User:Tohline/Appendix/Ramblings/OriginOfPlanetaryNebulae

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On the Origin of Planetary Nebulae

This chapter — initially created by J. E. Tohline on 19 November 2016 — is intended primarily to provide a summary of the research that has been undertaken following a discussion that took place on 3 July 2013 with Kundan Kadam (an LSU graduate student, at the time) regarding the stability of bipolytropes.

"This approximation is as dangerously crude as it is computationally economic."

— Drawn from §I of Härm & Schwarzschild (1975)


Whitworth's (1981) Isothermal Free-Energy Surface
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Context

Why do stars become red giants? In particular, why does a star on the main sequence — whose internal density profile is only moderately centrally concentrated — become a red giant — which has a highly centrally condensed structure — at the end of the core-hydrogen-burning phase of its evolution? It seems likely that this evolutionary transition is triggered by an instability associated with the Schönberg-Chandrasekhar mass limit.

Rationale:   As hydrogen fuel is exhausted at the center of the star and burning shifts predominantly to a surrounding, off-center shell location, the helium core that is left behind is inert and approximately isothermal because the helium, itself, is not hot enough to burn. Henrich & Chandraskhar (1941) and Schönberg & Chandrasekhar (1942) discovered that equilibrium stellar structures with isothermal cores can be constructed, but only if the fraction of the star's mass that is contained in the core is below a well-defined, limiting value. This so-called Schönberg-Chandrasekhar mass limit was initially identified as a "turning point" along a sequence of equilibrium bipolytrope models in which the effective polytropic index of the core (c) and envelope (e) were, <math>~(n_c, n_e) = (\infty, 3/2)</math>. Evolution along this equilibrium sequence — toward the turning point — is naturally associated with stellar evolution off the main sequence. Specifically, one expects to see a slow (secular) but steady increase in the fraction of the star's mass that is enclosed within the isothermal core as the hydrogen-burning shell slowly works its way outward from the center. An examination of the bipolytropic models along this sequence also reveals that, as the mass of the isothermal core increases, the star's equilibrium structure becomes more and more centrally condensed. As a result — as has been emphasized by, for example, Eggleton, Faulkner, and Cannon (1998) — the substantial structural change that occurs in a star as it evolves from the main sequence toward the red giant branch may be simply a natural consequence of evolution toward the Schönberg-Chandrasekhar mass limit.

Motivating Questions:   With this general scenario in mind, we began to wonder — as, almost certainly, other astrophysicists before us have wondered:

  • What type of instability — dynamical or secular (?) — is associated with the equilibrium sequence turning point that is synonymous with the Schönberg-Chandrasekhar mass limit?
    • If it is a secular instability, do stars normally find a way — via one or more secular mechanisms — to readjust their structure as they approach the turning point and avoid encountering the mass limit altogether?
    • If it is a dynamical instability, what is the — presumably catastrophic — result of encountering the Schönberg-Chandrasekhar mass limit? Does the core collapse on a free-fall time scale; is the envelope ejected instead? Or, perhaps envelope ejection occurs in concert with the core's collapse?
  • Might evolution toward the Schönberg-Chandrasekhar mass limit be hastened in situations where the hydrogen-shell-burning (bipolytropic) star has a binary companion? The natural, gradual expansion of the star's envelope as it evolves off of the main sequence may bring its surface into contact with the binary system's Roche lobe and, as a result, some of the star's mass will be transferred to its companion. This means that, even if the amount of mass contained within the inert helium (isothermal) core does not increase, the fraction of the star's mass that is contained in the core is destined to increase because the star's total mass is decreasing as a result of mass transfer. If the mass-transfer rate is high enough, perhaps the secular mechanisms that help an isolated star avoid the Schönberg-Chandrasekhar mass limit will not have sufficient time to operate and, as a result, the evolving star is pushed past the limit. How catastrophic is this?
  • Perhaps envelope ejection — and the consequential development of wonderfully photogenic planetary nebulae — is a natural outcome of evolving stars encountering the Schönberg-Chandrasekhar mass limit. And perhaps a star is more likely to be pushed to/past this limit if it has a binary companion.

Proposed Numerical Investigation:   The LSU astrophysics group ought to employ its three-dimensional hydrodynamic code to investigate what happens when a bipolyropic star (the donor) — with an isothermal (or nearly isothermal) core and a core-to-total mass ratio that is near the Schönberg-Chandrasekhar mass limit — fills its Roche lobe and transfers mass to its stellar companion (the accretor). After a fairly predictable amount of (envelope) mass has been transferred from the donor to the accretor, the donor should encounter the Schönberg-Chandrasekhar mass limit. What will the result be? Does the initially bipolytropic donor's internal structure readjust on a dynamical time scale in response to this encounter? Does its core collapse; and/or does its envelope rapidly expand?

Sub-Projects Undertaken

In order for the above proposed numerical investigation to provide informative results, it is important that we establish a firm understanding of a variety of related, but less complicated, concepts and problems. An emphasis has been placed on tackling problems that can described as fully as possible using analytic, rather than purely numerical, techniques. The following subsections provide a list, along with brief description, of related sub-projects that we have studied, to date.

Bipolytropes

Textbook Explanations

Clayton (1968)

Here, we consider the descriptions presented by D. D. Clayton (1968).

Evolution to the Red-Giant Branch: (§6-7, p. 485) "The core continues to contract as the hydrogen is exhausted, leaving a central region of helium plus heavier trace elements. This helium core will tend to be isothermal because nuclear energy generation has ceased …" As the star's evolution proceeds, the temperature of the inert (isothermal) core will continue to increase, as will "the temperature of a shell of hydrogen surrounding the core … The increased internal temperatures require the expansion of the stellar radius to keep the temperature gradient at a consistently low level. The star therefore reddens at a relatively rapid rate while the hydrogen-burning shell slowly increases the mass of the helium core."

Stellar Pulsation: (§6-10, p. 504) "By 1930 it was clear, thanks largely to the work of Eddington, that a pulsating star must in fact be some type of heat engine, in which some continuously operating mechanism transforms thermal energy into the mechanical energy of the oscillation." Analyses that attempt to explain the existence and properties of regular variable stars — such as the Cepheids and RR Lyrae variables — focus on stellar (envelope) structures that are dynamically stable, according to adiabatic stability analyses, but that harbor a tendency toward growing oscillatory amplitude when non-adiabatic effects are considered. Specifically referencing the three terms in equation (6-116) on p. 511, we can identify the principal "… physical effects contributing to the status of the stability of the zone."

  • Γ mechanism: "The first term always contributes to stability … [but its] influence is diminished in ionization zones."
  • κ mechanism: "The second term reflects the way in which the opacity varies during the pulsation. Positive values of <math>~\kappa_T</math> and <math>~\kappa_P</math> would imply that the opacity increases upon contraction, which would remove energy from the radiation flux … at the proper time to drive mechanical work."

Mass Loss: (§6-9, p. 501) "Mass loss is a self-descriptive term that is used to describe any process by which the main body of the star, defined as the gravitationally bound mass, reduces its mass by ejecting surface layers … Mass loss can occur in a variety of forms and can be initiated by a variety of physical mechanisms. Any catastrophic event in which a massive outer layer is lifted off into space by some internal instability must result in a drastically new structure for the remaining core. So special are these circumstances that they will not be discussed here.."

Rose (1998)

Here, we consider the descriptions presented by W. K. Rose (1998).

Evolution to the Red-Giant Branch:


Stellar Pulsation:


Mass Loss & Formation of Planetary Nebula:

See Also


(p. 663) "The parameters of the model close to dynamical instability were <math>~M = 1.3 M_\odot</math> … and core mass-fraction <math>~=0.20</math>. This model is the same as model 2.5 in Paper I = D. A. Keeley (1970, ApJ, Vol. 161, p. 643) … This model was calculated to test the possibility that an extreme red giant could eject its envelope and form a planetary nebula. G. O. Abell & P. Goldreich (1966, PASP, Vol. 78, P. 232) presented arguments in favor of this hypothesis. The problem was subsequently pursued by B. Paczynski (1968, Acta Astr., Vol. 18, p. 511), B. Paczynski & J. Ziólkowski (1968a) and (1968b) and the writer."

(p. 398) "We propose the following scheme for the late phases of stellar evolution. After the exhaustion of helium [sic] in the core the star evolves into the region of red supergiants and moves up on the H-R diagram very close to the Hayashi border … The mass of the helium and carbon core and the luminosity due to the helium and hydrogen-shell sources increase. A star with total mass smaller than about 4 <math>~M_\odot</math> will terminate this type of evolution with an outflow of hydrogen-rich matter as a result of the dynamical instability of the extended envelope. We suggest that planetary nebulae are formed in this way"

(p. 265) "We suggest that a planetary nebula is formed in this way. This process seems to be impossible for a star with larger mass. In the latter case the mass of the core in which all the nuclear energy sources are exhausted will finally exceed the Chandrasekhar limit for degenerate configurations. The dynamical instability of the core, followed by a supernova explosion may be expected …"

(p. 98) "Of the three mechanisms considered earlier for the separation of the planetary nebula, preference must clearly be given to the ejection of a massive shell at a single time; for it is free of many of the difficulties inherent in the hypothesis of gradual accumulation of matter in the neighborhood of the star that is the precursor of the planetary nebula and its central star."

(§ VI, p. 711) "… we summarize the major concepts which are currently accepted for the evolution through the PN stage."

  • i) "Renzini (1981) first considered that the very existence of PNs implies that, at a given time, the mass loss rate from asymptotic giant branch (AGB) stars must become much larger than the normal 'wind' rate (Reimers 1975). This 'superwind' must be at least several times <math>~10^{-5} M_\odot~\mathrm{yr}^{-1}</math>.
  • ii) "The precise phase during the AGB evolution at which the super wind sets in is also a fundamental parameter for the following evolution, as the helium shell burning is more or less efficient at different stages. Iben (1984) followed numerically a number of cases, showing, for instance, how the complete loss of the hydrogen layer may be achieved if, after the PN ejection, there is the possibility of igniting a final helium shell flash."
  • iii) "It is not clear at all whether PN ejection is a hydrodynamical event or not. The few existing hydrodynamic computations (e.g., Kutter and Sparks 1974) do not consider the resulting behavior of the hydrostatic remnant, whose further evolution to the blue is crucial for the appearance of the PN. many of the non hydrodynamical computations — starting from a fundamental paper by Härm and Schwarzschild (1975) and end with Iben (1984) — simulate the dynamics ejection by 'scaling' the models before ejection to smaller masses. It is not proven that this procedure is physically meaningful."
Whitworth's (1981) Isothermal Free-Energy Surface

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