. . . 201 D.2 Picture of aluminum foil tape being applied to an extruded

plank. . . 202 D.3 A silicon brush was used to smooth the foil and epoxy. . . . 203 D.4 View of planks being produced for the MLP. The black tube

contains the first plank while undergoing tests for detection efficiency. Stacks of completed planks sit to the right. Close to the right edge of the table shows 6 planks covered with dust shields as epoxy cures. . . 204 D.5 Two layers of the MLP. The chalky surface of the bottom

layer is caused by the BC-620 paint. . . 205 D.6 End view of the extruded planks before being wrapped in

black plastic. . . 206 D.7 The carriage used to hold the PMTs and connect the fibers

for readout. . . 207 D.8 The block used to secure the fibers from the MLP. . . 207

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D.9 Side view showing fibers from all 4 layers. Black plastic is being threaded between the layers to prevent cross talk. . . 208 D.10 The fibers from a single layer are shown being directed to

the PMT. Black plastic separates these fibers from the layers below. . . 209 D.11 Side view of the MLP showing the alignment of the CRDs

above and below the prototype. . . 210 D.12 The completed MLP shown from the opposite end of the

fiber readout. . . 210 D.13 The electronics rack used for these experiments. . . 211 D.14 CRD4 poised above the MLP. Also shown is the light tight

tube used for experiments upon the single plank as well as CRD1 and CRD3 configured in a CRT around the tube. . . 212 E.1 BCF-91A emission spectrum[13] . . . 213 E.2 mean free path of BCF-91A wavelength shifter dopant[14] . 214 E.3 reflectivity of BC-620 paint[14] . . . 215 E.4 optical reflectivity of undoped extruded polystyrene . . . 216 E.5 optical transmission of undoped extruded polystyrene . . . 217 E.6 emission spectrum of polystyrene[8] . . . 218 E.7 reflectivity of extruded scintillator capstock . . . 219 E.8 optical reflectivity of PMMA . . . 220

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E.9 optical transmission of PMMA . . . 221 E.10 emission spectrum of POPOP[35] . . . 222 E.11 optical extinction spectrum of POPOP[35] . . . 223 E.12 emission spectrum of PPO[35] . . . 224 E.13 optical extinction spectrum of PPO[35] . . . 225

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LIST OF TABLES

1.1 Summary of explosive nucleosynthesis in binary systems. . 2 2.1 Neutrino nucleosynthesis of rare elements [60]. . . 16 2.2 Table of leptons (without antiparticles). . . 22 2.3 Effects on the cosmic muon background. . . 34 2.4 A summary of fiber embedding experiments with an un-

known scintillator. . . 52 2.5 Summary of fiber connection experiments done on 25 June

2007. OF4 = non-index matched optical fiber, WMOF = index-matched optical fiber, WMHC = William and Mary hole- connector. . . 60 2.6 Summary of embedded fiber light output experiments. . . . 70 2.7 Timing resolution of the MLP (by layer) as measured with

CRT placed 2.3m away from fiber readout. The second column shows the time resolution measured by this setup while the third column takes the time resolution of the CRT into account via Equation 2.3. . . 89 2.8 Summary of material properties needed for stage 1 simula-

tions . . . 97 2.9 Parameters used to include the effects of the PMT in simu-

lation. . . 98

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2.10 Comparing the simulated and experimental sensitivity of the MLP to neutrons and gammas with ²⁴⁴Cm¹³C source. ‘exp.’

= experiment and ‘sim.’ = simulation. . . 102 3.1 A few acronyms used at TRIUMF. . . 108 3.2 Summary of masses/rest energies for this experiment[5]. . 137 3.3 Summary of beams where charge state distributions were

taken. ‘beam charge’ refers to the charge state of the beam as it was delivered by ISAC. The energy is in the lab frame.

The last column refers to the pressure in the gas target. . . 140 3.4 Effects on beam mixture by changing ISOL IMS4 Y-slits.

S7=beta monitor trigger rate, XslitM is the current on the mass slit. . . 146 3.5 Summary of detection calibration measurements on the beta

monitor. The last column indicates the charge state of the beam detected on the charge slits. . . 152 3.6 Measuring the beam energy (using charge state 8+) down-

stream from the gas target for a ²⁴Mg beam with E=512.25 keV/u (lab). . . 162 3.7 Measuring the beam energy after the gas target for²⁴Mg. . 163 3.8 Measuring the beam energy after the gas target for²³Mg. . 164 3.9 Summary of previous attempts to measure or predict the

astrophysically relevant resonances for ²³Mg(p,γ)²⁴Al. The last column shows the energy for the relevant reaction studied. Note that the Zegers[26] experiment was an inverse kinematic experiment. . . 165 B.1 Summary of fiber connection experiments on 19 June 2007. 187

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B.2 Summary of fiber connection experiments on 20 June 2007. 188 B.3 Summary of fiber connection experiments on 21 June 2007. 188 B.4 Summary of fiber connection experiments on 22 June 2007. 189 B.5 Summary of fiber connection experiments on 25 June 2007. 189 B.6 Summary of fiber connection experiments on 26 June 2007. 190 B.7 Summary of fiber connection experiments on 27 June 2007. 190 B.8 Summary of fiber connection experiments on 9 July 2007. . 191 B.9 Summary of fiber connection experiments on 10 July 2007. 191 B.10 Summary of fiber connection experiments on 11 July 2007. 192 B.11 Summary of fiber connection experiments 12 July 2007. . . 192 F.1 Charge state fractions measured during E810. . . 227 F.2 Charge state fractions measured during E810 (continued). . 228 F.3 Charge state fractions measured during E810 (continued). . 229 F.4 Charge state fractions measured during E810 (continued). . 230 F.5 Charge state fractions measured during E810 (continued). . 231 F.6 Charge state fractions measured during E810 (continued). . 232 F.7 Charge state fractions measured during E810 (continued). . 233

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F.8 Charge state fractions measured during E810 (continued). . 234 G.1 Abbreviated timeline of events during the²³Mg(p,γ)²⁴Al ex-

periment. . . 235 G.2 Abbreviated timeline of events during the²³Mg(p,γ)²⁴Al ex-

periment (continued). . . 236 H.1 Columns presented in Tables H.2 through H.6. . . 237 H.2 Expected yields by runs with ˜6.5T in the target. . . 238 H.3 Expected yields by runs with ˜6.5T in the target (continued). 239 H.4 Expected yields by runs with ˜8.0T in the target . . . 240 H.5 Measured detection efficiency of the NaI detectors by run. . 241 H.6 Measured detection efficiency of the NaI detectors by run

(continued). . . 242

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CHAPTER 1

SOURCES OF EXPLOSIVE NUCLEOSYNTHESIS

The formation of a nucleus from individual nucleons and/or other nuclei is referred to as nucleosynthesis. This is a large field of study and, depending on when and where the nucleus is created, a wide variety of physics play a role. The formation of the elements is a key component to such topics as big bang theory and the evolution and explosion of stars. The primary goal is explanation of the elemental abundances we see around us. This chapter contains a short description of some of the circumstances where nucleosynthesis is known to take place in our universe. Special attention is spent on those events relating to the experimental work presented later in this thesis.

1.1 Explosive Nucleosynthesis in Binary Systems

Nealy all main-sequence stars end their lifetimes as a degenerate object such as a white dwarf or neutron star. These stars have consumed all the fuel they are capable burning and now have begun a long cooling phase.

No longer producing energy, they also no longer participate in nucleosynthesis and the matter contained within the star is gravitationally trapped, unable to assist in the creation of new stars or planets.

However, these degenerate stars occasionally are partners of a more slowly developing star such as Sirius as shown in Figure 1.1. In these binary systems, the degenerate partner can then serve as a catalyst for some of the biggest thermonuclear explosions in the universe by using material from its partner. During these explosions, some of the newly created nuclides are ejected with enough energy to enter the interstellar medium. In this manner, the burned remains of old stars have helped to produce the elemental abundances detected in our own solar system and throughout

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Figure 1.1: Hubble image of the binary system composed of Sirius A (the bright main-sequence star) and Sirius B (the white dwarf in the lower left quadrant).

the Universe.

Astronomy provides several examples of explosive nucleosynthesis in accreting binary systems. These observations show there are different types of events depending upon the types of stars involved and the rate matter is passed between them. Table 1.1 gives a brief summary of the most understood categories of astronomical phenomena seen with these systems and the circumstances where they occur.

In all cases, matter is transferred from one star to the degenerate part-

degenerate type accretion rate event

white dwarf low novae

white dwarf high type Ia supernova neutron star low X-Ray burst neutron star high X-Ray pulser

Table 1.1: Summary of explosive nucleosynthesis in binary systems.

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Figure 1.2: Artist’s rendition of a binary system showing the accumulation of matter from a main sequence partner. Here the degenerate partner is a black hole.

ner by gravitational attraction. As the matter transfer occurs, it forms large structures referred to as Roche Lobes (see Figure 1.2) which are defined by the gravitational equipotential surface between the two obects. Eventu- ally, the transferred matter collects into an envelope or disk on the surface of the degenerate object. As more matter arrives, the temperature and pressure increase, eventually leading to the ignition of thermonuclear reaction sequences.

These explosions are complex and dynamic systems. In every example, numerous inputs are needed for accurate modeling and simulation. Cur- rently lacking is knowledge of many of the basic properties for the important nuclei and reactions, including their masses, half-lives, resonance strengths, resonance energies and nuclear structure. Additionally, better understanding of the convection and mixing occuring during these explosions is increasingly recognized as fundamental to accurately modeling these events.

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1.1.1 Novae

The present model for producing a nova assumes material is transferred from a star (typically a main-sequence star) to the suface of a white dwarf companion. It is believed the arriving matter accumulates in a thin layer of degenerate material, and a nova is produced by thermonuclear ignition. This scenario requires the accreting matter to remain relatively cool and to arrive slowly on the surface (otherwise a supernova is expected as described in Section 1.1.3).

As the density approaches 1000g/cm³, pp-chains start the fusion process near the surface of the white dwarf’s native material. But once the reactions begin, the temperature increases radically and seed nuclei that have mixed into the envelope start energy generation according to the hot-CNO cycle. Until the reactions raise the temperature above the Fermi tempera- tureTFthe pressure and density remain fixed while the temperature soars.

The Fermi temperature (in K) can be calculated by Equation 1.1 in which ρ is density in g/cm³andµerepresents the electron mean molar mass[57].

TF = 3.03 × 10⁵ ρ µe

!2/3

(1.1)

The delay in expansion caused by the degenerate state produces convection which quickly moves the hotter material to the surface of the envelope.

Because the hot-CNO process has been operating as part of a thermonuclear runaway, large amounts of high temperature radioactive nuclei are mixed throughout. The decay of these nuclei (mostly beta emitters) further accelerate energy production and the luminosity rises above the Edding- ton limit (10⁵L) and cause an expansion that propels the outer layers from the white dwarf[56].

Some observations show enhanced abundances for elements ranging from Ne to S in nova. Current speculation is that these observations originate from binaries where a white dwarf was created by a star that had advanced beyond the hot-CNO cycle¹. If correct, then the nucleosynthesis taking place during the thermonuclear runaway is sensitive to the type of white

1These stars are referred to as ONeMg white dwarfs. Stars that only advanced up to the CNO cycles are designated CO white dwarfs.

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dwarf involved. And indeed, recent models support this hypothosis[57].

For nova occuring on ONeMg stars, the reacting envelope will also contain the seeds for the NeNa cycle (such as ²⁰Ne) and the MgAl cycle (such as ²⁴Mg). Current models suggest significant amounts of ²²Na and ²⁶Al may be created by the reactions: ²⁰Ne(p,γ) ²¹Na(βν) ²¹Ne(p,γ) ²²Na(p,γ)

23Mg(p,γ) ²⁴Al(β⁺ν)²⁴Mg(p,γ) ²⁵Al(β⁺ν)²⁵Mg(p,γ) ²⁶Al. One of these reactions,²³Mg(p,γ)²⁴Al, will be discussed in greater detail in Chapter 3. Many of these reactions have large associated uncertainties which greatly affect the outcome of the current models. Based on the results of orbital satellites such as COMPTEL, the current models are over-estimating the abundances of²²Na and²⁶Al while underestimating the amount of material ejected during the explosion. Better knowledge of the key reactions listed above may help explain this discrepancy.

1.1.2 X-Ray Bursts

Otherwise very similar to novae events, X-Ray bursts occur when the degenerate partner is a neutron star. However, with the much higher mass of the neutron star, the accumulating envelope is more dense and easily sur- passes10⁶g/cm³on the surface. When ignition occurs under such extreme pressure, the explosion is powered by the energy production of pp-chains, the hot-CNO cycle as well as the triple-α-process. Additionally, the runaway quickly moves into the temperature regimes where the rp-process and (α,p) reactions can participate as well.

Present models for X-Ray bursts suggest a complicated explosion which can be clearly divided into phases as a different process becomes dominant. These phases are briefly discussed below, see [58] for more detail.

Stage 1: The high pressures and temperatures initiate the hot CNO cycles using seed nuclei existing on the surface of the neutron star. The origin of these nuclei can include ashes of previous bursts or may have ar- rived directly from the Roche Lobes. The extreme environment is not ideal for continued operation of the CNO cycles, however, and the seed nuclei are converted to more stable waiting point nuclei. However, the energy

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production raises the temperature to approximately2 × 10⁸K and promotes the explosion to stage 2.

Stage 2: With the higher temperatures, the triple-α process is activated.

Conversion of helium accelerates the runaway and the combination of triple-α and (α,p) reactions rapidly convert much of the matter into ⁵⁶Ni oceans which may span much of the star’s surface. The explosion peaks with temperatures of the order 10⁹K. But an equilibrum has now been reached where the population of ⁵⁶Ni is being maintained by a combination of (p,γ) and (γ,p) reactions. The respite allows the outer layers to expand as the degenerate condition has been lifted by the rapid temperature increases. With the consumption of the avaliable helium as well as the conversion of the waiting point nuclei, energy production drops and the temperature begins to decrease.

Stage 3: As the temperature decreases, the mass fraction of⁵⁶Ni moves away from equilibrum and the rp-process becomes dominant. Reactions convert the nuclei along the N=Z line and reach to the mass = 100 region of the table of the isotopes. However, there is substantial controversy today in regard how far the rp-process extends. While several suggestions have been made in the literature[70][50], it is clear more work is needed to examine where the rp-process would terminate during an X-Ray burst.

There is also significant debate regarding how much matter is ejected by these events[56]. Unlike novae, it is possible the ashes from the runaway are completely integrated into the neutron star and affect the seed nuclei distribution available for the next ignition.

1.1.3 Type Ia Supernova

Astronomers have observed other explosive phenomena which are at- tributed to explosive binary systems. The most important of which are type Ia supernovae²which are thought to be caused by complete destruc- tion of a white dwarf by a runaway thermonuclear explosion. Rather than

2These events are identified by the absence of hydrogen lines in their spectra.

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an explosion upon the surface of the star, these detonations are caused by energy produced by the dwarf during a core collapse.

For a white dwarf, accumulating enough matter to reach the Chandrasekhar limit requires the rapid accumulation of hydrogen. If the dwarf builds up a matter envelope too slowly then a nova event occurs and the star is un- likely to retain enough matter to cause a collapse. However, if the dwarf can accumulate matter quick enough to maintain a stable burn of the arriving hydrogen then the matter envelope doesn’t have a chance to accumulate for a nova explosion – instead the ashes of Carbon and Oxygen settle upon the surface[56].

Eventually, enough mass can be acquired in this way to collapse the highly degenerate core. Carbon within the core undergoes fusion and the temperature sets off a thermonuclear runaway that propagates outwards in a wavefront of burning nuclei. The degenerate state initially prevents expansion, but the delay causes convection and mixing which further accelerates the collapse. By the time degeneracy is lifted the star is of such high temperature it expands explosively and is destroyed.

Current models and simulations of type Ia supernova struggle with a lack of knowledge regarding the mass and composition required for a given star to cause a supernova. Again, there is great uncertainty regarding the effects of convection and the burning wavefront that occurs after collapse.

However, type Ia supernova are clearly an important part of the nucle- osythesis in the universe. These events are the primary source of Fe-peak elements including ⁵⁴Fe and ⁵⁸Ni. The neutronization of the core during collapse also produces many of the neutron-rich nuclei including⁴⁸Ca,⁵⁰Ti and⁵⁴Cr[56].

1.2 Type II Supernovae

Massive stars, late in their evolution, experience a violent core collapse followed by a gigantic outward explosion[12]. Referred to as a type II supernova, the explosion is epic³ and is one of the most violent phenomena in the universe. Remnants of the star are propelled outward with enough

3These explosions convert10⁴⁶joules of gravitational energy.

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energy to distribute them throughout their galactic environment. Present evidence suggests supernovae may be responsible for approximately half of the abundance of nuclides heavier than iron existing around us today.

The spectacle and importance of this energetic process has attracted attention for millennia but only recently have the fundamental physical pro- cesses involved become understood. We now know the amazing visuals account for only a very small fraction of the energy produced[36] [55] [12]

[67].

The supernova event can be thought of in terms of the following stages:

Stage 1: The core of a massive star collapses. These stars have mostly burned through their stores of hydrogen and helium, advanced past carbon burning, and have started producing heavier elements ranging all the way up to iron. This produces a dense degenerate core composed of iron, nickel and a few other elements of similar Z[60]. As the density of the star increases, gravitational forces begin to compress the iron rich core beyond the capability of the repulsive forces to resist. Initially, nuclear reactions increase the mass of the nuclei within the core[60]. But as the degenerate core continues to collapse, the temperature rises dramatically and photodisintegration fractures the core nuclei into alpha particles and free nucleons. No longer composed of tightly packed nuclei, the core goes into freefall. The collapse accelerates, temperature continues to rise, photodisintegration fractures the alpha particles into protons and neutrons[12].

Eventually the core reaches densities similar to those seen within the core of a nucleus and the collapse halts and then rebounds outwards.

Stage 2: The outer layers produce a shockwave. The core collapse occurs so quickly that the outermost layers of the star may never detect the change before the outward explosion occurs. Closer to the core, the star’s mass chases the collapse. Current theories suggest the mass im- mediately surrounding the core collapses inward, creating the wave-front of a massive shockwave. Attracted by gravitational forces, the shock-wave accelerates towards the core only to reverse course due to the rebounding core material. However, the outward propagating shockwave stalls as it lacks the energy to throw off the outer layers of the star. Meanwhile, free

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protons within the remaining core material undergo inverse beta decay to reduce the forces; eventually causing the core to neutronize on a grand scale[67]. For each neutron produced in this manner, a neutrino is emitted, thereby producing a large neutrino luminosity. The remains of the core have been converted into a proto-neutron star.

Stage 3: The star explodes. The stalled shockwave is re-energized (in a manner not fully understood) by the swell of neutrinos produced as the core becomes neutronized. Simulations have shown that without the massive neutrino flux the rebounding mass becomes stalled and the supernova fails to explode. In reality, the newly reenergized shockwave propagates outward creating the massive explosion. The outer layers are flung outwards at speeds in excess of 10,000,000m/s (1/30 the speed of light)[67].

By scattering and reacting with the nuclei in the shockwave, the neutrino wind alters the isotopic composition of the material of the exploding star.

The interaction between the non-core stellar material and the massive flux of neutrinos is a focus of study today. The neutrino, so weakly interacting it took decades to detect, somehow facilitates the supernova’s explosion.

Clearly the small neutrino-nucleus cross section suggests an extremely large number of neutrinos are required for this theory to be correct. And, in fact, today it is believed 99% of the energy converted during a supernova is carried away by neutrinos (approximately10⁵⁷ ν/s equivalent to 10⁴⁶ joules of energy).

Experiments have provided some evidence to support this theory. In 1987, a type II supernova in the Large Magellanic Cloud was detected on earth (this event was designated SN1987A; see Figure 1.3). In a 13 second window, the Kamiokande II experiment detected 11 neutrinos[40] – a count rate about 4 orders of magnitude above the normal count rate⁴(see Figure 1.4). Comparing the visual, X-Ray and neutrino luminosity of SN1987A suggests the number of neutrinos produced roughly agrees with theory.

Understanding the mechanics of stellar core collapse, and by extension the origin of much of the heavier nuclei in the universe, requires an accurate measurement of several neutrino-nucleus cross sections involved in driving the shock-wave outward from the star’s core. The ν-SNS ex-

4On average, the Kamiokande II experiment detected ∼14.5 neutrinos per day[41].

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Figure 1.3: Optical view of supernova 1987A[11].

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Enriching the universe with

stellar ashes (Supernova 1987A)

Hubble Telescope

Kamiokande

Figure 1.4: Kamiokande II detecting SN1987’s neutrinos.

periment, as described in the next chapter, is intended to measure these reaction rates on a few key materials.

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CHAPTER 2

THEν-SNS COSMIC RAY VETO

The neutrino is a difficult particle to study. As it rarely interacts with matter, the neutrino forces us to go to great lengths to detect its presence. How- ever, it is this same property that makes this particle so very attractive to astrophysicists. Neutrino detection offers one of the holy grails of nuclear astrophysics — the ability to observe the interior of a star. Apart from os- cillation between flavors, a neutrino produced by stellar fusion is likely to escape the star entirely without interaction.

However, neutrinos are reluctant to give up their secrets. In terms of nuclear astrophysics, knowledge of neutrino-nucleus cross sections are the most valuable missing ingredients. Computer models suggest the outcome of stellar core collapse is sensitive to the reaction rate of a number of neutrino-nucleus interactions, leading us to believe these cross sections are key to detailed understanding of supernova physics. Since the neutrino so rarely interacts with matter, measuring the neutrino-nucleus cross sections is a difficult, costly, and time consuming task. Fortunately, by knowing even a few of the most influential cross sections we could make tangible progress towards the accurate modeling of core-collapse supernovae.

With this motive, theν-SNS collaboration was formed. This group intends to place neutrino detectors at the Spallation Neutron Source (SNS) located at Oak Ridge National Laboratory (ORNL) to measure neutrino-nucleus cross sections. The SNS offers some unique advantages for such studies:

• The SNS produces the most intense flux of pulsed neutrinos on the planet.

• The energy spectra of these neutrinos overlap with those produced during supernova events.

• Multiple flavors of neutrinos are produced which allows several neutrino- nucleus interactions to be studied.

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• The pulsed nature of the SNS is helpful to distinguish neutrino interactions from background.

Because of these features, a relatively small detector (with mass as small as 10 tons) could produce a statistically significant measurement within 1 year¹. In order for cross sections to be accurately measured, the active and passive shielding surrounding a neutrino detector at the SNS must be extremely efficient at detecting and eliminating background. While the pulsed nature of the proton beam reduces backgrounds by a factor of 17,000, cosmic ray backgrounds must be reduced a further 100,000-fold.

This suppression will be achieved by a combination of active and passive shielding, combined with particle identification within the detectors them- selves. The Colorado School of Mines (CSM) nuclear physics group’s role was to design and implement the active component of the bunker’s shield (see Section 2.5).

2.1 Scientific Goals

For decades, our difficulties in detecting the neutrino has cast a shadow on some of the most interesting scientific questions of the our age. As described below, the neutrino’s interaction with matter plays an important role in a variety of fields, ranging from the smallest scales (e.g. lepton violation) to some of the grandest of astrophysical events (core collapse supernova). Greater detail can be found in theν-SNS proposal[60].

2.1.1 Supernovae

The primary goal of the ν-SNS collaboration is to gain insight into the dy- namics of stellar core collapse. The neutrinos produced as the core is neutronized play a pivotal role in propelling the outer layers of the star away from the collapsed core at high velocity. Early in this process, the collapse of the star’s core leaves the surrounding material falling inward towards what has now become a proto-neutron star. If this was the extent

1A 3.5 x 3.5 x 3.5 m homogeneous detector placed 20m from the SNS target could produce 1000 counts on Carbon during a year of full power operation by the SNS[60].

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of the story, no explosion would follow as the outer layers would stall after rebounding off the degenerate core. In reality, as the mass falls inward and rebounds the neutrinos interact with the star’s envelope and cause the supernova’s explosion[60].

Recent advances have diminished some of the uncertainties in modeling a supernova. For example, experiments and advanced shell models have increased the accuracy of the rate of inverse beta decay on the nuclei of special interest for modeling supernova events[60]. These results have been used to improve the precision of the predicted timing and strength of the neutrino flux produced by the collapsing core. As the knowledge of the neutrino flux has increased, the uncertainty of the neutrino-nucleus reaction rates have become more dominant. Computer models of stellar core collapse now show a strong dependence on the cross sections of neutrinos with stellar material. Current models suggest a 10% change in these rates can greatly affect the outcome of the star’s collapse. Addi- tionally, accurate measurements of these cross sections can be combined with neutrino astronomy (see below) to validate supernovae models[60].

Despite growing confidence in these models, some concern remains regarding the identification of which nuclear species play important roles in the process of neutronizing the core. Even though the star’s core is initially composed of iron, as the gravitational forces increases the core nuclei temporarily increase in mass as the degenerate forces can no longer prevent nuclear reactions. For a short time during the collapse a wide variety of nuclear species exists within the core (see Figure 2.1). Eventually the pressure increases to the point where these nuclei are fractured into alpha particles and later into closely packed nucleons.

Both of these issues are at the forefront of nuclear astrophysics research today.

In what manner does the massive neutrino flux re-energize the stalled shockwave? What is the precise flux of neutrinos and which electron capture reactions are responsible? To answer these questions, the collaboration intends to measure theνecapture rates on several nuclei up to A=100.

The few cross sections measured will guide theory towards reliable calcu- lation of neutrino cross sections across the table of nuclides.

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138Ba(νe, e⁻)¹³⁸La

139La(νx, νxn)¹³⁸La

180H f (νe, e⁻)¹⁸⁰Ta

181Ta(νx, νxn)¹⁸⁰Ta

12C(νx, ν⁰_xp)¹¹B

12C(νx, νxn)¹¹C(e⁺νe)¹¹B

12C(νx, νxd)¹⁰B

12C(νx, νxpn)¹⁰B

Table 2.1: Neutrino nucleosynthesis of rare elements [60].

2.1.2 Nucleosynthesis

While knowing these interactions to greater accuracy will be useful for modeling the core collapse of a star, neutrino interactions are also relevant for the study of nucleosynthesis within supernovae. Nucleosynthesis occurs as the star’s envelope presses down on the compressed core, producing the shock wave. The massive numbers of neutrinos affect the number and type of nuclei produced by heating the layers of stellar material.

In addition to scattering, some neutrinos are captured to produce a new nucleus, and this also alters the chemical composition of the non-core material. These capture reactions may be responsible for some of the rarest nuclei seen in nature (see Table 2.1).

Neutrino driven winds may also play a role in changing the reaction paths of the rapid neutron capture process (r-process). Today, we suspect much of the abundance of elements heavier than iron are products of the r- process. While observations have yet to confirm a location where the r-process is active, all of the current candidates happen to occur in the presence of large neutrino fluxes.

2.1.3 Astronomy

Operating a neutrino detector at the SNS also would provide useful information for neutrino astronomy. Today’s experiments to measure neutrinos from terrestrial and astronomical sources are limited by the uncertainties on interaction rates. Even though theν-SNS experiment is not directly suit-

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able for neutrino astronomy, the planned measurements would greatly increase the accuracy with which facilities such as Super-Kamiokande could measure the luminosity of a detected supernova. In this manner, the ν- SNS experiment will help in both the computer modeling of supernovae as well as the neutrino astronomy used to detect these events. Several cross sections are fundamental to understanding the entire Supernova process but the ones most material are neutrino interactions with hydrogen, oxygen, carbon, iron and lead. The following sections describe the relevant issues and interactions with each of these target nuclei. These passages are taken directly from theν-SNS proposal[60].

Oxygen: In water Cherenkov detectors like the Sudbury Neu- trino Observatory (SNO) and Super-kamiokande, the charged- current reaction ¹⁶O(νe, e⁻)¹⁶F is the principal channel for electron neutrino interactions for thermal sources in the rangeT_ν ≥ 4 − 5MeV. Its rate exceeds that of neutrino-electron scattering by an order of magnitude for T_ν ≥ 7 − 9MeV[38]. Moreover, the electron angular distribution is strongly correlated with the electron neutrino energy[22], providing a way to measure the incident neutrino energy.

In addition, the appearance of back-angle electron emissions from this reaction in, for example, Super-Kamiokande would result from energetic electron neutrinos, more energetic than predicted by supernova models, providing further evidence for flavor oscillations and thereby information about the µ and τ neutrino spectra emanating from supernovae[22].µ and τ neutrinos in the stellar core couple to the core material only via neutral currents, whereas electron neutrinos and antineutrinos couple via both neutral and charged currents. As a result, the former decouple at higher density and, therefore, temperature, and have harder spectra. Utilizing reactions on ¹⁶O, Langanke, Vogel, and Kolbe[45] have suggested a novel way of also unambiguously identifying µ and τ neutrino signatures in Super-Kamiokande. The larger average energies for these neutrino flavors may be sufficient to excite giant resonances via the neutral-current reactions ¹⁶O(ν_µ,τ, ν_µ,τ)¹⁶O^∗. These resonances are above particle threshold and subsequently de-

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cay via the emission of protons, neutrons, and γ rays. The γ rays would provide the µ and τ neutrino signatures. The two decay channels are: ¹⁶O^∗(γn)¹⁵O and ¹⁶O^∗(γp)¹⁵N. How- ever, potential channels for observing the µ and τ neutrinos from supernovae must be re-examined in light of recent work (see, e.g. [73],[66],[9]), which indicates that nucleon-nucleon bremsstrahlung and the effects of nuclear recoil in neutrino- nucleon scattering significantly soften theµ and τ neutrino spectra, lessening their energy excess over electron neutrinos.

Thus accurate measurements of both charged- and neutral- current neutrino cross sections on Oxygen would serve as a foundation for interpreting neutrino data from the next core collapse supernova in our galaxy and for using the data to poten- tially observeµ and τ neutrino spectra as they are emitted from the proto-neutron star. An experiment to measure the cross section for:

16O(νe, e⁻)¹⁶F

is a high priority. Further useful experiments would focus on the cross sections for:

16O(νµ, ν⁰_µnγ)¹⁵O

16O(ν_µ, ν⁰_µpγ)¹⁵N

Iron and Lead: The use of iron and lead in supernova neutrino detectors like the proposed ADONIS detector would provide another way of detecting the µ and τ neutrino spectra in core collapse supernovae[28]. Iron has a sufficiently high threshold for neutron production via charged-current neutrino interactions that such production is negligible, whereas in lead neutrons are produced by both charged- and neutral-current interactions. Oscillations between the more energetic µ and τ neutrinos and the electron neutrinos would boost the charged- current event rate while leaving the neutral-current rate roughly unchanged. Thus, the ratio of the event rate in lead to that in iron would serve as a further constraint on the extent of neutrino oscillations and the emittedµ/τ neutrino spectra, provided the

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reaction rates in iron and lead are well known. To further the development of a detector like ADONIS, experiments to measure the neutrino-iron and neutrino-lead cross sections are critical.

For iron, the neutral-current reaction:

56Fe(νx, ν⁰xn)⁵⁵Fe dominates.

For lead, cross sections for the following neutral- and charged- current channels are desirable:

APb(νx, ν⁰_xn)^A−1Pb

APb(νx, ν⁰_x2n)^A−2Pb

APb(νe, e⁻n)^A−1Bi

Since they are a both necessary calibration for future supernova neutrino experiments and are important to nucleosynthesis simulations, iron and lead cross section measurements are a very high priority.

Deuterium: Neutrino experiments that use heavy water, like SNO, can detect supernova neutrinos via four channels:

e⁻(νx, ν⁰_x)e⁻ d(νx, νxn)p d(νx, e⁻p)p d(νe, e⁺n)n

Measurement of the reaction d(νx, e⁻p)p which has been suggested as a calibration for the reactionpp(e⁺, νe)d (part of the pp chain of reactions powering the Sun), would also provide a calibration for heavy water neutrino detectors. Monte Carlo studies suggest that for the source brightness predicted for the SNS, two years of data in approximately fifteen fiducial tons of D2O would yield a cross section measurement with an accuracy of a few percent[6], which in turn may enable a more accurate interpretation of SNO data. This measurement would also serve as

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an important test case for the effective field theory approach to neutrino-nucleus interactions (see [34] and references therein).

Carbon: Large neutrino detectors that use liquid scintillator, such as KamLAND and MiniBoone, are ideally suited to detect antineutrinos from a supernova collapse viaν¯e+ p → e⁺+ n. In addition, a significant number of neutral- and charged-current interactions could be detected [25] via:

12C(νe, e⁻)¹²N

12C(νe, e⁻)¹²N^∗

12C(νx, ν⁰_x)¹²C^∗

It is clear that every recorded neutrino interaction from the next galactic supernova will be extremely important. Detailed measurements of neutrino interactions in multiple channels will help untangle intensities and temperature of individual neutrino species.

Althoughν+Carbon interactions have been measured by LSND and KARMEN,ν-SNS measurements of double differential cross sections can provide powerful additional information to assist the interpretation of supernova neutrino signals.

2.1.4 Other Examples

Other examples of possible experiments and their motivations include nuclear structure, coherent scattering[71], tests of the standard model, and enhanced understanding of muon decay (in terms of energy spectra and lepton number violation). Clearly, neutrino interaction rates play a vital role in many physics topics today.

2.2 Neutrino-nucleus Interactions

Neutrino scattering is an active field of study, taking place at several ac- celerator facilities around the world. In most cases the neutrinos involved have very high energies and are capable of energy transfers of the order

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1GeV. At these energies, a large number of neutrino interactions are possible, ranging from electron scattering to deep scattering off nucleons. Such high energy interactions can produce several particles during a single scattering event². These energetic (and complicated) reactions are excluded by the neutrino energies produced by stellar phenomena. Since the neutrinos produced at the SNS have at most 52.6MeV of kinetic energy[60], these lower energy neutrinos undergo a much more restricted range of interactions.

The Standard Model, with the electroweak extensions introduced by Glas- how, Weinberg and Salam, provides a theoretical framework for neutrino- nucleus interactions. According to this theory, neutrinos interact with nuclei via the weak force by exchanging heavy bosons between the two particles.

Depending on the type of neutrino and the energy available, this exchange may result in the neutrino being captured or scattered[16][59][60].

2.2.1 Charged Current Interactions

When the neutrino is captured by a nucleon within the nucleus and a new charged lepton is produced; this is referred to as a charged current (CC) interaction. Since the neutrino has been absorbed and the nucleon has undergone a conversion, conservation requires a new charged lepton to be produced and emitted. Since the type of lepton produced is limited by energy conservation, the neutrino involved must have enough energy to produce the corresponding particle (see Table 2.2). In most cases³, electron neutrinos are the only flavor capable of CC interactions.

Based on the Glashow-Weinberg-Salam Standard Electroweak Model, the Lagrangian for this interaction is dominated by the coupling between the W bosons and the charged fermion current (J)[48]. At low energies where the momentum transfered (q) is much less than the mass of the boson (MW) the process can be reduced to an effective four fermion contact interaction.

See Figure 2.2 for Feynman diagrams showing these interactions.

2A single deep inelastic scattering between a neutrino and a nucleon can produce a baryon and numerous mesons.

3Neutrinos from stellar fusion typically have energies less than 15MeV[41]. Nuclear reactors produce neutrinos with energies less than 8MeV[44]. Neutrinos from supernova have less than ∼60MeV[29].

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Figure 2.1: Example of the composition of the stellar core at one instant during collapse[24].

particle mass particle mass

electron 0.511MeV electron neutrino <3eV muon 105.6MeV muon neutrino <0.19MeV tau 1,777.0MeV tau neutrino <18.2MeV Table 2.2: Table of leptons (without antiparticles).