From Binaries to Asymmetric Outflows: The Influence of Low-mass Companions Around AGB Stars
by
Jason T. Nordhaus
Submitted in Partial Fulfillment of the Requirements for the Degree Doctor of Philosophy
Supervised by Professor Eric G. Blackman Department of Physics and Astronomy The College Arts and Sciences
University of Rochester Rochester, New York
2008 ii
To my mom and dad iii
Curriculum Vitae
The author was born November 24, 1980 in Concord, Massachusetts. He graduated from the University of Rochester with a Bachelor of Science degree in Physics and Astronomy and a Bachelor of Arts degree in Mathematics in 2003. Upon completion of his undergraduate education, he entered the doctoral program in the Department of Physics and Astronomy at the University of Rochester. In May 2004 the author received a Master of Arts degree in Physics from the University of Rochester. The author received a Department of Energy Frank J. Horton Fellowship (2004-2008) in addition to a Department of Education GAANN fellowship (2003-2006).
Selected Publications
• Nordhaus, J., Minchev, I., Sargent, B., Forrest, W., Blackman, E. G., De Marco, O., Kastner, J., Balick, B., Frank, A. 2008 MNRAS, submitted
• Edgar, R. G., Nordhaus, J., Blackman, E., Frank, A. 2008 ApJL, in press
• Nordhaus, J., Blackman, E. G., Frank, A. 2007 MNRAS, 376, 599
• Minchev, I., Nordhaus, J., Quillen, A. 2007 ApJL, 664, 31
• Watson, D. M., Leisenring, J. M., Furlan, E., Bohac, C. J., Sargent, B., Forrest, W. J., Calvet, N., Hartmann, L., Nordhaus, J. T., Green, J. D., Kim, K. H., Sloan, G. C., Chen, C. H., Keller, L. D., d’Alessio, P., Najita, J., Uchida, K. I., Houck, J. R. 2007, ApJS, submitted
• Nordhaus, J., Blackman, E. G. 2006, MNRAS370, 2004
• Blackman, E. G., Nordhaus, J., Thomas, J. H. 2006 New Astronomy, 11, 452 iv
• Nordhaus, J., Blackman, E. G. 2008 to appear in AIP Proceedings of the IXth Torino Workshop on AGB Nucleosyntheis
• Nordhaus, J., Blackman, E. G. 2007, in Asymmetric Planetary Nebulae IV, eds. R. L. M. Corradi, A. Manchado, N. Soker (in ASP Conference Series: San Francisco), in press
• Blackman, E. G., Nordhaus, J. 2007, in Asymmetric Planetary Nebulae IV, eds. R. L. M. Corradi, A. Manchado, N. Soker (in ASP Conference Series: San Francisco), in press v
Acknowledgments
This work would not have been possible without the support and encouragement of my thesis advisor, Dr. Eric Blackman. His guidance, patience and scientific insight were instrumental throughout my graduate career. I also wish to thank my fellow graduate compatriots, Dave Clader and Ivan Minchev for invaluable friendship, stim- ulating scientific & non-scientific discussions and for an endless supply of much needed distractions. I wish to acknowledge institutions which provided financial support and assis- tance. In particular, the Department of Physics and Astronomy at the University of Rochester. Financial support for this work was provided by the Laboratory for Laser Energetics through a U.S. Department of Energy Horton Fellowship and a U.S. Department of Education GAANN Fellowship. I would also like to thank my parents, Kurt and Sherri, and sisters, Miranda and Tiffany, for their love and support. Finally, I wish to thank my wife Lea for her unconditional love and encouragement. vi
Abstract
The study of intermediate mass, evolved stars is undergoing renewed interest due to recent observational and theoretical results suggesting that binarity is fundamental for shaping post-Asymptotic Giant Branch and Planetary Nebula outflows. Despite ex- tensive research, the physical mechanism responsible for transitioning from a spherical Asymptotic Giant Branch (AGB) star to an asymmetric post-AGB object is poorly understood. In an effort to understand how binaries may produce asymmetries, this thesis presents several theoretical studies which explore the effect of low-mass com- panions on evolved star outflows. This thesis consists of four separate projects: (1.) Close companions may become engulfed by the evolved star and in-spiral during a common envelope phase. Common envelope evolution can lead to three different consequences: (i.) equatorial ejection of material (ii.) spin-up of the envelope resulting in an explosive dynamo-driven jet and (iii.) tidal shredding of the companion into an accretion disk which ejects a poloidal wind. (2.) In addition, we study a dynamical, large-scale α − Ω interface dynamo oper- ating in an AGB star in both an isolated setting and a setting in which a low-mass companion is embedded inside the envelope. The back reaction of the fields on the shear is included and differential rotation and rotation deplete via turbulent dissipa- tion and Poynting flux. For the isolated star, the shear must be resupplied in order to sufficiently sustain the dynamo. Furthermore, we investigate the energy requirements that convection must satisfy to accomplish this by analogy to the Sun. For the com- mon envelope case, a robust dynamo results, unbinding the envelope under a range of conditions. (3.) Wide binaries can interact with the wind of the evolved primary. The grav- itational influence of the secondary focuses material in the equatorial plane. The vii companion induces spiral shocks which may anneal amorphous grains into crystalline dust. This work presents a physical mechanism to produce crystalline dust in AGB star binaries. (4.) We present a spectral modeling technique which constrains the geometry of evolved star nebulae. We apply our technique to HD 179821 which exhibits a dou- ble peaked spectral energy distribution (SED) with a sharp rise from ∼ 8 − 20 µm. Such features have been associated with dust shells or inwardly truncated circumstel- lar disks. In order to compare SEDs from both systems, we employ a spherically sym- metric radiative transfer code and compare it to a radiative, inwardly truncated disc code. As a case study, we model the broad-band SED of HD 179821 using both codes. Shortward of 40 µm, we find that both models produce equivalent fits to the data. However, longward of 40 µm, the radial density distribution and corresponding broad range of disc temperatures produce excess emission above our spherically symmetric solutions and the observations. For HD 179821, our best fit consists of a Teff = 7000
K central source characterized by τV ∼ 1.95 and surrounded by a radiatively driven, spherically symmetric dust shell. The extinction of the central source reddens the broad-band colours so that they resemble a Teff = 5750 K photosphere. We believe that HD 179821 contains a hotter central star than previously thought. Our results provide an initial step towards a technique to distinguish geometric differences from spectral modeling. viii
Contents
CurriculumVitae...... iii Acknowledgments...... v Abstract...... vi
Introduction 1 1.1Post-MainSequenceEvolution...... 1 1.1.1 AsymptoticGiantBranchEvolution...... 1 1.1.2 Post-AGBPhase...... 2 1.2 Support for the Binary Hypothesis ...... 4 1.2.1 ObservationalIndications...... 4 1.2.2 MagneticShaping...... 5 1.3ThesisStructure...... 5 References...... 7
Low-mass Binary Induced Outflows from Asymptotic Giant Branch Stars 9 2.1Abstract...... 9 2.2Introduction...... 10 2.3CommonEnvelopeEvolution...... 11 2.3.1 EnvelopeBindingEnergy...... 12 2.3.2 Orbital Energy and Angular Momentum Evolution ...... 14 2.4CommonEnvelopeEvolutionScenarios...... 17 2.4.1 SecondaryInducedEnvelopeExpulsion...... 17 2.4.2 Secondary Induced Envelope α − ΩDynamo...... 20 2.4.3 DiscDrivenOutflow...... 24 CONTENTS ix
2.5DiscussionofObservationalImplications...... 25 2.5.1 ObservationalConsequences...... 26 2.5.2 ApplicationstospecificPPNeandPNesystems...... 27 2.6Conclusions...... 29 References...... 31
Isolated versus Common Envelope Dynamos in Planetary Nebula Pro- genitors 35 3.1Abstract...... 35 3.2Introduction...... 36 3.3 Dynamos, Common Envelopes and Isolated AGB Evolution ...... 37 3.4DynamicalEquations...... 39 3.4.1 EvolutionofΩand∆Ω...... 41 3.4.2 Evolution of α ...... 43 3.5NumericalResults...... 44 3.5.1 IsolatedDynamoWithoutReseeding∆Ω...... 45 3.5.2 IsolatedDynamoWithReseeding∆Ω...... 47 3.5.3 CommonEnvelopeDynamo...... 49 3.6MorphologyofMagneticOutflows...... 54 3.7Conclusions...... 56 References...... 58
Towards a Spectral Technique for Determining Material Geometry Around Evolved Stars: Application to HD 179821 61 4.1Abstract...... 61 4.2Introduction...... 62 4.3 HD 179821: post-AGB or red supergiant? ...... 64 4.4PhotosphericModelsandExtinction...... 66 4.5InnerWall,Edge-onDiskModels...... 68 4.6 Spherical Shell Models ...... 72 4.7Summary...... 77 CONTENTS x
The Formation of Crystalline Dust in AGB Winds from Binary Induced Spiral Shocks 83 5.1Abstract...... 83 5.2Introduction...... 83 5.3NumericalStudy...... 85 5.4Results...... 85 5.5Discussion...... 87 5.5.1 GrainAnnealing...... 87 5.5.2 ShockTemperatureScaling...... 90 5.5.3 DustFormation...... 91 5.6Conclusion...... 91 xi
List of Tables
4.1PhotometricData...... 68 4.2 Color Corrected IRAS fluxes, Submillimeter Data for HD 179821 . . . 69 4.3ModelSummary...... 76 xii
List of Figures
1.1 Left (IRC +10216; V bank): Typically spherical AGB mass loss is revealed by spherical reflection nebulosities. Center (CRL 2688; [OI]): every time a reflection or neutral nebula is seen around a post-AGB star, a bipolar symmetry is present. Right (He 2-47; Hα): when post-AGB stars heat up, the young ionized PNe always have bipolar or multi-polar morphologies. [Credits: Mauron & Huggins (2000); Sahai et al. (1998); Sahai & Trauger (1998)] ...... 3
2.1 Left: Density and mass profiles for our model AGB star. The dotted line is the core-envelope boundary. Right: Mach number and sound speed as a function of radius. The Mach number is computed from the Keplerian motion of the planet inside the envelope. The motion is supersonic everywhere and thus justifies our choice of accretion radius (Bondi 1952)...... 13 2.2 Infall time as a function of position inside the envelope of the AGB star (left) and interpulse AGB star (right). The solid line represents a
companion of mass 0.02 M and the dotted line is a secondary of mass
0.2 M ...... 16 LIST OF FIGURES xiii
2.3 Three possible outcomes of our CE evolution. (a.) The companion be- comes embedded in the stellar envelope, orbital separation is reduced, eventually resulting in unbinding the envelope equatorially. (b.) The companion spirals in, the envelope is spun up causing it to differen- tially rotate. The presence of a deep convective zone, coupled with the differential rotation, generates a dynamo in the envelope. (c.) The companion is shredded into an accretion disc around the core. The disc then drives an outflow which, in principle, can unbind the envelope. . 18 2.4 For various efficiencies α (see Eq. 1), the solid line shows the radius at which the change in orbital energy equals the binding energy of the envelope for the beginning of the AGB star (left) and interpulse AGB star (right). The dotted vertical line marks the core-envelope boundary. The long-dashed line represents the radius at which the companion is tidally shredded by the core. The short-dashed line is where the com- panion first fills its Roche lobe, initiating mass transfer to the envelope. 19 2.5 The solid line depicts the energy required to unbind the envelope for the AGB star (left) and interpulse AGB star (right), if the secondary is not tidally shredded as it traverses the envelope. The dashed lines represent the amount of energy deposited into the envelope from the change in orbital energy of the secondary for efficiency parameter α
(Eq. 1). For α =1.0, a m2 =0.02 M brown dwarf delivers enough energy to blow off the AGB envelope at r ∼ 1010 cm. For α =0.3, the brown dwarf must traverse all the way to the core-envelope boundary before supplying enough energy to unbind the system. For smaller α,
a m2 =0.02 M companion cannot unbind the AGB envelope before spiraling down to a radius where an interface dynamo might participate
in unbinding the envelope. For the interpulse AGB star, a 0.02 M brown dwarf can supply enough orbital energy to unbind the envelope for α =1.0andα =0.3...... 21 LIST OF FIGURES xiv
2.6 Two rotation profiles for our 3.0 M AGB star. The solid curve repre- sents the spin up of an initial stationary envelope by an infalling 0.02
M brown dwarf. The dotted curve is the rotation profile generated in Blackman et al. 2001 in which a main sequence star exhibiting solid body rotation conserves angular momentum of spherical mass shells during its evolution onto the AGB. The solid vertical line marks the core boundary and the short-dashed line represents the base of the con- vective zone. The long-dashed line is the base of the differential rotation zone used in Blackman et al. 2001...... 23
3.1 A meridional slice of the dynamo geometry. The left figure shows the global geometry of the AGB star. The right figure is a close-up view of the dashed region on the left. The α-effect is driven by convection
and occurs in layer of thickness L1 above the differential rotation zone. The poloidal component of the field is pumped downwards into the differential zone, where it is wrapped torodially due to the Ω-effect. . . 39 3.2 The differential rotation energy is allowed to drain through field ampli- fication and turbulent dissipation. In this figure, k =5× 10−11 cm−1, z t cp =0.01, Q =5.0. We define, [PF,dis] ≡ 0 E[PF,dis] (t ) dt and label
M (PF)andT (dis) on the top right plot to distinguish between the thermal and magnetic contributions to the binding energy. For the left −4 −5 figure, cφ =10 while the right has cφ =10 . Peak field strengths are a factor of ∼ 5 − 10 less then those obtained in (Blackman et al. 2001). Differential rotation energy is drained in < 20 yrs. Lowering
cφ results in the differential rotation energy draining at a slower rate, allowing the field to sustain for longer periods of time (∼ 40 − 50 yrs). However,peakfieldstrengthsremainthesame...... 46 3.3 Results for reseeding differential rotation through convection. In the left figure, f = 1 corresponding to maximum convective resupply. Rotation is drained through Poynting flux but cannot sustain a dynamically im- portant dynamo. In the figure on the right, f = 0 (no resupply of ∆Ω). The rotation rate is fixed, corresponding to a buildup of Poynting flux intheinterfacelayer...... 48 LIST OF FIGURES xv
3.4 Convective resupply results in a steady-state differential rotation profile. For the left column, the envelope of the poloidal, toroidal and Poynting flux is plotted. The Poynting flux is sustained at ∼ 5 × 1034 erg/s. The sustained Poynting flux supplies enough energy to unbind the envelope 5 of our 3 M model at the end of the AGB phase (∼ 10 yrs). In this −5 −3 figure, cφ =10 and f =10 implying that only ∼ 0.1% of the cascade energy must be converted into differential rotation energy to supply the requisite Poynting flux. This model predicts a magnetically dominatedexplosion...... 50 3.5 Rotation profiles generated from the transfer of angular momentum from companion to envelope in our AGB star. In both figures, the solid curves represent the resulting profiles for companions of masses
0.05 (top), 0.02 and 0.01 (bottom) M . For the left figure, the dashed curve represents the Keplerian velocity while the dashed-dotted curve
is the sound speed. The 0.05 M companion initially spins up the envelope such that the inner region is rotating faster then the Keplerian velocity. Mass redistribution ensues and transfers matter outward until the rotation profile drops below Keplerian. The right figure presents the angular velocity corresponding to the left figure. The dash-dot vertical line is the approximate radius at which the companion is tidally shredded. The large-dash vertical line is the boundary of the shear layer in Blackman et al. 2001 while the small dash line is the base of the
convection zone. These profiles assume that αCE =0.3...... 52 LIST OF FIGURES xvi
3.6 Interface dynamo resulting from the in-spiral of a 0.02 M brown dwarf in the interior of our model AGB star. The differential rotation zone extends from the base of the convection zone to the radius at which the secondary is tidally shredded (Nordhaus & Blackman 2006). In this −4 −4 −3 model, PM =10 and Q =5,Ω0 =2.3×10 rad/s, ∆Ω0 =2.5×10 rad/s and δ/L = 1. In the left column, the envelope of the Poynting flux (top), toroidal field (middle) and poloidal field (bottom) are drawn with a solid line. The insets represent the time evolution from 0 to 0.2 yrs. The vertical scale of the insets are the same as the corresponding largerfigure...... 53
3.7 Interface dynamo resulting from the in-spiral of a 0.05 M brown dwarf. −6 −4 In this model PM =10 , Q =5,δ/L =1,Ω0 =5× 10 rad/s and −3 ∆Ω0 =2.5 × 10 rad/s. The insets represent the time evolution from 0 to 0.2 years. The vertical scale of the insets are the same as the correspondinglargerfigure...... 55
4.1 Top: ISO SWS spectra of HD 179821 and post-AGB object HD 161796. In both objects, there is a steep rise between the 10 and 20 µmfeatures possibly indicating a transition region to the optically thick outer wall
or shell. HD 179821 was de-reddened using AV = 2 while HD 161796
was corrected using AV =1.2. Bottom: ISO spectrum of the post-RSG IRC +10420...... 67 4.2 Our best fit wall models. The top figure corresponds to µ =0.25 while the bottom corresponds to µ =0.45...... 71 4.3 Another fit to the wall model for glassy olivine (top) and glassy bronzite (bottom). The inclination in both figures is µ =0.25...... 73 4.4 Our spherical shell model (dark line). The dust temperature at the inner and outer radii are 128 K and 41 K. Even though the central radiation source is a T = 7000 K photosphere, the detached shell reddens the central source so that it resembles a T = 5750 K photosphere (thin line). The silicate features are fit using a mixture of glassy, amorphous silicates with a small component of FeO...... 75 LIST OF FIGURES xvii
5.1 Volumetric density renderings of the wind emitted by a 1 M primary,
with a 0.25 M secondary in a 6 AU orbit. The view along the z axis is shown on the left, that along the x axis on the right. The binary orbits in the z =0plane...... 86
5.2 Temperature structure of the wind emitted by a 1 M primary, with a
0.25 M secondaryina6AUorbit...... 88 1
Chapter 1
Introduction
The study of intermediate mass, evolved stars is an important topic of active research due to recent observational and theoretical results suggesting that binarity is funda- mental for shaping post-Asymptotic Giant Branch and Planetary Nebula outflows. In this thesis, we present several theoretical studies which explore the effect of low-mass companions on an evolved star. In particular, close companions may be engulfed dur- ing post-main sequence evolution and eject material toroidally or via collimated bipolar outflows. Wider binaries can gravitationally focus evolved star winds and drive spiral shocks in equatorial outflows. In addition to the physics of binary interactions, we also present a spectral modeling technique that can be applied to unresolved, evolved stars to constrain the geometry of the circumstellar nebulae.
1.1 Post-Main Sequence Evolution
1.1.1 Asymptotic Giant Branch Evolution
After exhausting core hydrogen and helium burning, intermediate-mass stars (< 8 M ) enter the AGB phase. The structure of an AGB star consists of an electron degenerate C-O core surrounded by a He shell underneath a H-rich envelope. The star is luminous 3 4 3 (∼ 10 −10 L ) but cool (Teff < 3×10 K) and has expanded to ∼ 200 times its main- sequence radius. Initially, a quiet period in which nuclear burning is limited to the He shell defines the early-AGB phase. Hydrogen burning above the He-shell eventually dominates and marks the transition to the thermal pulsing AGB phase. As the H- CHAPTER 1. INTRODUCTION 2 shell burns, the mass of the underlying He-shell increases. Eventually, the He-shell reignites in a thermo-nuclear runaway and is commonly referred to as a thermal pulse (TP; Schwarzschild Harm 1965). The sudden increase in energy release extinguishes H-burning during which the convective zone can penetrate into the interior regions and dredge up C-rich material. As a result of the dredge-ups incurred by thermal pulses, the surface of the star is enriched with nuclear burning products. In particular, the initially oxygen-rich (C/O < 1) AGB star can become carbon-rich (C/O > 1) if it undergoes 3rd dredge-up (Iben 1975). If AGB evolution were strictly determined by nuclear burning processes, then the chemistry of the remnant nebula would depend only on the nuclear processing of the progenitor star. Unfortunately, mass-loss timescales exceed nuclear burning timescales by up to four orders of magnitude in the AGB phase. Thus, AGB termination results from strong mass-loss experienced during this phase. The mass-loss process is poorly understood and likely dependent on dust formation from the thermal pulse process (see Sedlmayr & Dominik 1995 for a review). If however, the AGB star is part of a binary, the picture is less clear and the evolution complicated. Average mass-loss rates −5 in isolated stars during the AGB phase are typically M˙ ∼10 M /yr, depleting the stellar envelope in ∼ 104 − 105 yrs. Typical AGB wind velocities are slow (∼ 10 − 15 km/s) and spherical. The depletion of the envelope eventually results in a detached nebula and signifies the end of the AGB star and the beginning of the post-AGB phase.
1.1.2 Post-AGB Phase
The post-AGB phase is short (∼ 103 − 104 yrs) and poorly understood. The effective temperature is increasing while the luminosity remains fairly constant. Once the central white dwarf reaches a temperature of ∼ 104 K, the circumstellar nebula (if present) ionizes and shines as PN. For a review on post-AGB stars see (van Winckel 2003). Originally, the ejected nebulae were thought to be spherical and originate from quasi-steady mass loss by an isolated star during the AGB phase. As the core heated, the surrounding ionized nebula would reveal spherical symmetry. This view has changed since high resolution optical telescopes (particularly the Hubble Space Tele- scope) revealed highly asymmetric structures consisting of disks, bipolar outflows and CHAPTER 1. INTRODUCTION 3 bullet-like ejecta (for a review see Balick and Frank 2002). The origin of the asym- metry and the physical mechanisms responsible for shaping have remained elusive for nearly three decades. Until recently, the textbook explanation of PNe shaping involved an interacting spherically symmetric “fast” wind preceded by an equatorial “slow” wind (Kwok et al. 1978). As the fast wind encounters the slow wind, it becomes collimated in the polar direction. Unfortunately, this scenario is only able to explain a subset of the older PNe and categorically fails to explain the younger, post-AGB progenitors and young PNe (Sahai 2002).
Figure 1.1 Left (IRC +10216; V bank): Typically spherical AGB mass loss is revealed by spherical reflection nebulosities. Center (CRL 2688; [OI]): every time a reflection or neutral nebula is seen around a post-AGB star, a bipolar symmetry is present. Right (He 2-47; Hα): when post-AGB stars heat up, the young ionized PNe always have bipolar or multi-polar morphologies. [Credits: Mauron & Huggins (2000); Sahai et al. (1998); Sahai & Trauger (1998)]
The post-AGB phase is likely fundamental as it marks the onset of extreme asym- metry (see Fig. 1.1; Sahai & Trauger 1998). However, post-AGB stars are short lived with < 400 objects presently known (Szczerba et al. 2007).Therefore, this phase of post-main sequence evolution is not well studied. A number of physical processes may be involved in shaping post-AGB outflows. These include: magnetic launch & collimation, binary interactions, and accretion disk formation & evolution. Constraining formation theories observationally is difficult as CHAPTER 1. INTRODUCTION 4 the bipolar engines are often shielded by the remnant reflection nebula. Nevertheless, a growing body of observational and theoretical work suggests that binary interactions may be the unifying ingredient to deciphering asymmetry in this elusive phase of stellar evolution.
1.2 Support for the Binary Hypothesis
Here we present some of the evidence suggesting binarity is a common theme in inter- mediate mass, evolved star evolution.
1.2.1 Observational Indications
Recently, a comprehensive and extensive survey of CO emission revealed that nearly all post-AGB objects with known CO emission (∼ 80%) feature large momentum excess over what can be supplied by radiation pressure - often 103 − 104 times larger (Bujarrabal et al. 2001). Clearly, an additional momentum source must be present and a natural candidate is a binary companion. Binaries are attractive as two-thirds of all stars are thought to reside in multiple systems (Duquennoy & Mayor 1991). In addition, momentum can be exchanged from the companion star to the primary star through tidal, wind or common envelope interactions. Radial velocity (RV) surveys have been carried out to determine the fraction of central stars of PN. Preliminary results indicate that RV variability is common and that the true binary fraction may be of order unity (De Marco et al. 2004; Sorensen and Pollacco 2004; Afˇsar & Bond 2005). However, wind and pulsation also contribute to RV variability in evolved stars, making this fraction uncertain (De Marco et al. 2007). Adding additional credence to the binary hypothesis are population synthesis stud- ies that predict more PNe than observed (Moe & De Marco 2006). If it is assumed that PNe only form from binary interactions (particularly common envelope interactions), then population synthesis predictions are comparable with observations (De Marco 2006; Moe & De Marco 2006). These studies suggest the formation rate of PNe is 1/3 the formation rate of white dwarfs and support the claim that observed PNe are the descendants of interacting binary systems (Soker 2006; Soker & Subag 2005). While CHAPTER 1. INTRODUCTION 5 this percentage is lower then the percentage of binary stars, it may suggest only close to intermediate companions are responsible for asymmetric shaping.
1.2.2 Magnetic Shaping
It was previously suggested that an isolated star with a strong toroidal magnetic field could create the equatorial density enhancements needed for bipolar collimation in post-AGB/PNe (Garc´ıa-Segura et al. 1999, 2005). While these simulations were successful in recreating a wide array of shapes, the origin and evolution of the magnetic field was neglected. In particular, as large-scale magnetic fields amplify, differential rotation energy is drained from the engine. This back-reaction can shut down the dynamo, after which the fields decay. Nordhaus et al. (2007) demonstrated that an isolated star cannot sustain the neces- sary field strengths to power bipolarity unless convectivion actively re-supplies differ- ential rotation. However, if convection fails, not all is lost as a binary companion can supply the additional energy and angular momentum needed to amplify and sustain sufficient magnetic fields. This is particularly important as magnetic fields appear to be responsible for jet collimation in post-AGB objects (Vlemmings et al. 2006) and are present in the central stars of PNe (Jordan et al. 2005).
1.3 Thesis Structure
The research presented in this thesis aims to better understand the effect of a binary in an evolved star system.
Chapter 2 presents a study of the in-spiral of a low-mass companion (< 0.3 M : planet, brown dwarf, low-mass main sequence star) embedded in an AGB envelope during a common envelope (CE) phase. Common envelopes form when the expanding primary overflows its Roche lobe and both the companion and primary core become immersed (Iben & Livio (1993); Paczynski (1976)). During in-spiral, the secondary transfers additional energy and angular momentum to the CE and can eject it. For low mass companions, a CE phase can lead to three different mass ejection consequences: (i) direct ejection of envelope material resulting in a predominately equatorial outflow, (ii.) spin-up of the envelope resulting in the possibility of powering CHAPTER 1. INTRODUCTION 6 an explosive dynamo-driven jet and (iii.) tidal shredding of the companion into a disk which facilitates a disk-driven jet. These scenarios can launch material both equatorially and poloidally and may naturally lead to the disk-jet lag of a few hundred years observed by Huggins (2007). The common envelope interaction for low-mass companions is particularly relevant as the first white dwarf + brown dwarf post-CE binary was recently discovered (Maxted et al. 2006). As a continuation of this work, a dynamo operating in a CE is presented in detail in Chapter 3. The interplay of strong shear and a deep convective zone can generate dynamically important, large-scale magnetic fields via an α − Ω dynamo. As the mag- netic field strengthens, differential rotation energy is drained to support this growth. When the rigorous back-reaction of the field amplification on the shear is included, the dynamo becomes a transient phenomena lasting ∼ 100 yrs. In particular, for an iso- lated star, the magnetic field is not strong enough to provide the extreme shaping seen in many post-AGB/PNe unless convection can re-supply differential rotation during the AGB phase. However, in the common envelope case, a robust dynamo results and is able to supply the necessary energy and momentum to eject the envelope for a wide range of binary systems. This work suggests that the inclusion of MHD processes is essential for understanding the CE phase and missing from current 3-D, hydrodynamic simulations of common envelopes. Since binary companions are expected to influence outflow shapes, it is important to observationally determine the geometry of post-AGB stars. In Chapter 4, a spectral modeling strategy and technique are presented that can be applied to evolved stars to constrain the geometry (disks or spherical shells) of their ejected nebulae whether or not the sources are resolved in images. We apply this technique to HD 179821 and show explicitly where significant spectral degeneracies exist between the two geometries. The models can be distinguished spectrally only with data above 40 µm. In general, the importance of considering both disk-like and spherical geometries for evolved stellar nebulae is exacerbated by the fact that disks are very natural when significant angular momentum and/or companions are present (de Ruyter et al. 2006). When the sources are not spatially resolved, this technique constrains the geometry. Chapter 5 presents 3-D, hydrodynamic simulations of the interaction of an AGB wind with a low-mass companion in a wide orbit. The gravitational influence of the CHAPTER 1. INTRODUCTION 7 secondary focuses material in the equatorial plane. The companion induces spiral shocks which may anneal amorphous grains into crystalline dust. This work presents a physical mechanism to produce crystalline dust in AGB binary systems. This study is relevant for post-AGB systems in which high degrees of crystallinity appear to be observationally associated with binarity (Molster et al. 2002).
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Chapter 2
Low-mass Binary Induced Outflows from Asymptotic Giant Branch Stars
2.1 Abstract1
A significant fraction of planetary nebulae (PNe) and proto-planetary nebulae (PPNe) exhibit aspherical, axisymmetric structures, many of which are highly collimated. The origin of these structures is not entirely understood, however recent evidence suggests that many observed PNe harbor binary systems, which may play a role in their shap- ing. In an effort to understand how binaries may produce such asymmetries, we study the effect of low-mass (< 0.3 M ) companions (planets, brown dwarfs and low- mass main sequence stars) embedded into the envelope of a 3.0 M star during three epochs of its evolution (Red Giant Branch, Asymptotic Giant Branch (AGB), inter- pulse AGB). We find that common envelope evolution can lead to three qualitatively different consequences: (i) direct ejection of envelope material resulting in a predom- inately equatorial outflow, (ii) spin-up of the envelope resulting in the possibility of powering an explosive dynamo driven jet and (iii) tidal shredding of the companion into a disc which facilitates a disc driven jet. We study how these features depend on the secondary’s mass and discuss observational consequences.
1Originally published as Nordhaus & Blackman 2006 MNRAS 370, 2004 CHAPTER 2. LOW-MASS BINARY INDUCED OUTFLOWS 10
2.2 Introduction
Deviation from spherical symmetry in planetary nebulae (PNe) and protoplanetary nebular (PPNe) can be pronounced (Soker 1997, Balick & Frank 2002, Bujarrabal et al. 2004). Understanding the origin of the asymmetries is an ongoing aim of current research. A variety of scenarios have been proposed to explain the transition from progenitor to planetary nebula. As low- and intermediate-mass main sequence stars evolve onto the Asymptotic Giant Branch (AGB), enhanced stellar wind mass loss leads to deple- tion of the hydrogen envelope around the central core. Recent surveys (Sahai 2000, Sahai 2002) of AGB and post-AGB stars have revealed spherical symmetry leading to the conclusion that any shaping process must occur in a relatively short time just before the birth of the PNe in the PPNe, or AGB phase (Bujarrabal et al. 2001). Bipolar outflows in either the AGB or post-AGB phase could produce such struc- tures. However, the mechanism by which the bipolar winds are produced remains to be fully understood. Binary interactions, large-scale magnetic fields and high rotation rates of isolated AGB stars may all play some role in explaining the observed struc- tures: Frank et al. 1994 appealed to a superwind induced by a thermal helium flash as a possible production mechanism for bipolar planetary nebulae. Soker 2002 argued that such a model does not account for the observational link between aspherical mass loss and asymptotic wind. Such an observational correlation could be explained by a binary system in which the companion is a low-mass star or brown dwarf (Soker 2004). AGB and post-AGB remnant central stars are known to possess magnetic fields (Bains et al. 2004 and Jordan et al. 2005). In addition, direct evidence of a magneti- cally collimated jet in an evolved AGB star has been detected (Vlemmings et al. 2006), further suggesting some dynamical role for magnetic fields. Magnetic outflows from single stars have been proposed as mechanisms for shaping PPNe and PNe (Pascoli 1993; Blackman et al. 2001a). However, single star models may be unable to sustain the necessary Poynting flux required to maintain an outflow through the lifetime of the AGB phase unless differential rotation is reseeded by convection or supplied by a binary (Blackman 2004, Soker 2006). A model in which a disc driven magnetic dy- namo driven outflow is sustained by accretion from a shredded secondary was pursued in (Blackman et al. 2001b). CHAPTER 2. LOW-MASS BINARY INDUCED OUTFLOWS 11
In this respect, it is noteworthy that recent studies support the claim that most, if not all, observed planetary nebulae are the result of a binary interaction (De Marco et al. 2004, Sorensen and Pollacco 2004, De Marco & Moe 2005, Mauron & Huggins 2006). While this conclusion is based on observations and population synthesis studies, Soker 2006b points out that the corresponding formation rate of PNe from such studies is 1/3 the formation rate of white dwarfs. This supports the claim that binary stars may produce the more prominently observed planetary nebulae (Soker & Subag 2005). The question of just how a binary shapes a PNe or PPNe remains a topic of active research. In this paper, we explore the effects of an embedded low-mass companion inside the envelope of a 3 M star during three epochs of its evolution off the main sequence (Red Giant Branch (RGB), AGB and interpulse AGB). A common envelope (CE) facilitates mass ejection in several ways: The in-spiral of the secondary toward the core deposits orbital energy and angular momentum in the envelope. This directly ejects and/or spins up the envelope. In the latter case, any enhanced differential rotation could aid in magnetic field generation, which in turn could drive mass loss. In section 2, we describe the stellar models used and derive the basic equations for in-spiral and for the transfer of energy and angular momenutum from the secondary to the envelope. Results for different evolutionary epochs are discussed in section 3. We present observational implications and applications to specific systems in section 4 and conclude in section 5.
2.3 Common Envelope Evolution
Under certain conditions, Roche lobe overflow in close binary systems results in both companions immersed in a CE (Paczynski 1976, Iben & Livio 1993). Once inside, ve- locity differences between companion and envelope generate a drag force that acts to reduce the orbital separation of the companion and core. Orbital energy is deposited into the envelope during the in-spiral process. Some of this energy is radiated away while the rest is available to reduce the gravitational binding energy of the envelope. The efficiency with which orbital energy unbinds envelope matter is of central impor- tance to CE evolution. This is commonly incorporated into a parameter, α,which CHAPTER 2. LOW-MASS BINARY INDUCED OUTFLOWS 12 represents the fraction of orbital energy available for mass ejection as follows:
Ebind = α∆Eorb, (2.1) where ∆Eorb is the change in orbital energy of the binary and Ebind is the energy required to unbind envelope material. In principle, knowledge of the binding energy and α determines how much material is ejected, and in the case of complete envelope ejection, final binary separation distances. Such studies have been performed for a variety of different systems under a range of conditions of which the following are a small sample (Yungelson et al. 1995, Dewi & Tauris 2000, Taam & Sandquist 2000, Politano 2004). We investigate the effect of embedding planets, brown dwarfs and low-mass main sequence stars into an envelope of a 3 M star during various epochs of its evolution off of the main sequence. That the secondary mass represents a small perturbation to the initial envelope configuration allows us to neglect detailed radiative and hydro- dynamical effects. We present as simplified a picture as possible in order to elucidate basic phenomonological consequences of the interaction.
2.3.1 Envelope Binding Energy
Our stellar model consists of a 3 M main sequence star whose evolution is followed through the AGB phase with X =0.74 (mass fraction of hydrogen), Y =0.24 (mass fraction of helium), Z =0.02 and no mass loss (S. Kawaler - personal communication, Fig. 2.1). A range of evolutionary models allows consideration of various positions and times at which the expanding envelope engulfs the orbiting brown dwarf. We focus on three main epochs in the evolution: (i) near the tip of the first Red Giant Branch, (ii) the beginning of the Asymptotic Giant Branch and (iii) the quiet period between thermal pulses on the AGB branch. In each case, we calculate the energy required to unbind the envelope mass above a given radius, r (measured from the center of the primary’s core) as follows: MT GM(r) Ebind(r)=− dm(r), (2.2) M r where MT is the total mass of the star and M is the mass interior to the companions orbital radius. Here it is assumed that the core and envelope do not exchange energy CHAPTER 2. LOW-MASS BINARY INDUCED OUTFLOWS 13
Figure 2.1 Left: Density and mass profiles for our model AGB star. The dotted line is the core-envelope boundary. Right: Mach number and sound speed as a function of radius. The Mach number is computed from the Keplerian motion of the planet inside the envelope. The motion is supersonic everywhere and thus justifies our choice of accretion radius (Bondi 1952). during the CE phase and that ejection of material has no bearing on core structure. The values we determine for the binding energy in all three epochs are comparable to results from an estimation method first proposed by Webbink (1984) and further refined by Dewi & Tauris (2000) and Tauris & Dewi (2001). Explicitly calculating the binding energy for each evolutionary epoch fixes our efficiency parameter α between 0and1. 9 For the RGB star, our model core radius rc ∼ 3.5 × 10 cm with the envelope 11 extending out to a radius r ∼ 7 × 10 cm. At the chosen time in the RGB phase, thecorecontains0.41M and the energy required to unbind the entire envelope (∼ 1048 ergs) is the largest of our three epochs. Once the star has ascended onto the 9 AGB, the core contracts to a radius of 2.9 × 10 cm and the envelope expands to r 12 =5.7× 10 cm. The core has increased its mass to 0.55 M and the energy required to unbind the entire envelope decreases to 1.3 × 1047 ergs. For the interpulse AGB 9 phase, the core has contracted to a rc ∼ 1.6× 10 cm while the envelope has expanded 13 to r∗ ∼ 1.3 × 10 cm. The envelope binding energy has been further reduced to 5.7 × CHAPTER 2. LOW-MASS BINARY INDUCED OUTFLOWS 14
46 10 ergs with the core containing 0.58 M . We find that the interpulse AGB phase is most favorable for binary induced envelope ejection since the range of masses and radii required to deposit favorable orbital energy into the envelope is greatest in this phase. We discuss these results in detail in section 3.
2.3.2 Orbital Energy and Angular Momentum Evolution
The immersion of the secondary in the envelope of the giant results in a reduction of the separation distance between core and companion. To calculate the in-spiral and angular momentum transfer, we need to equate the rate of energy lost by drag to the change in gravitational potential energy. The motion of a body under the influence of a central potential while incurring a drag force has been well studied and a general set of equations can be found in Pollard 1979. Here we limit ourselves to the case where orbital eccentricity is negligible, such that the planet exhibits approximate Keplerian motion at each radii. Under these conditions, the energy per unit time released by the secondary mass takes the following form:
2 − 3 Ldrag = ξπRaρ(v venv) , (2.3) where v =(vr,vφ, 0) is the companion velocity, ρ theenvelopedensity,venv the en- velope velocity, Ra the accretion radius measured from the center of the companion, and ξ is a dimensionless factor dependent upon the Mach number of the companions motion with respect to the envelope. For supersonic motion, ξ is greater than or equal to 2 (Shima et al. 1985). The orbital motion of the planet is supersonic everywhere inside the orbit (see Fig. 2.1) and for simplicity, we assume a value of ξ =4.The value of ξ acts only to slightly increase or decrease the in-spiral time of the secondary, while leaving the underlying physics unchanged. The accretion radius is then given by
∼ 2Gm2 Ra 2 , (2.4) (v − venv) and represents the region around the secondary inside of which matter is gravitation- ally attracted to the secondary as it passes through the envelope. If the companion moves close enough to the core, tidal effects can shred it. We estimate the shredding radius (measured from the center of the primay’s core) from balancing the differential gravitational force across the size of the companion R2 (measured from the center CHAPTER 2. LOW-MASS BINARY INDUCED OUTFLOWS 15
d GM Gm2 of the companion) with its self gravity, that is, 2 R2 2 , which yields dr r R2 3 2M 2 rs m2 R . To determine the companion size, R2, we separate our objects into three distinct groups: planets (m2 ≤ 0.0026 M ; Zapolski & Salpeter 1969), brown dwarfs (0.0026
M