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2017 Photometric and Spectroscopic Signatures of Superluminous Supernova Events The puzzling case of ASASSN-15lh

Welbanks Camarena, Luis Carlos

Welbanks Camarena, L. C. (2017). Photometric and Spectroscopic Signatures of Superluminous Supernova Events The puzzling case of ASASSN-15lh (Unpublished master's thesis). University of Calgary, Calgary, AB. doi:10.11575/PRISM/27339 http://hdl.handle.net/11023/3972 master thesis

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Photometric and Spectroscopic Signatures of Superluminous Supernova Events

The puzzling case of ASASSN-15lh

by

Luis Carlos Welbanks Camarena

A THESIS

SUBMITTED TO THE FACULTY OF GRADUATE STUDIES

IN PARTIAL FULFILLMENT OF THE REQUIREMENTS FOR THE

DEGREE OF MASTER OF SCIENCE

GRADUATE PROGRAM IN PHYSICS AND ASTRONOMY

CALGARY, ALBERTA

JULY, 2017

c Luis Carlos Welbanks Camarena 2017 Abstract

Superluminous supernovae are explosions in the sky that far exceed the luminosity of standard supernova events. Their discovery shattered our understanding of stellar evolution and death. Par- ticularly, the discovery of ASASSN-15lh a monstrous event that pushed some of the astrophysical models to the limit and discarded others.

In this thesis, I recount the photometric and spectroscopic signatures of superluminous super- novae, while discussing the limitations and advantages of the models brought forward to explain them. I show that a quark nova occurring in the wake of a supernova remnant of an Oxygen-type

Wolf-Rayet star can reproduce the light curve, photospheric radius and effective temperature evo- lution of ASASSN-15lh. This model was used to successfully simulate the spectrum of the event using SYNOW. Beyond being a successful explanation for ASASSN-15lh, the quark nova model is an appealing mechanism to power the most luminous events ever seen in the sky.

ii Acknowledgements

When I heard the learnd astronomer,

When the proofs, the figures, were ranged in columns before me,

When I was shown the charts and diagrams, to add, divide, and measure them,

When I sitting heard the astronomer where he lectured with much applause in the lecture-room,

How soon unaccountable I became tired and sick,

Till rising and gliding out I wanderd off by myself,

In the mystical moist night-air, and from time to time,

Lookd up in perfect silence at the stars.

-Walt Whitman

To my father for his selfless sacrifices, for his advice and his endearing inspiration to pursue my dreams. To my mother for giving me breath. To my grandparents for their care and support towards my career. To Stephanie for her commitment, love, laughter, care and feedback.

To my supervisor and mentor, Dr. Rachid Ouyed, for believing in me and pushing me to become the best version of myself. His patience, empathy and thoughtfulness made this work possible. To Dr. Nasser Moazzen-Ahmadi for teaching me that discipline and humour are the virtues of a scholar, and that kindness is the key to success. To Dr. Jo-Anne Brown for teaching me to knock on the doors of opportunity and being a considerate advisor. To the Quark-Nova group,

Amir, Nico, Zach and Matt, for being a trustworthy source of knowledge and encouragement.

iii Table of Contents

Abstract ...... ii Acknowledgements ...... iii Table of Contents ...... iv List of Tables ...... v List of Figures ...... vi List of Symbols ...... viii 1 Introduction ...... 1 1.1 Stellar Evolution: The Events Leading up to a Supernova ...... 2 1.2 A New Piece in the Puzzle: Superluminous Supernova Events ...... 6 1.3 Powering Mechanisms for Superluminous Supernovae ...... 9 1.3.1 Nuclear Decay ...... 10 1.3.2 Circumstellar Material Interaction ...... 11 1.3.3 Magnetars ...... 11 1.3.4 Tidal Disruption Events ...... 12 1.3.5 Advantages and Disadvantages ...... 12 2 The Quark Nova as a Mechanism Behind Superluminous Supernovae ...... 14 2.1 The Quark Nova Hypothesis ...... 14 2.1.1 Nuclear Deconfinement and Seeding ...... 14 2.1.2 Burning ...... 17 2.1.3 Energetics ...... 18 2.1.4 Astrophysical Scenarios and Their Observational Signatures ...... 19 3 ASASSN-15lh. The Most Luminous Supernova Event Ever Seen ...... 23 3.1 Detection History: Photometric and Spectroscopic Data ...... 23 3.2 Standard Powering Mechanisms ...... 26 4 The Quark Nova Model for ASASSN-15lh ...... 30 4.1 Quark Novae in Massive Star Binaries ...... 30 4.1.1 The Quark Nova Luminosity ...... 30 4.1.2 Spin-Down Interaction ...... 32 4.1.3 Black Hole Accretion Phase ...... 33 4.1.4 Light Curve ...... 33 4.1.5 Obstacles and Challenges ...... 35 4.2 Dual-shock Quark Nova With a Wolf-Rayet Star Supernova Progenitor ...... 35 4.2.1 Light curve ...... 40 4.2.2 Spectrum ...... 43 4.3 Discussion ...... 50 5 Summary and Conclusion ...... 52 Bibliography ...... 56

iv List of Tables

1.1 Advantages and disadvantages of models for superluminous supernovae ...... 13

4.1 Best fit parameters for the ASASSN-15lh light curve in the binary model ...... 34 4.2 Best fit parameters for the ASASSN-15lh light curve in the dual-shock model . . . 41 4.3 Elemental species parameter for synthetic spectrum of ASASSN-15lh using SYNOW. The envelope was assumed to be fully mixed giving all elements the same excita- tion temperature ...... 47

v List of Figures and Illustrations

1.1 Schematic structure of the burning stages of a star. This is the structure of a SN progenitor star...... 4 1.2 This Figure shows the light curves of SNe and SLSNe. While common SNe can reach peak luminosities of 2 × 1043 erg s−1(absolute magnitude > −19.5), SLSNe reach luminosities that are greater by a factor of 10. The Figure contains prototypi- cal events for each type of SN and compares them to SLSNe PTF09cnd, SN2006gy and SN2007bi. The specifics of the classification suggested by Gal-Yam are dis- cussed in the source. (Figure 1 from Gal-Yam [19].) ...... 8 1.3 This Figure shows the early spectra of some SLSNe-I. The absorption lines cor- respond to light elements like O, Mg, Si, and C. Other than the presence of these lines, a characteristic is the absence of H or He lines. Hα has an expected wave- length of 6562.8 A.˚ (Figure 4 from Gal-Yam [19].) ...... 9

2.1 A massive binary can evolve into a QN. A subsequent BH can accrete matter from the system rebrightening the light curve. The 1x and 10x in Panels g and G corre- spond to 1 and 10 times zoom respectively. (Figure 1 from Ouyed et al. [52].) . . . 22

3.1 Figure showing the rest-frame absolute magnitude light curve of ASASSN-15lh near peak compared with other type I SLSNe events. The comparison contains the most luminous events previously known showing that ASASSN-15lh is more luminous than the other events by roughly more than 1 magnitude. Dong et al. [16] contains a full description of the corrections performed to obtain this image. (Figure 4 from Dong et al. [16].) ...... 24 3.2 Multi-band light curve of the rebrightened ASASSN-15lh. The magnitudes are in the Vega system and were not corrected for extinction or contribution from the host. For more information on methods read Godoy-Rivera et al. [21]. (Figure 1 from Godoy-Rivera et al. [21].) ...... 25 3.3 Top Panel: Spectroscopic evolution of ASASSN-15lh at the rest frame. Differ- ent curves represent different stages of the light curve evolution. (Figure 2 from Godoy-Rivera et al. [21]). Bottom Panel: Spectroscopic evolution of ASASSN- 15lh near Hα. The spectra cover a period of time before and after the fading stage at around 196 days. The green line at 6562.8 A˚ shows the expected position of Hα. For both panels, offsets were added to the spectra for clarity. (Figure 3 from Godoy-Rivera et al. [21].) Godoy-Rivera et al. [21] explain in detail the dates of observation and instruments used for observation...... 27 3.4 Top: Effective photospheric temperature of ASASSN-15lh. Bottom: Apparent photospheric radii of ASASSN-15lh. Data obtained from Godoy-Rivera et al. [21]. 28

4.1 Light curve for ASASSN-15lh using the QN model in a binary. Obtained light curve for ASASSN-15lh. The purple stars represent the observations taken from Godoy-Rivera et al. [21]. The red line shows the result of the model while the contributions of each engine are shown separately. In blue the QN, in green the SpD contribution, and in black the BH accretion...... 34

vi 4.2 Schematic representation (not to scale) of the sequence of events in our model. The top left panel shows a WO star that later undergoes a SN in the top right panel. The middle left panel shows the two-component configuration of the WO-SNR composed of a dense core and a less dense envelope. The middle right panel shows the QN that results in the birth of a QN and a QN ejecta. The QN ejecta shocks the core and the envelope. The envelope later cools and recedes. The bottom left panel shows the photosphere receding through the envelope resulting in the first hump in the light curve. The bottom right panel shows the photosphere receding through the hot core which rebrightened the light curve resulting in a second hump. . . . . 38 4.3 Results for the features of ASASSN-15lh using the dsQN model are shown with solid red lines. The blue line (black line) shows the envelope (core) contribution to the light curve. The data is represented using coloured stars. Top: The dsQN model fit to the light curve. Middle: The results for the photospheric radius evolution. Bottom: Effective temperature evolution. Data obtained from Godoy-Rivera et al. [21]...... 42 4.4 Synthetic spectrum of ASSASN-15lh using SYNOW at ∼ 26 days post peak lumi- nosity assuming an envelope with OII only. The optical depth used was τ(OII)= 0.2. The inlet shows an amplified look at the double feature of OII. The absorption feature at ∼ 4400 A˚ is not present in the observed spectrum. Data obtained from Leloudas et al. [34]...... 46 4.5 Synthetic spectrum of ASSASN-15lh using SYNOW at ∼ 26 days post peak lu- minosity. The top panel shows the total simulation of the model using an O com- postion and an O and Mg compostion. The bottom panel shows the individual contributions of each elemental species assumed in the model. The data was ob- tained from Leloudas et al. [34]...... 48 4.6 Synthetic spectrum of ASSASN-15lh using SYNOW at ∼ 180 days post peak lu- minosity. The spectrum is practically featureless and it is modelled with a shell of O and Mg. The data was obtained from Leloudas et al. [34]...... 49

vii List of Symbols, Abbreviations and Nomenclature

Symbol Definition ASASSN All Sky Automated Survey for SuperNovae

BH Black hole CE Common envelope

CO Carbon and oxygen

CSM Circumstellar material dsQN Dual-shock quark nova

KE Kinetic Energy

NS Neutron star PISN Pair instability supernova

QGP Quark-gluon plasma

QN Quark nova

QS Quark star

SLSN Superluminous supernova

SMBH Supermassive black hole

SN Supernova

SNR Supernova remnant

SQM Strange quark matter

TDE Tidal disruption event

U of C University of Calgary

viii Chapter 1

Introduction

The traditional picture of stellar evolution, the one in textbooks and taught in lecture halls, tells us

that the life of a star ends once it runs out of fuel to burn. A massive enough star dies in an extreme

explosion known as a supernova (SN) and a neutron star (NS) is left behind. However, recent

observations point towards a not-so-straight-forward death. New findings show supernovae (SNe)

explosions that are much brighter than what we traditionally thought was possible, with hints of

what could be a second explosion or even rebrightening following the first one. Traditional models

do not account for superluminous supernovae (SLSNe). SLSNe are explosions in the sky releasing

> 10 times more energy than a typical SN, and some of them bang twice. This is the puzzle of

SLSNe.

First seen at the turn of the millennium, SLSNe puzzled the scientific community because they challenged our understanding of stellar evolution and death. Not only are their energetics extreme, but their spectrum and time evolution are significantly different from any other event we have seen to this day. Understanding what is the mechanism behind these extreme explosions and why some of them explode twice, will help us form a more complete picture of stellar evolution and explore new venues that can explain these exceptional events.

The main focus of my work is to explain the photometric and spectroscopic features of SLSNe as applied to the most energetic event to this day, ASASSN-15lh. My main contribution is to expand on an existing model for SLSNe and determine which parameters explain the photometry and spectroscopy of ASASSN-15lh. Here, I explore the explosive transition of a NS to a quark star

(QS), called a quark nova (QN) [49], following the SN explosion of a Wolf-Rayett type star (i.e. a star that is very hot and massive and has strong winds) rich in oxygen. This work is the first in depth exploration of the features of this controversial event.

1 Think of this work as a detective novel. We are trying to find the smoking gun behind the

very violent, mysterious and puzzling death of a star. We have different suspects which we will

investigate further along with their modus operandi. Then we will put all this information together to present a plausible reconstruction of the events. Like Sir Arthur Conan Doyle wrote in “The

Adventure of the Blanched Soldier” [17]:

When you have eliminated all which is

impossible, then whatever remains,

however improbable, must be the truth.

Sherlock Holmes

This dissertation will be structured as follows: In Chapter 1, we start by looking at what we

know. We will go over the basics of stellar evolution and the events that lead up to a SN. Later, we

will discuss what we know about SLSNe, their classification and some of the hypothesized power-

ing mechanisms. In Chapter 2, I present the QN hypothesis and its ability to power SLSNe. Chapter

3 is a presentation of ASASSN-15lh, the most luminous SLSN seen to date. I begin by present-

ing its detection history and the available photometric and spectroscopic data. Then I present the

possible standard powering mechanisms along with their limitations. Chapter 4 presents the QN

scenarios for ASASSN-15lh, their advantages and challenges. These results, namely the photo-

metric and spectroscopic fits to ASASSN-15lh from our model, shed some light into the structure

of the mechanism behind this event. To close this work, Chapter 5 presents the conclusion and

summary.

1.1 Stellar Evolution: The Events Leading up to a Supernova

Standard stellar evolution, that is the birth and the death of a star, is thought to be well under-

stood and generally taught through undergraduate and graduate courses. Here, I present a brief

introduction to stellar evolution and the events that lead up to a SN. This is not a complete de-

tailed description but more of an introduction and refresher for the reader. I direct the interested

2 reader to “An Introduction to the Theory of Stellar Structure and Evolution” by Prialnik [59] and

“Introduction to High-Energy Astrophysics” by Rosswog and Bruggen¨ [63], for informative and

introductory textbooks. This section is loosely based on these books.

A star’s first fuel is hydrogen. This hydrogen is in the core of the star and it is being burnt

into helium. As long as this process is sustained, it is said that the star stays in a main sequence

phase. After some time, few billion years, the star contracts and heats up enough to ignite the

hydrogen shell around a helium core. This is now called a red giant star. More time passes by and,

eventually, the star runs out of hydrogen to burn. This makes the star contract and heat up again

allowing for helium to ignite and start burning, mainly into carbon and oxygen with some other

more rare products like magnesium, neon and silicon.

The fate of the star heavily depends on its mass. If the star is low-mass, that is less than 8 M

(where M is the mass of the Sun), its self gravity does not allow for enough contraction to ignite carbon and oxygen. As a result, a low-mass star will not go beyond this point and it will end its life in what we call a white dwarf, a star with no source of energy which slowly cools and fades away.

What happens with high-mass stars? Well, these are the ones relevant to this thesis. Stars that are more massive than 8 M can proceed and burn carbon to produce more oxygen, magnesium and neon. If the star is more massive, say 9 M to 10 M , then it will be able to burn these products into silicon, sulfur, calcium and argon all the way to the strongest-bound nuclei of all: iron.

The burning process will not proceed beyond iron as the fusion of lighter elements into heavier elements releases energy only for elements before the iron peak. As a result, we have a star that burns in stages and with distinctive layers. The structure of a burning star has layers very much like an onion or an everlasting gobstopper. Figure 1.1 shows the schematic structure of a burning star.

The inevitable contraction of the iron core due to electron capture leads to the collapse of the star into what we call a core-collapse SN explosion. The name supernova comes from nova, new in Latin, after the notion of the sudden appearance of a “new star” in the sky. These outbursts of

3 H burning

He burning H-rich envelope C burning He O burning C O Si burning Si

Fe core

Figure 1.1: Schematic structure of the burning stages of a star. This is the structure of a SN progenitor star.

4 energy have luminosities comparable to that of an entire galaxy. While a typical star like our Sun

has a luminosity of a few times ∼ 1033 erg s−1, a standard SN has a luminosity of ∼ 1044 erg s−1.

SNe are classified into two types based on their observed properties, specifically their spectrum.

Type I are those SNe without hydrogen lines in the spectrum. SNe that have hydrogen lines in their spectrum are called Type II.

Type I SNe are further classified into Ia, Ib and Ic. Type Ia1 are believed to be the result of collapsing white dwarfs that have reached the limiting Chandrasekhar mass due to accretion or coalescence. The possible source of mass for the star to accrete from is thought to be another star.

Although the nature of this companion star remains unclear, it seems that binary systems are the best venue for mass accretion and subsequent detonation of white dwarfs. On the other hand, Type

Ib and Ic SNe are believed to come from a single star. Type Ib SNe show no hydrogen in the spectrum but have hints of helium. Type Ic have no hydrogen nor helium lines.

Type II SNe are associated with the collapse of iron cores of massive stars. The hydrogen lines in the spectrum of Type II SNe is believed to come from their large hydrogen-rich envelopes (see

Figure 1.1). These events are mostly observed in the gas and dust rich arms of spiral galaxies, regions of active star formation and where young stars are abundant. In comparison, Type I SNe are found in all types of galaxies. The luminosity of a SN is dependent on its type. In general, the luminosities range from ∼ 1042 erg s−1 for Type II to ∼ 1043 erg s−1 for Type Ia.

A reasonable question to ask at this point is, what is left behind a SN? To answer that, we need to look at what happens to the iron core that could not keep undergoing nuclear fusion. After undergoing all burning stages, the iron core cannot burn into any other subsequent product. As any inert stellar core, the iron core contracts until the electrons in it become a degenerate gas. The contraction of the iron core goes on rapidly as the degenerate electron pressure is incapable of opposing self gravity. Eventually, the density becomes high enough for the free protons in the core

1Type Ia SNe are fundamentally different objects from the ones considered in this work as they are the result of a binary system and not a singular star. Type Ib/Ic SNe are more closely related to Type II SNe as they all seem to have a single star as a progenitor. The only significant difference between Type II and Type Ib/Ic is the lack of hydrogen lines in the spectrum of the later.

5 to absorb electrons and turn into neutrons. This leads to a degenerate neutron gas that occurs at a density of about 1015 g cm−3; for a comparison, the nuclear matter saturation density is 2.8×1014 g cm−3. This new degenerate gas offers enough pressure to help overcome contraction and the collapse halts.

Then, after a SN explosion we are left behind with an ejected envelope by the explosion and a neutron-rich core that eventually becomes a NS. The NS was first hypothesized in 1932 by Lev

Landau [31, 78]. These objects are compacts stars made of neutrons. Their characteristics like mass and size are dependent on how we model the matter in them (i.e. which equation of state of neutron matter is used) and their structure. The maximum theoretical mass of NS must lie somewhere between 1.4 M and 2.4 M while their radius is about 12 km, roughly the size of Calgary. NSs are interesting and intellectually challenging objects that have been a constant topic of research.

The path of stellar evolution can take a final turn when a massive star leaves behind a too massive neutron core that surpasses the limiting mass for a neutron star and collapses into a black hole (BH). For completeness, I must state that a BH is an object based on the theory of general relativity that results when the radius of a star of a given mass is so small that the escape velocity surpasses the speed of light.

1.2 A New Piece in the Puzzle: Superluminous Supernova Events

Given that the mechanism behind SNe is thought to be well understood, we have constraints on the possible energy output from these types of events. Since the turn of the millennium we have encountered objects that greatly exceed the luminosity of the SNe we had seen, which we call

SLSNe. These are events in the order of ∼ 1044 erg s−1 or larger; recall that a typical type II

SN releases on average ∼ 1042 erg s−1 in luminosity. Understanding them means revisiting our understanding of stellar evolution and of astrophysics. SLSNe are an opportunity for scientists to explore new physics. They are new astronomical laboratories for new powering mechanisms.

6 The study of SLSNe is motivated by their association with the death of the most massive of all

stars, their possible role in element formation in galaxies and the fact that they suggest there are

mechanisms that release a lot of energy that we have not yet discovered. In this section I present

a description of what we know of SLSNe and their characteristics. These events are not included

in standard text books at the moment, so I encourage the interested reader to review Gal-Yam [19]

for a broader picture.

Perhaps, instead of talking about erg s−1, it is easier to think of SN classification by their magnitude. We remind the reader that the magnitude system2 suggests that the more negative the number, the brighter the event is. Using that scale, the cutoff between traditional SNe and SLSNe is about -21 in absolute magnitude. Figure 1.2 presents the luminosity evolution of SNe and SLSNe.

These graphs of luminosity evolution are called light curves and they are our photometric probes for the study of these events. As it is evident from the Figure, some of the SLSNe last for hundreds of days.

Similar to traditional SNe, SLSNe are classified based on the presence (Type II) or absence

(Type I) of hydrogen lines in their spectrum. Here I present the main characteristics of Type I and

Type II SLSNe.

• Type I SLSNe: These SLSNe lack H in their spectra. The early spectrum is relatively featureless

and the lines that do appear are absorption lines from elements like O, Mg, Fe and Ca [19].

Figure 1.3 shows the early spectra of several type I SLSNe events. Examples of this class, and

the first two reported events, are SN 2005ap [60] and SCP 06F6 [2]. Photometrically the events

have extreme peak luminosities (few times ∼ 1044 erg s−1) and large and persistent UV fluxes.

That means that compared to other SLSNe, their UV flux is larger and lasts for several weeks.

2 The magnitude system is useful to measure the apparent luminosity of an object. The main characteristics are that it is a logarithmic measure of brightness, the scale factor is negative and such that a difference of five magnitudes corresponds to an intensity ratio of 100. A star with a magnitude of 1.0 is 100 brighter than a star of magnitude 6.0. There are two types of magnitude, apparent (the brightness of an object as it appears in the sky) and absolute (the brightness of an object as if it was placed at a specific distance, conventionally 10 parsecs). “Astrophysics volume I” by Richard Bowers and Terry Deeming has a deeper and more detailed explanation of the concept of magnitude [7].

7 Figure 1.2: This Figure shows the light curves of SNe and SLSNe. While common SNe can reach peak luminosities of 2 × 1043 erg s−1(absolute magnitude > −19.5), SLSNe reach luminosities that are greater by a factor of 10. The Figure contains prototypical events for each type of SN and compares them to SLSNe PTF09cnd, SN2006gy and SN2007bi. The specifics of the classification suggested by Gal-Yam are discussed in the source. (Figure 1 from Gal-Yam [19].)

8 Figure 1.3: This Figure shows the early spectra of some SLSNe-I. The absorption lines correspond to light elements like O, Mg, Si, and C. Other than the presence of these lines, a characteristic is the absence of H or He lines. Hα has an expected wavelength of 6562.8 A.˚ (Figure 4 from Gal-Yam [19].)

• Type II SLSNe: Just like their SN equivalent, SLSN II have strong H lines in their spectra. These

events seem to be the most commonly observed. One of the events of this type is SN 2006gy, an

event that in 2007 was the most luminous SLSN ever recorded to that date [64].

The nature of these events is highly speculative. Different mechanisms and models have been suggested to power SLSNe. Their strengths are diverse, but with new discoveries, most of them have been pushed to their limits. The main models are described in upcoming subsections.

1.3 Powering Mechanisms for Superluminous Supernovae

Three main mechanisms have been suggested to explain the extremely large luminosities of SLSNe:

SLSNe powered by radioactive decay of 56Ni, SLSNe powered by rapidly rotating magnetars, and

9 SLSNe powered by the collision of the ejecta of a SN with a dense circumstellar material (CSM).

A fourth mechanism called tidal disruption events (TDEs), which is different in nature and not

related to SNe, has recently been suggested as an alternative explanation to the photometric and

spectroscopic signatures of SLSNe. In this section, I introduce the reader to each one of these

mechanisms and refer them to the appropriate sources for further information

1.3.1 Nuclear Decay

56 It is estimated that the explosions of massive stars can synthesize about 0.1 M of radioactive Ni. The γ-ray and positron emission resulting from the radioactive decay chain

56Ni → 56Co → 56Fe is thought to be the energy deposition mechanism. The emission of the

decay chain is thermalized and converted to optical radiation by the expanding SN ejecta [19]. The

decay of this material is thought to energize the SN ejecta accounting for the luminosity of some

SLSNe. However, the nature of the explosion of the massive star in the first place is controversial.

Two mechanisms have been suggested: core collapse SNe or pair instabilities SNe.

A core collapse mechanism could be responsible for the explosion of the star. This is, the iron

core of a very massive star could collapse giving rise to the production of 56Ni [19, 39]. This

56 42 −1 mechanism produces about 0.1 M of Ni which can account for a luminosity of ∼ 10 erg s .

A pair instability SN (PISN) [61] is considered to occur when very massive stars (> 60 M )

9 3 6 −3 have cores of He at high temperature (& 10 K) and low density (∼ 10 −10 g cm ) that undergo rapid production of electron-positron pairs [10]. These pairs allow for the conversion of internal gas energy into their rest mass, reducing the radiation pressure in the core and resulting in its contraction [23]. This contraction, in turn, allows for explosive nuclear burning that produces a strong shock making the star explode. After a PISN, models [20] suggest that huge amounts (3 M -

56 10 M ; 30 to 100 times the values expected in standard SNe) of radioactive Ni are synthesized by the explosion which could account for luminosities of up to ∼ 1043 erg s−1.

10 The study of the spectrum of these events, by modelling it, allows for the reconstruction of the

elemental composition of the ejected mass illuminated by 56Ni decay. These abundances allow for

the estimation of the mass of radioactive 56Ni and the size of the progenitor star. SN 2007bi has been considered an example of this type of event and the model has been explained in detail by

Gal-Yam et al. [20] and Moriya et al. [39].

1.3.2 Circumstellar Material Interaction

This powering mechanism relies on the presence of dense circumstellar material (CSM) surround- ing the event. The main idea is that the ejecta of a SN interacts with a dense CSM. The shock from the collision converts some of the kinetic energy (KE) of the SN ejecta into thermal energy. This thermal energy is eventually released as radiation energy [40].

The nature of this material is unknown and depends on the event. The origin of the CSM has been considered to be mass stripped from a companion star [41] or circumstellar wind of the SN

[40, 77]. This is one of the most versatile mechanisms for SLSNe, primarily because the CSM interaction could happen a long time after the SN allowing for the explanation of different shapes of light curves. Because the CSM can be extended, the interaction can be active for many years allowing to explain long-lasting events [19, 62]. Furthermore, the CSM can explain the presence of some element lines in the spectrum of these events. That is why CSM interaction is one of the methods most commonly used to explain type II SLSNe [42, 40, 19].

1.3.3 Magnetars

The premise behind this mechanism is that a NS that is rapidly rotating could power a SLSN by means of magnetic interaction [76]. Some NSs left behind by a SN are rapidly rotating and have strong magnetic fields. Magnetar is the name given to a NS with a magnetic field strength exceeding 1013 −1014 G [37]. The origin of the magnetic field has been the subject of controversy.

One possibility is that the magnetic flux of the parent star is trapped in the NS. Alternatively, a dynamo effect due to a differentially rotating NS could be the generator of the magnetic field [68].

11 Assuming the existence of such a magnetic field, the energy radiated by the NS depends on the

52 −2 period of its rotation. The rotational energy of the NS is ≈ 2×10 Pms erg where Pms is the period of rotation in milliseconds. The energy loss can be estimated using the Larmor formula

dE 2 2π 4 = (BR3 Sin α)2 ≈ 1049B2 P−4 erg s−1 (1.1) dt 3c3 P 15 ms 6 15 where R is the radius of the neutron star (∼ 10 cm), B15 is the surface dipole field in 10 G, and α is the inclination angle between the magnetic and rotational axes. Assuming an arbitrary angle

the spin-down luminosity can be in the order of 1044 erg s−1 [76].

In a simple model, a regular SN takes place ejecting an envelope of material and leaves behind

a magnetar. As the envelope expands, the magnetar loses rotational energy by magnetic dipole

rotation and this energy is injected into the SN ejecta providing the observed luminosity [27].

1.3.4 Tidal Disruption Events

TDEs are not the explosive end of a star’s life. They are not the powering mechanism behind

SLSNe or SNe. However, the recent observation of SLSN ASASSN-15lh has challenged all mech-

anisms. It has been suggested that this event is a TDE instead of a SLSN [34] and that is why I

briefly review TDEs.

The principle behind these events is that a star can be disrupted (i.e. torn apart) when it passes

within the tidal radius of a supermassive black hole (SMBH), resulting in a luminous flare of

1042 − 1043 erg s−1 [70]. This model relies heavily on the physical properties of the black hole

such as mass and spin. One of the candidate events for a TDE is ASASSN-14ae [26].

1.3.5 Advantages and Disadvantages

Each mechanism in this chapter addresses some important characteristics of SLSNe: the presence

of some spectral lines, the energetics required to power the event, the duration of the event and their

rates. The disadvantages are event dependent. Each observation is different and there is not, so far,

a unique model that explains every single characteristic of all events. Table 1.1 summarizes Section

12 1.3 and shows the standard powering mechanisms and their advantages and disadvantages in their application to SLSNe. This section serves as an introduction to the models and the big picture of how they work. These models are later discussed and analyzed in the case of ASASSN-15lh

(Chapter 3), the most luminous event seen to date.

Table 1.1: Advantages and disadvantages of models for superluminous supernovae Model Advantages Disadvantages

• SLSNe require large amounts of radioactive 56Ni. • A standard SN can produce radioactive 56Ni. Nuclear Decay • The nature of the first explosion is not fully known. • Emission of the decay chain is a known product. • Cannot “bang” twice. • The emission can be thermalized.

• The origin of the CSM is not known. • The time delay between the SN and the collision with the CSM can explain different lightcurve features. • SLSNe require extremely dense ejectas. CSM Interaction • The CSM composition could explain absorption • The interaction with a not extended CSM could result in its lines in the observed spectrum. expansion, rather than its rebrightening.

• SLSNe can require magnetars with extreme values. • Provides a rapidly rotating object with a strong magnetic field. Magnetars • The origin of the strong magnetic field is disputed. • Energy loss due to magnetic dipole rotation energizes • Cannot explain double peaked light curves. the SN ejecta.

• They are not related to the death of a star in its evolution.

• Photospheric radius and effective temperature are • They are not required to explain most SLSNe. different to other SLSNe. Tidal Disruption Events • Their lightcurve is different to that of SLSNe. • The disruption of a star provides a large luminosity. • Large luminosity events might require SMBH spinning in prograde motion.

13 Chapter 2

The Quark Nova as a Mechanism Behind Superluminous

Supernovae

An interesting alternative to power SLSNe and other high energy astronomical events is the QN, the explosive transition of a NS into a QS. In this chapter I present a summary of the QN hypothesis and the relevant models to explain SLSNe.

2.1 The Quark Nova Hypothesis

The existence of strange quark matter (matter made up of up, down, and strange quarks) has been hypothesized as the most stable state of matter in the universe [5, 75]. This hypothesis allows for the existence of compact objects constituted entirely of quark matter; these objects are called (u, d, s) quark stars (QS). The core of a NS could, in principle, reach a high enough density to initiate a phase transition from hadronic matter (protons and neutrons) to quark matter, giving birth to a QS.

The theoretical explosive transition of a NS to a QS is called a QN [49].

The phenomenology of a phase transition in a NS is a fertile area of research. The study of such events could have very important observational and theoretical consequences in astrophysics. Be- cause the phase transition is an exothermic process, the energy could manifest itself observationally in a number of ways. These energetic events could be the explanation behind some astrophysical observations that are currently not well understood [51].

2.1.1 Nuclear Deconfinement and Seeding

The Hadronic-to-Quark-Matter phase transition occurs when hadronic (nucleated) matter under high temperatures and/or densities deconfines into what is called a quark-gluon plasma (QGP) [9].

Under high density and temperature conditions, many neutrons can occupy a reduced space. With

14 their proximity, they start sharing quarks and quark matter is formed [43]. Since the densities

14 −3 inside a NS far exceed the nuclear density, ρnuclear ∼ 2.8 × 10 g cm [73], it is reasonable to assume that these objects might contain matter in a quark phase.

One method in which Hadronic-to-Quark-Matter phase transition can be initiated is through the interaction of strange quark matter (SQM) with hadronic matter [28]. The SQM hypothesis asserts that bulk matter comprised of deconfined u, d and s-quarks at zero pressure is more energetically favorable than the most stable atomic nuclei [6]. To obtain SQM inside a NS, supra-nuclear densi- ties are not enough: that is, nucleons cannot spontaneously dissolve into their constituent quarks.

The presence of s-quarks, adds an extra degree of freedom, lowering the overall energy per baryon and resulting in the SQM phase of matter [56].

The initial amount of strange quarks is referred to as a seed. Currently, it is assumed to proceed via either clustering of lambda baryons1, higher-order neutrino “sparking” reactions2, seeding from the outside [43], or self-annihilation of dark matter particles gravitationally accreted by NSs [58].

The interaction between a SQM seed and hadronic matter can lead to the Hadronic-to-Quark-

Matter transition. In the boundary region surrounding the SQM seed, there will be an overlap with the hadronic matter and it will dissolve into its constituent u and d quarks. This interface region will attempt to equilibrate chemically by producing more s-quarks. The new s-quarks then will diffuse into hadronic regions and more SQM will be formed [56]. This situation can be described as an interface where SQM is on one side and hadronic matter on the other [43]. This is why the presence of s-quarks and their amount at the beginning of the reaction play a very important role in the activation of the phase transition process.

1Lambda baryons are subatomic particles composed of three quarks: up, down, and either strange, charm, bottom, or top. Lambda baryons with a strange quark could be the source of strange quark seeding. For further information on Lambda baryons, refer to a particle physics book or to Astroparticle Physics by Claus Grupen, an introductory book on particles in an astrophysical context. 2Neutrino “sparking” refers to neutrino reactions that produce strange quarks. Further information can be found in Ouyed et al. [56].

15 The equations describing the equilibrium of the interface are:

− d → u + e + νe (2.1)

− u + e → d + νe (2.2)

− s → u + e + νe (2.3)

− u + e → s + νe (2.4)

u + d → u + s, (2.5)

where u, d, and s are the up, down, and strange quarks respectively, ν¯e and νe are electron antineu- trinos and electron neutrinos respectively, and e− is an electron.

At the core of the NS, before the appearance of any seed, there are no s-quarks and process (2.3) does not occur. Simultaneously, process (2.4) is suppressed because of momentum conservation.

Process (2.5), however, can occur because it satisfies momentum conservation and it produces s-quarks [1].

There is a minimum number of s-quarks required for reaction (2.3) to occur. If met, the

Hadronic-to-Quark-Matter phase transition will start, resulting in reaction (2.4). If the number of s-quarks is less than required, then no Hadronic-to-Quark-Matter phase transition can occur.

The presence of a seed by means of an external process can ease the activation of reaction (2.3).

The equations of equilibrium are presented here to give an idea to the reader of how burning works, however, the process is more complicated and goes beyond the scope of my research.

Niebergal [43] explores the transition in detail and is a recommended resource for the interested reader.

An alternative to seeding is what is called a two-step process. In this mechanism hadronic matter is compressed by a shock to high densities that trigger deconfinement into u and d quark matter. The two-flavored quark matter can then decay into three-flavored (i.e. SQM) in a heat releasing process [4].

16 2.1.2 Burning

The specifics of how the transition of hadronic-to-quark-matter takes place are still an active area of research. Different approaches provide different results and here I present a short overview of the different possibilities under which burning could take place. A common approach is to think of this transition as a burning process, where the hadronic matter acts as the fuel which after the transition leaves behind ashes, the quark matter. In this transition, there is a burning front. If the speed of this burning front is higher than the speed of sound we have a detonation, otherwise the combustion is a deflagration.

The transition has been studied as a weak deflagration as a result of the diffusion of seed quark matter [24, 45, 46] with some studies finding the timescale of the transition in the order of 0.1 seconds to a few minutes [45]. Other studies suggest that if the transition were to happen as a two- step step transition, the transition from hadronic to two-flavored quark matter will happen in the order of milliseconds while the conversion to three-flavored quark matter will happen in hundredths of seconds [4]. Other authors disagree and suggest that the transition occurs as a detonation [35].

In any case, the reality is that there is no consensus as of how this process happens and whether it is an explosive transition or not. That is why other groups have decided to develop highly sophisticated tools to compute the hydrodynamics of the combustion of a NS to a QS [25, 36, 44].

The Quark-Nova group has developed Burn-UD, a hydrodynamic code that simulates the micro- physics of the hadronic-to-quark-matter transition, to study this transition in detail and understand how this process takes place. Some simulations show that explosive and core-collapse QNe are possible (see Ouyed et al. [48] and Welbanks et al. [74] for a review).

The assumption that I make for this work is that the combustion of the NS is explosive and that

QNe are a possibility.

17 2.1.3 Energetics

One of the main benefits of the QN hypothesis is the amount of available energy after the transition

of a NS to a QS. This energy comes from basic principles of nuclear physics and particle physics.

The derivations for the energetics of the event have been previously performed by Ouyed et. al.

from a theoretical standpoint [49] and later confirmed with hydrodynamical simulations [44]. Here,

I present an introductory explanation of the principles used to explain the energetics of the event.

There are two main sources of energy after a QN event, the accretion energy that comes to the

object as more and more material is compressed into higher densities and the energy associated

with the conversion of neutrons into quark matter. We can represent that as

MNS 2 EQN ≈ M c (ηacc. + ηconv.) (2.6) M

where MNS is the mass of the ejected envelope by the QN while ηacc. and ηconv. are the the accretion efficiency and conversion efficiency respectively. The accretion efficiency must lay somewhere

between the accretion efficiency of a NS (ηacc. ∼ 0.15) and of a BH (ηacc. ∼ 0.1). The conversion efficiency depends on what considerations are taken for the transition of hadronic-to-quark-matter.

Taking both efficiencies together, we can define ηQN as the efficiency of the QN. Thus the energy released in the QN is

MNS 2 53 MNS ηQN EQN = M c (ηQN) ≈ 2 × 10 erg (2.7) M M 0.15 which exceeds ∼ 1053 erg.

More advanced calculations include other sources of energy as neutrino and photon emission

53 and the effects of the hydrodynamics [57, 71]. These results show that EQN > 10 erg. If only 10%

KE 52 of the energy of a QN was converted to kinetic energy we would have an energy of EQN > 10 KE erg which in turn is > 10 ESN . The calculations shown here serve as an indication for the reader as of the power of the QN model and the reasonability of its assumption.

18 2.1.4 Astrophysical Scenarios and Their Observational Signatures

Assuming that QNe are possible and that they release in excess of 1052 erg of energy in kinetic en- ergy, allows us to consider the observational signatures of such a powerful event. As the aftermath of a NS, QNe could happen in similar scenarios to those of other stars: binaries, collapsing stars, long-lived stars, etc. (i.e. QNe are naturally associated with massive stars). Here, I discuss two compelling scenarios for the presence of a QN that could help us explain SLSNe. A dual-shock

QN and a QN occurring in a binary.

Dual-shock quark nova

First explored more than 15 years ago [32], this model has been revisited and successfully applied

to different SLSNe [29, 50]. The dual-shock QN (dsQN) is the result of a QN happening some

time after a traditional SN.

After an initial SN, an ejecta is released and a NS is allowed to continue its evolution. If the

NS undergoes a QN, it can release a second ejecta which travels close to the speed of light, much

faster than the first ejecta (i.e. the SN ejecta). Depending on the time delay between the SN and the

QN, the QN ejecta catches up with the SN remnant (SNR), collides with it, and reheats it through

a shock. Effectively what happens is that the SN provides the material for the QN to reheat and

re-energize causing a rebrightening of the SN [50]. We are effectively harnessing > 1052 erg of

QN kinetic energy into radiation.

Different time delays offer different photometric signatures in the light curve of the event. If the time delay between the SN and the QN is small (∼ days), the radius of the SN ejecta will be relatively small. If the QN ejecta catches up during the rise of the original SN, the shock can be hidden from direct observation. If the time delay is larger (∼ days to weeks) the radius of the SN ejecta will be larger and it could be optimal for extreme rebrightening. If the time delay is too large

(∼ months) the SN ejecta is too far and diffuse to experience any rebrightening. This is like a QN in isolation and the kinetic energy of the QN is not harnessed.

19 In a way, the photometric signatures of an astronomical event can tell us about the time delay

between the SN and the QN with there being a Goldilocks time, which is just right for maximum rebrightening like there is in ASASSN-15lh. A different photometric signature is a double-hump light curve. The SN will create a hump in the light curve, and the subsequent ejecta shock caused by the QN can create a second hump.

Lastly, the spectrum of such an event would tell us about the environment in which the QN occurred and the composition and type of the original star that underwent a SN. The QN shock can ionize the elements in the SNR, help produce some unique spallation products [29, 47, 50], and allow us to take a look at the element composition of the SN ejecta.

Quark Novae in Binaries

Any undergraduate level textbook tells us that binary stars are very common in our universe [59,

63]. Since QNe are naturally linked to massive stars, we expect them to occur in massive binaries.

The consideration here is, what would happen if one of these stars in a massive binary experiences

a QN? For that, let us consider a series of steps.

First, let there be two massive stars in a binary configuration with masses of approximately

20 M to 25 M [59]. The more massive component of the binary (component A) can undergo its regular stellar evolution path leading to a SN. This will give birth to a NS of ∼ 1.4 M . The remaining star (component B) continues its evolutionary path and turns into a red giant, a star in

its later phases of stellar evolution. Component B will have a developed He core as discussed in

Chapter 1.

The NS that resulted from component A will continue to be in the binary and will be spiralling

into component B. As the NS spirals into the red giant, the H envelope surrounding the giant will

be ejected and our NS will accrete matter during this phase (roughly 0.1 M ). This is known as the ejection of the first common envelope (CE). See Panels A, B, C and D in Figure 2.1.

After the first CE has been ejected, a He-core-NS binary is left behind in a close orbit with a

20 period of a few hours. Our He core continues its regular stellar evolution path and gives us a second

CE phase which is now rich in He. With the new envelope, the system helps the fast in-spiralling

NS accrete more matter. This is shown in Panels E and F of Figure 2.1.

The CE continues to evolve and grow in size and with this, the NS can accrete the required

∼0.3 M - 0.4 M to reach critical density in its core and experience quark deconfinement and a QN [52]. By harnessing the ∼ 1052 erg of kinetic energy of the QN, the system can unbind the CE

and eject it. This will result in a large luminosity that could explain some humps in the observed

light curves of SLSNe. See Panel G and H in Figure 2.1.

Now that the second CE has been ejected, a QS is left behind orbiting quite closely a carbon

and oxygen core (CO core), which was left behind from component B. The separation between

these two objects could result in the QS providing spin-down power to the system if they are far

enough or in a fast merger into a BH if they are close. For illustration see Panel I of Figure 2.1.

After the merger, the bulk of the carbon and oxygen (CO) core forms a disc around the BH. The

accretion of the CO disc into the BH results in a luminosity that can power a long lasting hump of

a light curve. See Panels J, K and L in Figure 2.1.

Finally, the ejected He CE can collide with the previously ejected first CE. This would give

rise to a late CSM interaction. Depending on the distance at which the first CE is and its density,

this interaction can result in a rebrightened light curve or some late hints of H in the spectrum

potentially explaining SLSNe with late H emission in their spectrum.

Ouyed et al. [52] offer a detailed description of the steps in this binary model. Figure 2.1 shows

a simplification of the steps explained above in the binary system.

The benefits of this model are multiple. First, there is no exotic physics behind this model.

These are processes that are well established in stellar evolution. Second, the ejection of the first

CE (H rich) could explain the lack of H lines in the early spectrum and potentially the appearance

of H lines in late spectra (type I SLSNe). Lastly, this sequence of events very well could explain

multiple rebrightening epochs (i.e. bumps in the light curve).

21 Figure 2.1: A massive binary can evolve into a QN. A subsequent BH can accrete matter from the system rebrightening the light curve. The 1x and 10x in Panels g and G correspond to 1 and 10 times zoom respectively. (Figure 1 from Ouyed et al. [52].)

22 Chapter 3

ASASSN-15lh. The Most Luminous Supernova Event Ever

Seen

All models for SLSNe were tested and pushed to the extreme when ASASSN-15lh was discovered.

Not only was this event the most luminous ever seen, but some of its features were not possibly described by traditional models. Its late rebrightening in the UV, its energetics and spectrum are three pieces of a puzzling challenge that still remain unsolved. In this chapter I present the story of ASASSN-15lh, its discovery and data.

Explaining ASASSN-15lh is key in the understanding of high energy explosive astrophysics.

This is the first extremely energetic event, of presumably many more, that pushes theoretical mod- els to the extreme. It has forced some previously successful models to go beyond the limits of their parameters and others to be completely ruled out. An event like this was needed to help us find the successful models for SLSNe. The advantages and disadvantages of the models explained in

Chapter 1 will be revisited in Section 3.2.

3.1 Detection History: Photometric and Spectroscopic Data

On June 14th, 2015 the All-Sky Automated Survey for SuperNovae (ASAS-SN) triggered on a new source which was later reported by Dong et al. [16] as the most luminous SN yet, reaching a peak

45 -1 bolometric luminosity of Lbol. = (2.2±0.2)×10 erg s . With the light curve, an optical spectrum was published. The spectrum contained wavelengths from 3700 A˚ to 9200 A.˚ The event was mostly featureless except for a deep broad absorption near ∼5100 A˚ in the observer frame which seemed to resemble an O II absorption line at roughly 4100 A˚ in the rest frame [16]. This spectrum and the lack of H or He lines made this event consistent with the type I SLSNe classification. Figure

23 Figure 3.1: Figure showing the rest-frame absolute magnitude light curve of ASASSN-15lh near peak compared with other type I SLSNe events. The comparison contains the most luminous events previously known showing that ASASSN-15lh is more luminous than the other events by roughly more than 1 magnitude. Dong et al. [16] contains a full description of the corrections performed to obtain this image. (Figure 4 from Dong et al. [16].)

3.1 shows the reported light curve (absolute Rest-frame u-band AB Magnitude) for ASASSN-15lh and compares it with other known type I SLSNe events. It is possible to observe that the event lasted more than 80 days and it is evident from this picture how much brighter the event was compared to other events.

On May 2016, Godoy-Rivera et al. [21] published a paper on the unexpected, long-lasting, UV rebrightnening of ASASSN-15lh. The rebrightening began after roughly 90 days on the observer frame and was followed by a 120 day plateau in the bolometric luminosity. The event started to fade again after 210 days. Figure 3.2 shows the light curves of ASASSN-15lh in the observer frame from pre-maximum dates throughout 302 days after maximum. The peak, decline, valley, rebrightening and fading features seem evident in this image.

24 Figure 3.2: Multi-band light curve of the rebrightened ASASSN-15lh. The magnitudes are in the Vega system and were not corrected for extinction or contribution from the host. For more information on methods read Godoy-Rivera et al. [21]. (Figure 1 from Godoy-Rivera et al. [21].)

25 To make things even more puzzling, the rebrightening and later decline in the spectra did not

show any hints of an interaction with a dense circumstellar material (i.e. no CSM signatures).

There were some hints of a weak Hα emission line but the feature was not conclusive and could be an artifact. Figure 3.3 shows the rest frame spectroscopic evolution of ASASSN-15lh on the top panel and on the bottom the rest frame spectroscopic evolution near Hα. A feature around

∼ 4100 A,˚ which was thought to be an O II absorption line, is definitely present on the spectrum

at 19 days (peak) while the Hα line is quite weak and could be argued to be an instrumental

artifact. Nonetheless, the possible presence of H in the spectrum would challenge the classification

of ASASSN-15lh as a type I SLSN event. With this information it seems clear that ASASSN-15lh

is the most extreme event seen yet. The spectrum and UV rebrightening are the hints that will help

discern the successful models.

Other pieces in our understanding of the event are the evolution of the apparent photospheric

radii and effective photospheric temperature of the event. These allow us to compare them to other

known events and distinguish its evolution not only in terms of its spectrum or light curve. Figure

3.4 show the data reported by Godoy-Rivera et al. [21].

3.2 Standard Powering Mechanisms

Since its discovery, ASASSN-15lh has invited scientists to explain the powering mechanism of the

event with all possible theories. Here I present a summary of the work up to date.

In the original publication about the discovery of ASASSN-15lh [16] it is suggested that the

decline rates of the luminosity after the peak seem to be too fast to be explained by the radioactive

56 decay of Ni. To produce such a light curve an enormous mass, anywhere from & 30M [16] 56 up to ∼ 1500 M [30], of Ni would be required. Such a large progenitor is very unlikely as it would be more likely to collapse directly into a BH rather than undergo a SN [12].

26 Figure 3.3: Top Panel: Spectroscopic evolution of ASASSN-15lh at the rest frame. Different curves represent different stages of the light curve evolution. (Figure 2 from Godoy-Rivera et al. [21]). Bottom Panel: Spectroscopic evolution of ASASSN-15lh near Hα. The spectra cover a period of time before and after the fading stage at around 196 days. The green line at 6562.8 A˚ shows the expected position of Hα. For both panels, offsets were added to the spectra for clarity. (Figure 3 from Godoy-Rivera et al. [21].) Godoy-Rivera et al. [21] explain in detail the dates of observation and instruments used for observation.

27 4.5

4.4

4.3 ) K /

T 4.2 ( 0 1 g o

l 4.1

4

3.9 Photospheric Temperature Data

0 50 100 150 200 250 Day since maximum (rest frame)

16

15.8 ) 15.6 m c / R ( 0

1 15.4 g o l 15.2

15 Photospheric Radii Data 14.8 0 50 100 150 200 250 Day since maximum (rest frame)

Figure 3.4: Top: Effective photospheric temperature of ASASSN-15lh. Bottom: Apparent photo- spheric radii of ASASSN-15lh. Data obtained from Godoy-Rivera et al. [21].

28 As for the possibility of a magnetar powered event, ASASSN-15lh would require a magnetar

with a period P ' 1 ms or less (i.e. extreme spin), a magnetic field1 of B ' 1014 G, and assume

that all the spindown power is completely thermalized in the SN envelope [16]. Other models

suggest a weaker magnetic field B∼ 1013 G but cannot agree on a value for the ejecta from the

SN with values ranging from 3 M [38] to 6 M [3]. There are other models that suggest even

more extreme values with periods of P= 0.7ms and ejecta masses of 14 M [65]. Other authors challenge this scenario because they find that the energy from the magnetar is transformed into

kinetic energy rather than luminosity [72]. The largest problem with these models is that they fail

to reproduce the rebrightening in the lighturve [12]. The magnetar model cannot easily explain

ASASSN-15lh and current models require the most extreme parameters.

Other work has been done considering a combination of scenarios, the so called hybrid models.

The combination of several luminosity inputs has been studied and it is capable of reproducing the

early phases of the light curve and the rebrightening [12]. These models use a combination of

radioactive decay, magnetar spin-down and CSM interaction. The weakness of this approach is the

amount of fine tunning required to obtain the light curve of the event and their failure to reproduce

the spectrum of the event.

Lastly, there is the idea that ASASSN-15lh is not a SLSN but a TDE [34]. This model considers

8 a rapidly prograde spinning black hole with a mass of > 10 M that traps and disrupts a star with

∼ 1 M outside the event horizon of the event. The main reason for ASASSN-15lh as a possible TDE rather than a SLSN is the temperature evolution of the event resembling that of a TDE rather than a SLSN. However, the light curve resembles a SLN more than a TDE (see Leloudas et al. [34] for a comparison of light curves, temperature evolution and photospheric radii). The extreme BH mass and the fact that is spinning in a prograde direction is a limitation to the model as this type of

BH is not common. Lastly, it is important to remark that no other SLSN to this date has required the TDE model to explain its features.

1 A typical NS has a field of 1012 G [63].

29 Chapter 4

The Quark Nova Model for ASASSN-15lh

In this chapter I introduce my work and that of the Quark-Nova group in trying to explain the observations of ASASSN-15lh. We need to look at the sequence of events that could help us explain the extreme photometric and spectroscopic signatures of the most luminous SN event ever seen. To begin, I will pose my contribution to the binary model and my work including spin- down power to the QN in massive star binaries’ sequence of events. I present the light curves generated and then I discuss some of the problems with this model. Then, I stage the dsQN model in the wake of a WO-star SN. I also offer a simulated spectrum of the event using SYNOW, a synthetic spectrum simulator software. The dsQN model for ASASSN-15lh as presented here was developed by Rachid Ouyed and my contribution was to explore the implications of this model to the spectroscopic signatures of the event.

4.1 Quark Novae in Massive Star Binaries

The binary model for a QN is explained in Section 2.1.4. Here I present a detailed description of the model as it was applied to ASASSN-15lh.

4.1.1 The Quark Nova Luminosity

To calculate the luminosity of an event we require more than the energy of the event; we need to know how the photosphere of the QN shocked envelope behaves. An extensive look at the behaviour of a QN photosphere has been performed in the appendix of Ouyed et al. [53]. First, we need to look at the behaviour of the photons in the photosphere. As the photosphere diffuses

30 outwards, it (i.e. the light) seems to move inwards. The photosphere radius is then:

Rphot.(t) = RCE(t) − D(t), (4.1)

where RCE(t) = RCE,0 +vQNt, t is time, and RCE,0 is the CE radius at the QN shock breakout (t = 0 in our model). D(t) corresponds to the photon diffusion length and it is given by

2 2 c D(t) = D0 + t, (4.2) nCE(t)σth where nCE(t) is the number density in the CE, σth is the Thomson cross section, and D0 is the initial diffusion length defined by setting nCE,0 σth D0 ≈ 1. The initial volume and number density in the CE is defined as nCE,0 [32]. The model has a free parameter which is the initial QN shock heating per particle, and it is defined as (3/2)kBTQN,sh, where TQN,sh is the QN shock temperature and kB is the Boltzmann constant. Then, the heat is redistributed into gas and radiation giving a temperature TCE,0 for the CE after the QN shock has crossed the envelope. The relevant equation is

3 3 k n T + a T 4 = k n T (4.3) 2 B CE,0 CE,0 rad CE,0 2 B CE,0 QN,sh

with nCE,0 being the number density of electrons and ions, and arad the radiation constant.

2−α After the shock, the temperature of the CE evolves as Tcore(t) = TCE,0(RCE,0/RCE(t)) QN,sh. .

The model introduces αQN,sh., which parameterizes shock physics to account for a non-uniform initial temperature.

The corresponding luminosity for a QN event is then given by

dD(t) L (t) = c (T)∆T (t)n (t)4πR2 (t) , (4.4) QN V,tot. core CE phot. dt

where cV,tot. represents the total specific heat, the addition of the specific heat from gas and radia-

tion effects. We also assume that ∆Tcore ∼ Tcore since the photosphere cools promptly.

31 4.1.2 Spin-Down Interaction

The double humped light curve of ASASSN-15lh seemed to require an extra powering mechanism,

beyond the QN and the BH accretion. A QN by itself could not provide such a long lasting first

hump, and the second hump could not be explained with CSM interaction due to the lack of H

lines in the spectrum. As a result, my contribution was to include the spin-down interaction of

the QS to power the first hump of the light curve. As the NS in-spirals into the second CE, it

accretes matter, angular momentum, and its rotation period increases. Eventually, the NS gains

enough mass to undergo a QN explosion. After the QN, we obtain a QS with a large magnetic field

which is slowing down rapidly as it is no longer accreting matter. The QS is free to release up to

52 −2 ESpD = 2 × 10 PQS,ms erg in rotational energy [52]. The light curve of a spin-down QS was based on the work of Chatzopoulos et al. [13] which

states that the light curve of an event powered by a spin-down input is given by

Z x 0 2ESpD −[x2+wx] [x02+wx] x + w 0 L(t) = e e 0 2 dx , (4.5) tp 0 (1 + yx ) where x = t/td, w = R0/(vtd) and y = td/tp with R0 being the initial radius of the progenitor, td an effective diffusion time and tp the characteristic timescale for spin-down which is dependant on the strength of the magnetic field. Section 2 of Chatzopoulos et al. [13] explains in detail how to estimate tp given a moment of inertia. The code implemented uses this prescription and assumes

2 −2 tp ∝ PQSBQS where PQS and BQS are the period and magnetic field of the rotating QS. When the QS merges with the CO core giving birth to the BH, the spin-down power has to stop as there is no longer an object spinning down. The solution for times after spin-down power has stopped (tmax) is given by

−β L(t > tmax) = e L(tmax) (4.6) where β is the ratio defined as (t −tmax/td) and L(tmax) is the luminosity from equation 4.5 evalu- ated at tmax.

32 4.1.3 Black Hole Accretion Phase

Different models have been put forward to explain the luminosity of an accreting BH. This work

is a combination of the model by Ouyed et al. [54] complemented with the equations presented in

the appendix of Dexter et al. [15].

The resulting QS from the QN event spirals towards the CO core of the companion. It keeps

getting closer until they merge and form a BH with a CO disc around it. The resulting BH can re-

energize the ejected He envelope and can power a long lasting hump in the light curve by accreting

matter from the CO disc that was formed. The life of the BH accretion phase depends on the time

it takes for the BH to consume the mentioned CO disc. This time is given by

Z t +∆t  −nBH 0 BH t 2 L0 dt = ηBHMCOc , (4.7) t0 t0

where ηBH is the accretion efficiency, L0 is the BH initial accretion luminosity, and nBH is a param-

−n eter defining the injection power L0(t/t0) BH . Dexter et al. [15] give the semi-analytic solution for t > t0 for accretion powered light curves of this form as

 nBH  n/2 t0 −t2/2t2 1 L(t) = L0 e d − td 2   2   2  nBH t0 n t γ 1 − ,− 2 − γ 1 − ,− 2 , (4.8) 2 2td 2 2td

where γ(s,x) is the lower incomplete Gamma function. Notice that in this case, t0 corresponds to t = 0 which is the birth of the BH. There is a natural delay between the QN event and the birth of

the BH, which I have called tBH and I refer to it as the BH delay. In general, it is a free parameter and determines spin-down power.

4.1.4 Light Curve

Using the model described, I performed a best fit to the light curve data obtained for ASASSN-

15lh. The best parameters are described in Table 4.1 and Figure 4.1 shows the light curve obtained

33 Table 4.1: Best fit parameters for the ASASSN-15lh light curve in the binary model QN SpD BH accretion MHe (M ) RCE,0(R ) P(ms) B15 tBH,delay (days) L0 (erg/s) nBH 8.0 1500 0.70 0.075 84 1.5 × 1046 1.8 Note: The magnetic field in the spin-down power prescription is given in units of 1015 G.

Observations 11.75 Model luminosity BH contribution QN contribution 11.5 SPD contribution ) ⊙ L /

L 11.25 ( 0 1 g o l 11

10.75

10.5 0 50 100 150 200 250 300 Time since QN (days)

Figure 4.1: Light curve for ASASSN-15lh using the QN model in a binary. Obtained light curve for ASASSN-15lh. The purple stars represent the observations taken from Godoy-Rivera et al. [21]. The red line shows the result of the model while the contributions of each engine are shown separately. In blue the QN, in green the SpD contribution, and in black the BH accretion. by the model.

A code was written by the Quark-Nova group to solve the equations of each part of the model and the parameters were modified in a manual parameter space looking for the optimal minimized

χ2. The best set of parameters give a minimized χ2 of 1.069.

The spin-down contribution of the QS is capable of reproducing the high luminosity of ASASSN-

15lh while the BH accretion provides a long lasting event that powers the light curve for over 100 days. Although this is a considerably good fit, there were some complications that this model could not solve.

34 4.1.5 Obstacles and Challenges

Although this fit might seem to be a convincing explanation to ASASSN-15lh, we need to consider all the other puzzle pieces in the event. There were other considerations for the event.

An attempt was made to self consistently calculate the values of the photospheric radius evolu- tion and effective photospheric temperature evolution of the event. The results were inconclusive and this prescription provided a photospheric evolution that was in disagreement with the observa- tions.

The parameters found to provide the best fit, although appropriate, were very similar to other models like the magnetar model. Such a short period QS is almost as difficult to obtain as is the case of the NS in the magnetar model, and in principle, the QS would be virtually indistinguishable from the magnetar. In other words, this model does not create a better explanation to the prob- lem. Considering that one of the main criticisms of the magnetar model is the unlikelihood of its parameters, it seems counterproductive to put forward a model that uses the same type of values.

Lastly, although the model reproduces the light curve, it does not have any direct way of ex- plaining the observed spectrum of the event. Perhaps if different elements were abundant in the envelope, or if the CO core was disrupted by the QN in order to eject part of the oxygen faster than the shock so that O was present in the material surrounding the event, we would be able to explain the spectrum. However, either explanation is unlikely. At this point, it was necessary to think that the simplest solution is probably the correct solution.

4.2 Dual-shock Quark Nova With a Wolf-Rayet Star Supernova Progenitor

A natural way of explaining the presence of oxygen was found in the dsQN scenario. If the pro- genitor star was one with a high content of oxygen, then the envelope surrounding the QN would have the chemical composition to explain the puzzling spectrum of ASASSN-15lh. As a result, two parallel endevours were started: the energetics of a dsQN and the spectrum simulation for such an event. Here, I explain the model of the dsQN and our spectrum simulation, which was my main

35 contribution to this research. Later, I present the results and their implications.

A Wolf-Rayet star is a star that is very hot and massive and has strong winds [63]. Named after

the French astronomers Charles Wolf and Georges Rayet, these stars are classified based on their

strong broad spectral lines: WN, WC and WO for nitrogen, carbon and oxygen lines respectively.

A WO type star is known for its oxygen absorption lines and is considered to be an extremely rare

type of star [69] (see de Loore et al. [14] for a review on Wolf-Rayet stars).

As the WO-type star evolves, it experiences mass loss of hydrogen and helium from its outer

layers and shows a large amount of He-burning products, oxygen being the main one [14]. Like any

other star, at the end of its life, it can experience a SN. The SN will eject a mass MWO with a kinetic

SN 51 energy EKE ∼ 10 erg typical for a SN. The velocity at which the ejecta travels can be obtained 2 given the kinetic energy of the event. Start with the definition of kinetic energy EKE = 1/2 mv . Then, for this SN ejecta in a spherical geometry, we obtain

1 Z ESN = v2 ρ4πr2dr (4.9) KE 2 SN where ρ is the density of the material ejected, vSN is the velocity of the SN ejecta and r is the dummy variable to integrate the mass of the material ejected.

Solving the integral and substituting MWO we obtain

3 ESN = M v2 (4.10) KE 10 WO SN q SN which implies that the velocity of the ejecta can be calculated as vSN = (10/3)EKE/MWO once the kinetic energy of the explosion and the ejected mass are known.

The time delay (tdelay) between the SN and the QN allows for the SN envelope to be extended and receive the energy from the QN (rather than losing energy due to PdV losses, that is work used

in extending the envelope). With the velocity of the ejecta and time delay between the events, it is

36 possible to calculate how extended the envelope of the SNR is when the QN occurs. It is

Renv.,0 = RWO + vsntdelay ' vsntdelay (4.11)

where RWO is the radius of the WO star at the moment of the SN and it is assumed to be much

smaller than the extended envelope at the time of the QN (RWO << Renv.0). The main assumption here is evidently that a QN can occur. If that is assumed, a QN occuring

after the SN explosion of a WO-type star can account for the intriguing features of ASSASN-15lh.

The second assumption here, is that the SNR has a two-component configuration composed of a

dense inner core and a lower density extended envelope. The QN will heat the dense core first

and the envelope second. What the observer will see is a hot, shocked envelope which then cools

down as the photosphere recedes. Later, once the photosphere reaches the deeper core, we see a

rebrightening in the light curve as the core is now exposed and radiating. That is, the core remains

insulated until it is exposed by the receding photosphere. This assumption is not so extreme if

we consider that the envelope and the core have different densities. Figure 4.2 shows a schematic

representation of the sequence of events in our model.

To quantify the two step density, we employ a two-piece stepwise density profile constituted of

two uniform density components: a core of mass Mcore and an envelope of mass Menv. = MWO −

Mcore. This choice gives us one more parameter which is the initial size of the core at the time of

the QN wake (Rcore,0). Then the core’s expansion velocity, assuming a homologously expanding ejecta is simply given by

Rcore,0 vcore = vSN . (4.12) Renv.,0 My work presented here assumes that the two-step configuration does not have an effect on the

1/2 ejecta velocity vSN. Ouyed et al. [55] show that the ejecta velocity is divided by a factor of ζ

37 Figure 4.2: Schematic representation (not to scale) of the sequence of events in our model. The top left panel shows a WO star that later undergoes a SN in the top right panel. The middle left panel shows the two-component configuration of the WO-SNR composed of a dense core and a less dense envelope. The middle right panel shows the QN that results in the birth of a QN and a QN ejecta. The QN ejecta shocks the core and the envelope. The envelope later cools and recedes. The bottom left panel shows the photosphere receding through the envelope resulting in the first hump in the light curve. The bottom right panel shows the photosphere receding through the hot core which rebrightened the light curve resulting in a second hump.

38 where

 2 M R , M ζ = env. + core 0 core . (4.13) MWO Renv.,0 MWO

When choosing fiducial parameters Menv./MWO ≈ 1 and Rcore,0/Renv..0 << 1, then ζ < 1 which

means that a two-component profile will yield a slightly larger vSN . Once the QN occurs, and understanding that the velocity of the QN ejecta is much faster (i.e.

relativistic) than the velocity of the SN ejecta [55], the extremely dense and relativistic ejecta

catches up very quickly (within hours) with the SN ejecta which has been expanded for a time tdelay (days to weeks). The energy of the QN is efficiently transformed into thermal energy in the SN ejecta rather than expanding even further the envelope. This implies that after the shock, the

velocity of the envelope is unchanged venv. = vSN. As the shock from the QN travels through the SNR material, part of the kinetic energy of the

QN QN ejecta (EKE ), is converted into thermal energy in the core (Ecore) while the rest is deposited in QN the overlaying envelope (Eenv. = EKE − Ecore). The energy deposited in these two components is diffused in the envelope first, giving rise to the first hump of the light curve, and later it is diffused

in the envelope, producing the second hump in the light curve. The analytic form of the luminosity

of each component comes from solving the general full output luminosity of the photosphere of a

SN ejecta1.

The solution takes the form of L = E/t where L is the luminosity of the component, E is the

energy of the component and t is the shock duration or the time it takes for the shock to traverse

the medium. That way, if the shock has not broken out from the envelope, the luminosity of the

event comes from both the envelope and the core. If the shock has traversed the envelope but the

photosphere has yet to recede deep enough to the core radius, the contribution comes from the

core which acts effectively as a “hot plate”. Once the photosphere crosses the core, the envelope is

optically thin and the core radiates its remaining energy.

1See Section 2.1 in Chatzopoulos et al. [11] and the full derivation of the general light curve model for centrally located power sources in Appendix A of the same paper.

39 In this particular model, the photospheric radius in the envelope (and similarly in the core) is

calculated as

Renv.,ph. = (Renv.,0 + vSNt) − αλ λenv. (4.14)

where λenv. = 1/(κenv.ρenv.) is the photon mean-free path for an envelope with density

3 ρ (t) = M (R + v t)−3 (4.15) env. env. env.,0 SN 4π

and κenv. being the mean opacity in the envelope. The factor αλ is meant to emulate the spherical geometry of an extended SN ejecta. Other considerations of the model like diffusion times and the self consistent derivation of the break out shock from the envelope are omitted here as they are fully explained in Ouyed et al. [55].

4.2.1 Light curve

To summarize, the free parameters of this model are:

SN • The SN parameters: These are the kinetic energy of the SN (EKE) and the mass of the ejecta

(MWO). The kinetic energy of the SN is used to calculate the velocity of the ejecta as shown in Equation 4.11.

QN • The QN parameters: The kinetic energy of the QN ejecta (EKE ) and the time delay between the

QN and the SN (tdelay).

QN • The envelope parameters: The amount of energy deposited in the envelope (Eenv./EKE ), the

envelope mass (Menv.) and the envelope’s mean opacity κenv..

• The core parameters: The core initial radius in relation to the initial radius of the envelope

(Rcore,0/Renv,0), and its mean opacity in the core (κcore).

The results obtained using this model are shown in Figure 4.3. Table 4.2 shows the parameters employed to generate these results. The values for each parameter were determined by doing

40 Table 4.2: Best fit parameters for the ASASSN-15lh light curve in the dual-shock model SN QN ENVELOPE CORE SN QN QN 2 −1 2 −1 EKE (erg) MWO(M ) EKE (erg) tdelay Eenv./EKE Menv.(M ) κenv. (cm g ) Rcore,0/Renv.,0 κcore (cm g ) 2.5 × 1051 3.5 1.7 × 1052 8.5 0.68 1.0 1.2 0.09 2.4

a best fit by observation and modifying the parameters visually. Presumably, performing a χ2 minimization parameter search would yield similar results. The red line in Figure 4.3 shows the output luminosity of the event which is first dominated by the contribution of the shocked envelope.

As the event cools down, the photosphere recedes and exposes the core which then powers the luminosity of the event.

The middle panel in Figure 4.3 shows the photospheric radius evolution, while the bottom panel shows the effective temperature evolution. The evolution of the photospheric radius captures the receding photosphere which then seems to grow as the exposed core radiates the energy deposited by the QN. The effective temperature is derived from the blackbody luminosity as

Lout Teff. = 2 (4.16) 4πσSBReff. where σSB is the Steffan-Boltzmann constant. An evident feature in our model is the presence of sharp features that show the transition be- tween energy output from the envelope and the core. These features do not seem natural and are an artifact of the imposed piece-wise function for the luminosity of the powering engine. However, this is not necessarily a flaw of the model. A more rigorous hydrodynamical simulation of the SN ejecta would be required to capture the nature of the density difference between the core and the envelope. Furthermore, the approach here is to provide a sensible explanation to the event and in- cite further research on this type of explosions. Despite the simple approach, the model is capable of capturing the general photometric signatures of ASASSN-15lh. The remaining puzzle piece is the spectral signatures of the event.

41 Observations 11.75 Model luminosity Core contribution Envelope contribution 11.5 ) ⊙ L /

L 11.25 ( 0 1 g o l 11

10.75

10.5 0 50 100 150 200 250 300 Time since QN (days)

16

Observations Model luminosity 15.8 Core contribution Envelope contribution

) 15.6 m c / R ( 0 1 g

o 15.4 l

15.2

15 0 50 100 150 200 250 300 Time since QN (days)

4.5 Observations Model luminosity Core contribution Envelope contribution 4.4 ) K /

T 4.3 ( 0 1 g o l 4.2

4.1

4 0 50 100 150 200 250 300 Time since QN (days)

Figure 4.3: Results for the features of ASASSN-15lh using the dsQN model are shown with solid red lines. The blue line (black line) shows the envelope (core) contribution to the light curve. The data is represented using coloured stars. Top: The dsQN model fit to the light curve. Middle: The results for the photospheric radius evolution. Bottom: Effective temperature evolution. Data obtained from Godoy-Rivera et al. [21].

42 4.2.2 Spectrum

As discussed in Section 3.1, the spectroscopic data was an odd piece of the puzzle. An absorption line was detected near 4100 A˚ in the rest frame and it did not exactly match any other features previously seen. In general, astronomers do not know what absorption lines they are observing by pure inspection. Different factors like redshift, observation conditions, host galaxy and type of event could give similar absorption lines. Generally what is done is that different simulations are run, and are then matched with the observations [8].

In the case of ASASSN-15lh, it was assumed that given its initial classification as a type I SLSN

(i.e. no hydrogen) the observed line would be a result of an OII doublet seen in other SLSNe [19].

Simulating the event as a SN with a higher temperature, it was impossible to reproduce the feature using only OII as it is always a doublet while the one observed in ASASSN-15lh is a single line

[34].

This highlights a considerable problem when studying the spectrum of the event of interest.

First, it is necessary to assume the nature of the event, SLSN, TDE, etc. Once we assume the powering mechanism, we should simulate the spectrum given the specific conditions of the mech- anism. However, to this day, there is no specialized software to simulate the spectrum of SLSNe or TDEs [34]. As a result, further assumptions need to be made.

For my work, I assume that ASASSN-15lh’s spectrum can be modelled with the chemical com- position of a SN which is reheated by a QN shock. I also assume that to first order approximation the spectrum will be similar to that of a SN. This is not a completely unreasonable assumption if we consider that the QN will not change dramatically the chemical composition of the ejecta and that instead it will re-energize it.

Starting from the assumption that the QN occurs in a WO-star SNR, I consider that there is a high abundance of oxygen. The oxygen in the SNR is then hit by the QN shock and allows for the ionization of the species present. As for the species present, I assume there to be OII and its ionized product OIII. The reasoning for this is that a WO-type star exposes oxygen to its surface,

43 and there are strong absorption lines of oxygen in its SN spectrum [14]. Consequently, the SNR

will be abundant in OII providing the material to be ionized. Alternatively, even if the source

of the oxygen is unknown (i.e. not a WO-star), there are significant observations that show that

hydrogen-less SLSNe have OII absorption features in their spectrum [19]. As a result, if a high

energy shock was present, as in the case of the QN model, the OII in a SLSN could be ionized to

OIII.

To test this assumption, I decided to use SYNOW [8]. SYNOW is a software traditionally used

for the spectrum of SNe. Since from our fits we know that ASASSN-15lh is SN related, I grant that

SYNOW’s considerations are suitable. SYNOW assumes spherical symmetry, sharp photosphere,

LTE excitation and homologous expansion, all of which are true for our model. Its main use

is to take multiple line scattering into account and help the scientist make line identification by

comparing the simulated spectrum to the obtained data. Although other pieces of software exist

which are more complex (i.e. SYN++ [66] or SYNAPPS [67]) SYNOW is preferred for this study

due to its relative simplicity to operate and fast computational speed.

SYNOW approximates continuum transport using a sharp photosphere that travels at a parametrized

velocity (vphot). This photosphere acts as a blackbody with a temperature (TBB). It also uses the Sobolev approximation, the assumption that the rapid expansion of the atmosphere of the star dom- inates radiative transport [18]. That is, random motion due to thermal energy can be neglected as the macroscopic flow velocities are larger [22]. This assumption allows for radiative transfer to simplify as two local problems. The first is to find the probability of escape for a photon at a particular location and the second is to find the amount of light received at that location.

To solve the first problem, SYNOW receives a parameter called optical depth (τ), which is the negative common logarithm of the probability of escape of the photon at a given location. This parametrized optical depth is given for a reference line for each ionization stage of each element

(i.e. species). Then, SYNOW determines the relative optical depths of other transitions of the same species by assuming that the transitions are in LTE at some parametrized temperature (Texc) given

44 by the user. To solve the second problem, SYNOW assumes that the number of photons entering

a location is the same number of photons exiting that location, hence ignoring any sophisticated

model for photon creation or destruction. Another relevant assumption for SYNOW is that it

calculates the radial dependence of the optical depth of a transition, once the reference optical

depth was given, using a power law, exponential or Gaussian function as required by the user.

However, different functions give similar results [18]. Lastly, SYNOW receives a database of 40

million atomic transitions.

For each species (i.e. specific ions), the user needs to provide SYNOW with Texc. and the

velocity range in which the ion is present (vmin. and vmax.). The velocity range allows SYNOW to calculate whether a line is detached or not. For instance, any ions at a velocity lower than the

velocity of the photosphere will not absorb any photons as the photosphere is beyond it. If the

velocity is larger than the maximum velocity of the photosphere, then these ions are detached.

This description of SYNOW is not an extensive or in-detail work because this software is

independent of my research. It is presented here as a guide for the reader and to demonstrate that

the main assumptions of SYNOW are not in conflict with the assumptions in our model. For the

interested reader, a complete and detailed description of SYNOW (and SYNNEW a more recent

version of SYNOW) along with the source code is provided Fisher [18].

To begin our simulation of the spectrum, I focused on the spectrum at ∼ 26 days post maximum as it evidently shows the absorption feature near ∼ 4100 A.˚ I begin by simulating an envelope composed mainly of OII. Figure 4.4 shows the result of an envelope with only OII with an opacity small enough to have the OII doublet present. A smaller opacity would make the feature disappear and it would not match the ASASSN-15lh spectrum.

The results of this first simulation show that modelling the spectrum of ASASSN-15lh as an event with a large amount of OII (and as a result a large optical depth) is not possible as a com- panion feature at ∼ 4400 A˚ is always present. Figure 4.4 shows an inlet with a closer look at the double absorption feature. The second absorption feature at ∼ 4400 A˚ is not present in the

45 1 QN model (OII only) ASASSN-15lh spectrum 0.8 0.25

x

u 0.2 l 0.6 F 0.15 d e

z 0.1 i l 0.4 a 0.05 m r

o 0

N 0.2 4000 4200 4400 4600

0

3000 4000 5000 6000 7000 8000 Wavelength (Å)

Figure 4.4: Synthetic spectrum of ASSASN-15lh using SYNOW at ∼ 26 days post peak luminosity assuming an envelope with OII only. The optical depth used was τ(OII)= 0.2. The inlet shows an amplified look at the double feature of OII. The absorption feature at ∼ 4400 A˚ is not present in the observed spectrum. Data obtained from Leloudas et al. [34].

46 Table 4.3: Elemental species parameter for synthetic spectrum of ASASSN-15lh using SYNOW. The envelope was assumed to be fully mixed giving all elements the same excitation temperature

OII OIII MgII Texc. vmin. vmax. km km τ τ τ K s s 0.02 1.0 0.2 12,000 19,000 23,000

observed spectrum. This result is in agreement with the supplementary material to Leloudas et al.

[34]. The other parameters used for this simulation were deducted from the available data. The

temperature of the photosphere in the model is TBB = 21,000 K based on the spectral energy dis- tribution evolution of ASASSN-15lh [21]. The velocity of the photosphere vphot. = 16,000 km/s can be calculated from the slope of the photospheric radius evolution. Other parameters like the

profile for the optical depth radial dependence (i.e. power law, Gaussian or exponential) and the

maximum velocity of the envelope (vmax = 50,000 km/s) are weakly constrained by our model. With these results, I proceed to consider the presence of OII a possibility but in small quantities

and have it mostly ionized to OIII. With this possibility, I consider as well the presence of Mg as

it is a product of C and O burning and it has been suggested to be present by other studies of

ASASSN-15lh in the context of TDEs [34]. The same parameters for TBB, vphot. and vmax are used. The result is shown in Figure 4.5.

The parameters used for each elemental species is shown in Table 4.3. The parameters for the

maximum and minimum velocity at which the ions are present were determined with the intention

of having the species after the photosphere and before the end of the photosphere in an arrangement

that resembles a shell. Nonetheless, these parameters are weakly constrained by our model and are

presented here for completion.

In this model, the high optical depth of OIII compared to that of OII is justified due to the

higher ionization of OII compared to normal SNe. The results of the simulation show that OIII

does provide the absorption line near ∼ 4100 A.˚ Furthermore, the low optical depth of OII allows

for the doublet to not be noticeable. In principle, and as it can be seen in the individual contribution

of each species, OII could not be part of the model, which would mean all OII has been ionized.

47 1 QN model (O and Mg) QN model (O only) 0.8 ASASSN-15lh spectrum x u l 0.6 F

d e z i l

a 0.4 m r o

N 0.2

0

3000 4000 5000 6000 7000 8000 Wavelength (Å)

1.75 OIII OII t

n 1.5 MgII a

t ASASSN-15lh spectrum s n o c

+

x 1 u l F

d e z i l

a 0.5 m r o N

0

3000 4000 5000 6000 7000 8000 Wavelength (Å)

Figure 4.5: Synthetic spectrum of ASSASN-15lh using SYNOW at ∼ 26 days post peak luminos- ity. The top panel shows the total simulation of the model using an O compostion and an O and Mg compostion. The bottom panel shows the individual contributions of each elemental species assumed in the model. The data was obtained from Leloudas et al. [34].

48 1 ASASSN-15lh late spectrum QN model (O and Mg) 0.8 x u l F 0.6 d e z i l

a 0.4 m r o

N 0.2

0

3000 4000 5000 6000 7000 8000 Wavelength (Å)

Figure 4.6: Synthetic spectrum of ASSASN-15lh using SYNOW at ∼ 180 days post peak lumi- nosity. The spectrum is practically featureless and it is modelled with a shell of O and Mg. The data was obtained from Leloudas et al. [34].

Furthermore, WO-type SNe are characterized by the presence of absorption lines of high ionization

states of oxygen [14], which is consistent with our model. As for the suggested presence of Mg

by Leloudas et al. [34], this idea cannot be completely discarded. The presence of MgII is a

possibility. However, the stronger features of that species happen below 3000 A˚ and there is no

available data to study whether they were present in ASASSN-15lh or not.

I performed a last test. I simulated the spectrum at a later time. I chose the spectrum at ∼ 180

days because it is mostly featureless and there is an associated TBB = 18,000 K. Figure 4.6 shows the results obtained. The model is the same with the same ion species. Here the optical depths of

OII and MgII remain unchanged as they were already small. The optical depth of OIII was changed to τ(OIII) = 0.5. The velocity of the photosphere was approximated to vphot. = 40,000 km/s and

the weakly constrained species velocities were selected as vmin. = 8,000 km/s and vmax. = 40,000 km/s to try to capture the expansion of the envelope after more than 100 days.

The late spectrum result is mostly shown for completion and to demonstrate the power of the

49 model. A strong consideration is that the late spectrum of ASASSN-15lh is considered to be featureless [21] and that is why no other research group has invested time in modelling it. In principle, a blackbody continuum at 18,000 K should demonstrate to be a good reproduction of the spectrum.

The question then arises, why would OIII not be seen in the late spectrum? And, as the photo- sphere recedes and the core/envelope cools down, why do we not see signatures of recombination or OII? A simple answer is that after the QN, and after more than 100 days, the envelope is largely extended and the species have diluted (i.e. their density has decreased and they have become trans- parent). As the photosphere recedes through the core we see the photons from the cooling core.

These photons do not interact with any other species, as the surrounding medium has become mostly transparent.

4.3 Discussion

ASASSN-15lh challenges all preconceived models of SLNe and stellar evolution. Its unique fea- tures and the fact that it is, so far, the only event with that combination of photometric and spec- troscopic signatures make it the perfect case to test new models for astrophysics.

There is no traditional mechanism that fully explains our photometric and spectroscopic signa- tures of ASASSN-15lh. Some models are capable of addressing the early photometric evolution of the event but fail to explain the late UV rebrightening. Others present a possible explanation for the spectrosopic signatures of the event, but consider extreme events and fail to explain the photometric data.

The first model discussed here, the QN binary model, presents a simple explanation to the light curve of ASASSN-15lh. Its simplicity and the fact that the sequence of events are a natural progression of stellar evolution make it an appealing model. However, a strong deficiency of the model is the lack of explanation for the absorption line present in the early spectrum of ASASSN-

15lh. The inability to capture the photospheric radius evolution and effective temperature evolution

50 signals that this model is, perhaps, incapable of explaining ASASSN-15lh.

The dsQN model, presents a simple explanation to both spectroscopic and photometric signa- tures of ASASSN-15lh. A two-piece density SNR of a WO-type star allows for a double peaked light curve. Furthermore, the presence of oxygen in a WO-SNR, which has been previously docu- mented, allows for an explanation of the spectrum of the event.

The synthetic spectrum simulation presented here, although rudimentary, opens the possibility for the event to be explained by the presence of OIII. This is a new idea that is worth exploring and that has been overseen by other researchers. The results presented here show that OII could be ionized increasing the presence of OIII. Understanding the composition of the ejecta of ASASSN-

15lh through its spectrum, will allow us to understand the conditions of its progenitor star and discriminate whether it is a SLSN or a TDE.

51 Chapter 5

Summary and Conclusion

Our understanding of stellar evolution was challenged with the discovery of SLSNe in general and of double humped SLSNe in particular. These new events puzzled the scientific community because of the lack of a model that can explain the energetics required to power them. With the discovery of these events, new questions arose. What is the mechanism powering SLSNe? How should we categorize them? How do they differ from the traditional model for SNe? Our curiosity was further sparked with the discovery of ASASSN-15lh, the most luminous event to this day.

In this work I presented an overview of SLSNe. Their suggested characterization, classification and their photometric and spectroscopic signatures. Later, I discussed the theoretical model for

QN events, which offer an avenue to explain the high energy required to power SLSNe. Finally,

I applied the QN model to ASASSN-15lh while I discussed its photometric and spectroscopic signatures.

My contributions to the study of ASASSN-15lh were:

• Expansion of the binary model: I applied the QN in a massive star binaries model to ASASSN-

15lh and found that spin-down power would be required to explain the light curve of the event.

• Spectrum simulation: I generated a spectrum simulation for ASASSN-15lh considering a SN

ejecta abundant in oxygen and magnesium.

The QN model for SLSNe offers a creative and simple mechanism to power high energy events in astrophysics. A strength of the model is that it relies on our understanding of nuclear and particle physics. This strength allows us to test our assumptions in laboratories here on Earth, while we look for signatures of these events in space. Quarks exist and quark deconfinement has been observed in laboratories on Earth like the Large Hadron Collider and the Relativistic Heavy

52 Ion Collider. Furthermore, the QN model works as an addendum to our understanding of standard stellar evolution. Our model does not rely on reinventing our understanding of the universe, but allowing for stars to undergo a different path than their death as dwarfs, NSs or BHs.

Of the models presented here, the dsQN model in a two-component density configuration of the

SN ejecta of a WO-type star successfully captures the behaviour of the light curve, photospheric radius evolution and effective temperature evolution of ASASSN-15lh. A QN occurring ∼ 8 days after the SN explosion of a WO-type star reproduces the energetics and photometric signatures of ASASSN-15lh. The first hump in the light curve is powered by the shock resulting from a

QN travelling through a dense core and low-density envelope. As the envelope cools down, the photosphere recedes. Later, the hot dense core is revealed by the receded photosphere powering the second hump in the light curve.

I presented a basic simulation of the synthetic spectrum of ASASSN-15lh. The simulation performed takes most of its parameters from what we know of ASASSN-15lh: its blackbody temperature, the velocity of its photosphere and the idea of a homologously expanding ejecta. The

QN occurring in the center of a WO-type SNR allows us to assume that there are copious amounts of oxygen in it. Hence, I simulated the spectrum with an oxygen abundant SNR reshocked by the QN ejecta and successfully captured the absorption line seen in the early spectra of the event.

Furthermore, the model is promising in the explanation of the late spectrum of ASASSN-15lh.

The work presented here has also some limitations that require further analysis, consideration and discussion:

• The QN hypothesis: The idea that a NS can undergo an explosive transition to a QS has been

deeply studied and analyzed in previous work referred to throughout this thesis. Theoretical and

numerical work suggest that QNe are feasible and that with more simulations and the advent of

new extreme events, such as ASASSN-15lh, it is a matter of time before we conclusively know

of their existence. Quark matter is a reality and nuclear science tells us that QSs should exist.

Work like this and the extensive work done by the Quark-Nova group referred here serve as

53 encouragement for other scientists to think of this possibility.

• The two-component configuration: The assumed two-density profile remains to be confirmed.

This profile requires further work with hydrodynamical simulations that properly capture the

transition between an envelope and a core left behind such an event. If proved correct, this

model tells us that the SN in ASSASN-15lh leads to a unique remnant that is not very common

in our universe.

• The number of parameters in the model: The nine parameters in our model might seem to

be enough to fit any type of event. However, these parameters are constrained by the known

characteristics of WO-stars. None of the parameters suggest extreme behaviours.

• The use of SYNOW: Although it has been used to study SLSNe (see Leloudas et al. [33])

SYNOW is not designed to account for the possible differences between traditional SNe and

SLSNe. The development of a new tool that is capable of simulating the spectrum of this new

class of events would be advantageous. The work here is done under the assumption that the

behaviour of the envelope in our model is relatively similar to the assumptions in SYNOW. The

main contribution of our model is offering an explanation for the possible chemical composition

of the SNR and the reason why OIII is present. More detailed spectrum synthesis codes are

planned to be used in future work to offer a greater confidence in our results and the identifica-

tion of ASASSN-15lh’s spectroscopic features. If correct, our model could be applied to other

progenitor stars that could have significantly different absorption lines in their spectrum.

This work comes at a relevant time when the mechanism behind ASASN-15lh remains un- known. The implications of this work go beyond understanding ASASSN-15lh. The study of such energetic events allows us to understand the fate of massive stars as they evolve, the possibility of quark matter and QSs, along with their astronomical signatures and the processes through which matter interacts. It is a bridge between astrophysics and quantum chromodynamics.

54 In conclusion, the dsQN model not only a compelling mechanism for SLSNe in general and

ASASSN-15lh in particular, but if true, we might be at the verge of demonstrating that the most energetic events in the universe are fundamentally linked to the formation of the tiniest particles namely, quarks.

55 Bibliography

[1] J. D. Anand, A. Goyal, V. K. Gupta, and S. Singh. Burning of two-flavor quark matter into

strange matter in neutron stars and in supernova cores. ApJ, 481(2):954, 1997.

[2] K. Barbary, K. S. Dawson, K. Tokita, et al. Discovery of an unusual optical transient with

the Hubble Space Telescope. ApJ, 690(2):1358, 2008.

[3] M. C. Bersten, O. G. Benvenuto, M. Orellana, and K. Nomoto. The unusual super-luminous

supernovae SN 2011kl and ASASSN-15lh. ApJL, 817(1):L8, 2016.

[4] A. Bhattacharyya, S. K. Ghosh, P. S. Joarder, R. Mallick, and S. Raha. Conversion of a

neutron star to a strange star: A two-step process. PhRvC, 74(6):065804, 2006.

[5] A. R. Bodmer. Collapsed nuclei. PhRvD, 4:1601–1606, 1971.

[6] J. Boguta and A. R. Bodmer. Relativistic calculation of nuclear matter and the nuclear sur-

face. NuPhA, 292(3):413–428, 1977.

[7] R. L. Bowers and T. Deeming. Astrophysics: Stars, volume 1. Jones and Bartlett, 1984.

[8] D. Branch, J. Parrent, M.A. Troxel, et al. Probing the nature of type I supernovae with

SYNOW. In AIPConf, volume 924, pages 342–349. AIP, 2007.

[9] N. Cabibbo and G. Parisi. Exponential hadronic spectrum and quark liberation. PhLB,

59(1):67–69, 1975.

[10] E. Chatzopoulos, D. R. Van Rossum, W. J. Craig, et al. Emission from Pair-instability Super-

novae with Rotation. ApJ, 799(1):18, 2015.

[11] E. Chatzopoulos, J. C. Wheeler, and J. Vinko. Generalized semi-analytical models of super-

nova light curves. ApJ, 746(2):121, 2012.

56 [12] E. Chatzopoulos, J. C. Wheeler, J. Vinko, et al. Extreme supernova models for the super-

luminous transient asassn-15lh. ApJ, 828(2):94, 2016.

[13] E. Chatzopoulos, J. C. Wheeler, J. Vinko, Z. L. Horvath, and A. Nagy. Analytical Light Curve

Models of Superluminous Supernovae: χ2-minimization of Parameter Fits. ApJ, 773(1):76,

2013.

[14] C. de Loore and A. J. Willis. Wolf-Rayet stars: observations, physics, evolution, volume 99.

Springer Science & Business Media, 2012.

[15] J. Dexter and D. Kasen. Supernova Light Curves Powered by Fallback Accretion. ApJ,

772(1):30, 2013.

[16] S. Dong, B. J. Shappee, J. L. Prieto, et al. ASASSN-15lh: A highly super-luminous super-

nova. Sci, 351(6270):257–260, 2016.

[17] A.C. Doyle and C. Morley. The Complete Sherlock Holmes. Doubleday, 1930.

[18] A. K. Fisher. Direct analysis of Type Ia supernovae spectra. PhD thesis, University of

Oklahoma, 2000.

[19] A. Gal-Yam. Luminous Supernovae. Sci, 337(6097):927–932, 2012.

[20] A. Gal-Yam, P. Mazzali, E. O. Ofek, et al. Supernova 2007bi as a pair-instability explosion.

Natur, 462(7273):624–627, 2009.

[21] D. Godoy-Rivera, K. Z. Stanek, C. S. Kochanek, et al. The Unexpected, Long-Lasting,

UV Rebrightening of the Super-Luminous Supernova ASASSN-15lh. arXiv preprint

arXiv:1605.00645, 2016.

[22] W. R Hamann. Line formation in expanding atmospheres-On the validity of the Sobolev

approximation. A&A, 93:353–361, 1981.

57 [23] A. Heger and S. E. Woosley. The nucleosynthetic signature of population III. ApJ,

567(1):532, 2002.

[24] H. Heiselberg, G. Baym, and C. J. Pethick. Burning of strange quark matter and transport

properties of QCD plasmas. NuPhS, 24(2):144–147, 1991.

[25] M. Herzog and F. K. Ropke.¨ Three-dimensional hydrodynamic simulations of the combustion

of a neutron star into a quark star. PhRvD, 84(8):083002, 2011.

[26] T. W.-S Holoien, J. L. Prieto, D. Bersier, et al. ASASSN-14ae: a tidal disruption event at 200

Mpc. MNRAS, 445(3):3263–3277, 2014.

[27] D. Kasen and L. Bildsten. Supernova light curves powered by young magnetars. ApJ,

717(1):245, 2010.

[28] P. Keranen,¨ R. Ouyed, and P. Jaikumar. Neutrino emission and mass ejection in quark novae.

ApJ, 618(1):485, 2005.

[29] M. Kostka, N. Koning, D. Leahy, R. Ouyed, and W. Steffen. Quark nova signatures in super-

luminous supernovae. RMxAA, 50(2):167–180, 2014.

[30] A. Kozyreva, R. Hirschi, S. Blinnikov, and J. den Hartogh. How much radioactive nickel

does ASASSN-15lh require? MNRASL, 459(1):L21–L25, 2016.

[31] L. D. Landau. On the theory of stars. Physikalische Zeitschrift der Sowjetunion, 1(285):152,

1932.

[32] D. Leahy and R. Ouyed. Supernova SN2006gy as a first ever Quark Nova? MNRAS,

387(3):1193, 2008.

[33] G. Leloudas, E. Chatzopoulos, B. Dilday, et al. SN 2006oz: rise of a super-luminous super-

nova observed by the SDSS-II SN Survey. A&A, 541:A129, 2012.

58 [34] G. Leloudas, M. Fraser, N. C. Stone, et al. Corrigendum: The superluminous transient

ASASSN-15lh as a tidal disruption event from a Kerr black hole. NatAs, 1, 2016.

[35] G. Lugones, O. G. Benvenuto, and H. Vucetich. Combustion of nuclear matter into strange

matter. PhRvD, 50(10):6100, 1994.

[36] M. A. A. Manrique and G. Lugones. Hydrodynamic Simulations of the Combustion of Dense

Hadronic Matter into Quark Matter. BrJPh, 45(4):457–466, 2015.

[37] S. Mereghetti. The strongest cosmic magnets: soft gamma-ray repeaters and anomalous X-

ray pulsars. A&ARv, 15(4):225–287, 2008.

[38] B. D. Metzger, B. Margalit, D. Kasen, and E. Quataert. The diversity of transients from

magnetar birth in core collapse supernovae. MNRAS, 454(3):3311–3316, 2015.

[39] T. Moriya, N. Tominaga, M. Tanaka, K. Maeda, and K. Nomoto. A core-collapse supernova

model for the extremely luminous Type Ic Supernova 2007bi: an alternative to the pair-

instability supernova model. ApJL, 717(2):L83, 2010.

[40] T. J. Moriya, S. I. Blinnikov, N. Tominaga, et al. Light-curve modelling of superluminous

supernova 2006gy: collision between supernova ejecta and a dense circumstellar medium.

MNRAS, 428(2):1020–1035, 2013.

[41] T. J. Moriya, Z. W. Liu, J. Mackey, T. W. Chen, and N. Langer. Revealing the binary origin of

Type Ic superluminous supernovae through nebular hydrogen emission. A&A, 584:L5, 2015.

[42] T. J. Moriya and N. Tominaga. Diversity of Luminous Supernovae from Non-steady Mass

Loss. ApJ, 747(2):118, 2012.

[43] B. Niebergal. Hadronic-to-Quark-Matter Phase Transition: Astrophysical Implications.

PhD thesis, University of Calgary, 2011.

59 [44] B. Niebergal, R. Ouyed, and P. Jaikumar. Numerical simulation of the hydrodynamical com-

bustion to strange quark matter. PhRvC, 82:062801, 2010.

[45] M. L. Olesen and J. Madsen. Burning a neutron star into a strange star. NuPhS, 24(2):170–

174, 1991.

[46] A. V. Olinto. On the conversion of neutron stars into strange stars. PhLB, 192(1-2):71–75,

1987.

[47] A. Ouyed. Spallation in Dual-Shock Quark Nova: A Robust Nucleosynthetic Process. Mas-

ter’s thesis, University of Calgary, 2013.

[48] A. Ouyed, L. Welbanks, N. Koning, and R. Ouyed. The Burn-UD code for the numerical

simulations of the Hadronic-to-Quark-Matter phase transition. Conference Proceedings of

the Compact Stars in the QCD Phase Diagram IV (CSQCD IV), 2015.

[49] R. Ouyed, J. Dey, and M. Dey. Quark-nova. A&A, 390(3):L39–L42, 2002.

[50] R. Ouyed, N. Koning, and D. Leahy. SN 2009ip and SN 2010mc as dual-shock Quark-Novae.

RAA, 13(12):1463, 2013.

[51] R. Ouyed, M. Kostka, N. Koning, D. Leahy, and W. Steffen. Quark nova imprint in the

extreme supernova explosion SN 2006gy. MNRAS, 423(2):1652–1662, 2012.

[52] R. Ouyed, D. Leahy, and N. Koning. Quark-Novae in massive binaries: a model for double-

humped, hydrogen-poor, superluminous Supernovae. MNRAS, 454(3):2353–2359, 2015.

[53] R. Ouyed, D. Leahy, and N. Koning. The Superluminous SN DES13S2cmm as a Signature

of a Quark-nova in a He-HMXB System. ApJ, 809(2):142, 2015.

[54] R. Ouyed, D. Leahy, and N. Koning. Quark-novae Occurring in Massive Binaries: A Uni-

versal Energy Source in Superluminous Supernovae with Double-peaked Light Curves. ApJ,

818(1):77, 2016.

60 [55] R. Ouyed, S. Leahy, L. Welbanks, and N. Koning. The Superluminous (Type I) Supernova

ASASSN-15lh: A case for a Quark-Nova inside an Oxygen-type Wolf-Rayet supernova rem-

nant. arXiv preprint arXiv:1611.03657, 2016.

[56] R. Ouyed, B. Niebergal, and P. Jaikumar. Explosive Combustion of a Neutron Star into a

Quark Star: the non-premixed scenario. in Conf. Proceedings of the Compact Stars in the

QCD Phase Diagram III (CSQCD III), ed. L. Paulucci, J. E. Horvath, M. Chiapparini, & R.

Negreiros (Guaruja, Brazil), 12-15, 2013.

[57] R. Ouyed, R. Rapp, and C. Vogt. Fireballs from quark stars in the color-flavor locked phase:

application to gamma-ray bursts. ApJ, 632(2):1001, 2005.

[58] M. A.´ Perez-Garc´ ´ıa, F. Daigne, and J. Silk. Short gamma-ray bursts and dark matter seeding

in neutron stars. ApJ, 768(2):145, 2013.

[59] D. Prialnik. An Introduction to the Theory of Stellar Structure and Evolution. Cambridge

University Press, 2009.

[60] R. M. Quimby, G. Aldering, J. C. Wheeler, et al. SN 2005ap: A most brilliant explosion.

ApJL, 668(2):L99, 2007.

[61] G. Rakavy and G. Shaviv. Instabilities in highly evolved stellar models. ApJ, 148:803, 1967.

[62] A. Rest, R. J. Foley, S. Gezari, et al. Pushing the Boundaries of Conventional Core-collapse

Supernovae: The Extremely Energetic Supernova SN 2003ma. ApJ, 729(2):88, 2010.

[63] S. Rosswog and M. Bruggen.¨ Introduction to High-Energy Astrophysics. Cambridge Univer-

sity Press, 2011.

[64] N. Smith, W. Li, R. J. Foley, et al. SN 2006gy: discovery of the most luminous super-

nova ever recorded, powered by the death of an extremely massive star like η carinae. ApJ,

666(2):1116, 2007.

61 [65] T. Sukhbold and S. E. Woosley. The most luminous supernovae. ApJL, 820(2):L38, 2016.

[66] R. C. Thomas. SYN++: Standalone SN spectrum synthesis. ascl soft, 2013.

[67] R. C. Thomas, P. E. Nugent, and J. C. Meza. SYNAPPS: data-driven analysis for supernova

spectroscopy. PASP, 123(900):237, 2011.

[68] C. Thompson and R. C. Duncan. Neutron star dynamos and the origins of pulsar magnetism.

ApJ, 408:194–217, 1993.

[69] F. Tramper, G. Grafener,¨ O. E. Hartoog, et al. On the nature of WO stars: a quantitative

analysis of the WO3 star DR1 in IC 1613. A&A, 559:A72, 2013.

[70] S. Van Velzen, G. R. Farrar, S. Gezari, et al. Optical discovery of probable stellar tidal

disruption flares. ApJ, 741(2):73, 2011.

[71] C. Vogt, R. Rapp, and R. Ouyed. Photon emission from dense quark matter. NuPhA, 735(3-

4):543–562, 2004.

[72] L. J. Wang, S. Q. Wang, Z. G. Dai, et al. Optical Transients Powered by Magnetars: Dynam-

ics, Light Curves, and Transition to the Nebular Phase. ApJ, 821(1):22, 2016.

[73] S. Weissenborn, I. Sagert, G. Pagliara, M. Hempel, and J. Schaffner-Bielich. Quark matter in

massive compact stars. ApJ, 740(1):L14, 2011.

[74] L. Welbanks, A. Ouyed, N. Koning, and R. Ouyed. Simulating Hadronic-to-Quark-Matter

with Burn-UD: Recent work and astrophysical applications. in Conf. Proceedings of the

Compact Stars in the QCD Phase Diagram V (CSQCD V), ed. M. Mannarelli & I. Bombaci

(Gran Sasso, Italy), 2017.

[75] E. Witten. Cosmic separation of phases. PhRvD, 30:272–285, 1984.

[76] S. E. Woosley. Bright supernovae from magnetar birth. ApJL, 719(2):L204, 2010.

62 [77] S. E. Woosley, S. Blinnikov, and A. Heger. Pulsational pair instability as an explanation for

the most luminous supernovae. Natur, 450(7168):390–392, 2007.

[78] D. G. Yakovlev, P. Haensel, G. Baym, and C. Pethick. Lev Landau and the concept of neutron

stars. PhyU, 56(3):289, 2013.

63