DOCSLIB.ORG

On Constructing Horizontal Branch Models

Total Page：16

File Type：pdf, Size：1020Kb

Download full-text PDF Read full-text
On Constructing Horizontal Branch Models A&A 442, 1041–1048 (2005) Astronomy DOI: 10.1051/0004-6361:20053390 & c ESO 2005 Astrophysics On constructing horizontal branch models A. Serenelli1,2 and A. Weiss1 1 Max-Planck-Institut für Astrophysik, Karl-Schwarzschild-Str. 1, 85748 Garching, Germany e-mail: [email protected] 2 Institute for Advanced Study, Einstein Drive, Princeton, NJ 08540, USA Received 9 May 2005 / Accepted 28 July 2005 ABSTRACT We investigate different methods used to construct (zero-age) horizontal branch models and compare the resulting horizontal branch evolution with that of models resulting from the calculation of the complete stellar evolution from the main sequence and through the core helium flash. We show that the approximate methods may lead to small but discernible effects, but that some methods, which are as simple, reproduce the complete evolution very well. Key words. stars: evolution – stars: horizontal-branch – methods: numerical primarily by the composition (the “first parameter”) and a “sec- 1. Introduction ond parameter”, the nature of which is so far unknown, al- The horizontal branch (HB) in the Hertzsprung-Russell- though age is considered to be a likely candidate (see, for exam- Diagram (HRD) has a number of important applications. For ple Lee et al. 1994; Sweigart & Catelan 1998; Catelan 2000). example, it allows distance-independent age determinations One should add that here secondary age effects are considered, of globular clusters, when the difference between the time- because age itself is a primary parameter, due to the fact that dependent turn-off brightness and the almost time-independent the turn-off mass is already influencing the HB morphology be- one of the HB is used (e.g., Salaris & Weiss 1997). On the cause younger populations, i.e. more massive stars will become other hand, its brightness or that of RR Lyr stars (Feast 1997) more massive, therefore cooler, HB stars. populating it can be used to determine distances. The number The evolution to an HB star goes through the violent core of stars on the HB relative to the number of stars on the Red helium flash. For several decades, only very few calculations Giant Branch (RGB) above the HB is an indicator for the he- have been able to follow stellar models through this event. The lium content of the cluster (Buzzoni et al. 1983; Salaris et al. reason is that typical timesteps decrease to the level of seconds 2004). because of the rapid change in helium luminosity, which may 10 Physically, HB stars are in the core helium burning phase flare up to 10 L for a few days. Because other parts of the following the RGB. The ignition of helium burning takes place star remain completely unchanged during this time, the differ- under degenerate conditions and is a very energetic and par- encing schemes become numerically unstable and no longer > tially dynamic event. Stars of initial mass Mi ∼ 2.3 M ig- converge. Therefore, full flash calculations were successful in nite Helium under non-degenerate conditions and populate the most cases only for higher masses, where degeneracy of the equivalent of the HB, the so-called clump, which is observa- RGB stellar core is lower and the flash milder (see Thomas tionally a region of increased star density close to or on top of 1967, for a pioneering study). With a few exceptions, complete the RGB at the level of the HB. stellar evolution tracks into the HB phase were absent. The distribution of post-flash low-mass stars along a hor- For this reason most HB models were calculated by starting izontal sequence can be explained by a constant core mass, new sequences on the HB, where the initial structure is taken which is closely related to the luminosity of the star, and a from that at the tip of the giant branch, and solutions of the spread in envelope and thus total mass (Faulkner 1966; Dorman structure equations are sought where the effect of the helium 1992). This determines the effective temperature, but only flash is taken into account. This implies a reduced core density, weakly the luminosity. In a globular cluster, due to the fact that hydrostatic helium burning, a lower surface luminosity and an all stars have the same age, the initial mass of the stars now effective temperature depending on the envelope mass. Details populating the HB must have been the same and the spread of such approximate methods will be given later in the paper. along the HB must be the result of different mass loss histories. Overall, such starting models represent the true zero-age HB The detailed distribution of HB stars as function of Teff, (ZAHB) and the following evolution quite well, but there is the so-called HB morphology, appears to be determined no systematic investigation of the effects of the approximation. Article published by EDP Sciences and available at http://www.edpsciences.org/aa or http://dx.doi.org/10.1051/0004-6361:20053390 1042 A. Serenelli and A. Weiss: On constructing horizontal branch models The underlying assumption of such methods is that during the Table 1. List of models for which we did full evolutions through the helium flash no events take place that modify the internal chem- helium flash. The columns list metallicity, initial mass, the parameter η ical structure significantly. in Reimers’ mass loss formula, total mass Mf after the core helium Since a few years an increasing number of stellar evolu- flash, core mass Mc, and convective core mass Mcc at the ZAHB. tion codes have been improved numerically to a level where / η / / / full flash calculations have become possible. Still, the compu- ZMM Mf M Mc M Mcc M tational effort is large, such that grids of HB models derived 0.0001 0.80 0.0 0.7971 0.5011 0.131 from complete evolution sequences are not truely practicable. 0.3 0.6788 0.5014 0.134 However, it is now possible to compare full with approximate 0.5 0.5863 0.5015 0.133 HB models and to design schemes for the construction of HB 0.85 0.0 0.8471 0.4996 0.132 models which are both economical and accurate. Given the fact 0.5 0.6574 0.4999 0.131 that the requirements of accuracy of stellar models is steadily 0.7 0.5632 0.5001 0.132 0.003 0.82 0.0 0.8170 0.4892 0.121 increasing, either from observational data about stars, or – as in 0.2 0.6650 0.4895 0.121 the case of the cluster ages – because of more detailed questions 0.3 0.5734 0.4897 0.122 to be answered, we must consider the accuracy of HB models. 0.90 0.0 0.8970 0.4878 0.121 Recently, Piersanti et al. (2004) published a similar 0.3 0.6906 0.4879 0.120 and independent investigation testing their own method for 0.5 0.5085 0.4880 0.119 ZAHB model construction. They showed results for two cases, to which we will refer. In Sect. 2 we will discuss briefly our stellar evolution code obviously violate the basic assumption of the approximate and present full core helium flash calculations and the resulting models of an unmodified composition profile between the onset ZAHB models and HB evolution. In Sect. 3 we present sev- of the flash and the ZAHB. eral typical schemes to create approximate HB models without We note that convection is always treated in the most following the complete flash. Section 4 is devoted to critically simple way, i.e. using mixing length theory and employing analyzing results for approximate models by using results of the Schwarzschild criterion for convective stability. We thus Sect. 2 as a reference. We will not investigate methods that rely consider neither overshooting nor semiconvection. Both ef- on a completely new construction of models, since today they fects play a role during the HB evolution, but need addi- are of only historical interest (Rood 1970; Sweigart & Gross tional assumptions and introduce free parameters, which usu- 1976). All calculations have been done with the same code to ally are determined from comparison with observations. The avoid differences resulting from details in the numerical pro- inclusion would not change our results significantly, because grams. Section 5 provides a short summary. we are mainly interested in the quality of the constructed ZAHB models. 2. Complete evolution from the ZAMS to the HB 2.1. Code 2.2. Calculations . Throughout this paper we use the Garching stellar evolution We have evolved low-mass stars between 0 8and09 M,and = . code, which has been described in some detail in Weiss & compositions of Z 0 0001 and 0 003. The initial helium con- = . Schlattl (2000). The numerical aspects of this code have been tent was always Y 0 25. A list of cases is given in Table 1. ff improved such that the evolution of Asymptotic Giant Branch With these parameters we cover turn-o ages between 8 and (AGB) stars can be followed through many thermal pulses 12 Gyr years and the lower metallicity range of galactic glob- without convergence problems and intervention (Wagenhuber ular clusters. The evolution was followed from the zero-age & Weiss 1994). Due to the same improvements in the algo- main-sequence (ZAMS), starting from a homogeneous model rithms controlling the temporal and spatial discretization the up to the lower AGB, that is, prior to the onset of thermal helium core flash also became feasible (Wagenhuber 1996), be- pulses. For mass loss we used the standard Reimers formula cause the numerical problems encountered in both phases are (Reimers 1975), of a similar nature.
Load more

Recommended publications
  • The Variable Stars and Blue Horizontal Branch of the Metal-Rich Globular Cluster NGC 6441
    View metadata, citation and similar papers at core.ac.uk brought to you by CORE provided by CERN Document Server The Variable Stars and Blue Horizontal Branch of the Metal-Rich Globular Cluster NGC 6441 Andrew C. Layden1;2 Physics & Astronomy Dept., Bowling Green State Univ., Bowling Green, OH 43403, U.S.A. Laura A. Ritter Department of Astronomy, University of Michigan, Ann Arbor, MI 48109-1090, U.S.A. Douglas L. Welch1,TracyM.A.Webb1 Department of Physics & Astronomy, McMaster University, Hamilton, Ontario L8S 4M1, Canada ABSTRACT We present time-series VI photometry of the metal-rich ([Fe=H]= 0:53) globular − cluster NGC 6441. Our color-magnitude diagram shows that the extended blue horizontal branch seen in Hubble Space Telescope data exists in the outermost reaches of the cluster. About 17% of the horizontal branch stars lie blueward and brightward of the red clump. The red clump itself slopes nearly parallel to the reddening vector. A component of this slope is due to differential reddening, but part is intrinsic. The blue horizontal branch stars are more centrally concentrated than the red clump stars, suggesting mass segregation and a possible binary origin for the blue horizontal branch stars. We have discovered 50 new variable stars near NGC 6441, among ∼ them eight or more RR Lyrae stars which are highly-probable cluster members. Comprehensive period searches over the range 0.2–1.0 days yielded unusually long periods (0.5–0.9 days) for the fundamental pulsators compared with field RR Lyrae of the same metallicity. Three similar long-period RR Lyrae are known in other metal-rich globulars.
    [Show full text]
  • (GK 1; Pr 1.1, 1.2; Po 1) What Is a Star?
    1) Observed properties of stars. (GK 1; Pr 1.1, 1.2; Po 1) What is a star? Why study stars? The sun Age of the sun Nearby stars and distribution in the galaxy Populations of stars Clusters of stars Distances and characterization of stars Luminosity and flux Magnitudes Parallax Standard candles Cepheid variables SN Ia 2) The HR diagram and stellar masses (GK 2; Pr 1.4; Po 1) Colors of stars B-V Blackbody emission HR Diagram Interpretation of HR diagram stellar radii kinds of stars red giants white dwarfs planetary nebulae AGB stars Cepheids Horizontal branch evolutionary sequence turn off mass and ages Masses from binaries Circular orbits solution General solution Spectroscopic binaries Eclipsing spectroscopic binaries Empirical mass luminosity relation 3) Spectroscopy and abundances (GK 1; Pr 2) Stellar spectra OBFGKM Atomic physics H atom others Spectral types Temperature and spectra Boltzmann equation for levels Saha equation for ionization ionization stages Rotation Stellar Abundances More about ionization stages, e.g. Ca and H Meteorite abundances Standard solar set Abundances in other stars and metallicity 4) Hydrostatic balance, Virial theorem, and time scales (GK 3,4; Pr 1.3, 2; Po 2,8) Assumptions - most of the time Fully ionized gas except very near surface where partially ionized Spherical symmetry Broken by e.g., convection, rotation, magnetic fields, explosion, instabilities, etc Makes equations a lot easier Limits on rotation and magnetic fields Homogeneous composition at birth Isolation (drop this later in course) Thermal
    [Show full text]
  • Hydrogen-Deficient Stars
    Hydrogen-Deficient Stars ASP Conference Series, Vol. 391, c 2008 K. Werner and T. Rauch, eds. Hydrogen-Deficient Stars: An Introduction C. Simon Jeffery Armagh Observatory, College Hill, Armagh BT61 9DG, N. Ireland, UK Abstract. We describe the discovery, classification and statistics of hydrogen- deficient stars in the Galaxy and beyond. The stars are divided into (i) massive / young star evolution, (ii) low-mass supergiants, (iii) hot subdwarfs, (iv) cen- tral stars of planetary nebulae, and (v) white dwarfs. We introduce some of the challenges presented by these stars for understanding stellar structure and evolution. 1. Beginning Our science begins with a young woman from Dundee in Scotland. The brilliant Williamina Fleming had found herself in the employment of Pickering at the Harvard Observatory where she noted that “the spectrum of υ Sgr is remarkable since the hydrogen lines are very faint and of the same intensity as the additional dark lines” (Fleming 1891). Other stars, well-known at the time, later turned out also to have an unusual hydrogen signature; the spectacular light variations of R CrB had been known for a century (Pigott 1797), while Wolf & Rayet (1867) had discovered their emission-line stars some forty years prior. It was fifteen years after Fleming’s discovery that Ludendorff (1906) observed Hγ to be completely absent from the spectrum of R CrB, while arguments about the hydrogen content of Wolf-Rayet stars continued late into the 20th century. Although these early observations pointed to something unusual in the spec- tra of a variety of stars, there was reluctance to accept (or even suggest) that hydrogen might be deficient.
    [Show full text]
  • The Deaths of Stars
    The Deaths of Stars 1 Guiding Questions 1. What kinds of nuclear reactions occur within a star like the Sun as it ages? 2. Where did the carbon atoms in our bodies come from? 3. What is a planetary nebula, and what does it have to do with planets? 4. What is a white dwarf star? 5. Why do high-mass stars go through more evolutionary stages than low-mass stars? 6. What happens within a high-mass star to turn it into a supernova? 7. Why was SN 1987A an unusual supernova? 8. What was learned by detecting neutrinos from SN 1987A? 9. How can a white dwarf star give rise to a type of supernova? 10.What remains after a supernova explosion? 2 Pathways of Stellar Evolution GOOD TO KNOW 3 Low-mass stars go through two distinct red-giant stages • A low-mass star becomes – a red giant when shell hydrogen fusion begins – a horizontal-branch star when core helium fusion begins – an asymptotic giant branch (AGB) star when the helium in the core is exhausted and shell helium fusion begins 4 5 6 7 Bringing the products of nuclear fusion to a giant star’s surface • As a low-mass star ages, convection occurs over a larger portion of its volume • This takes heavy elements formed in the star’s interior and distributes them throughout the star 8 9 Low-mass stars die by gently ejecting their outer layers, creating planetary nebulae • Helium shell flashes in an old, low-mass star produce thermal pulses during which more than half the star’s mass may be ejected into space • This exposes the hot carbon-oxygen core of the star • Ultraviolet radiation from the exposed
    [Show full text]
  • White Dwarfs in Globular Clusters 3 Received Additional Support from the Theoretical Investigations of [94, 95, 96]
    White Dwarfs in Globular Clusters S. Moehler1 and G. Bono1,2,3 1 European Southern Observatory, Karl-Schwarzschild-Str. 2, 85748 Garching, Germany, [email protected] 2 Dept. of Physics, Univ. of Rome Tor Vergata, via della Ricerca Scientifica 1, 00133 Rome, Italy, [email protected] 3 INAF-Osservatorio Astronomico di Roma, via Frascati 33, 00040 Monte Porzio Catone, Italy We review empirical and theoretical findings concerning white dwarfs in Galactic globular clusters. Since their detection is a critical issue we describe in detail the various efforts to find white dwarfs in globular clusters. We then outline the advantages of using cluster white dwarfs to investigate the forma- tion and evolution of white dwarfs and concentrate on evolutionary channels that appear to be unique to globular clusters. We also discuss the usefulness of globular cluster white dwarfs to provide independent information on the distances and ages of globular clusters, information that is very important far beyond the immediate field of white dwarf research. Finally, we mention pos- sible future avenues concerning globular cluster white dwarfs, like the study of strange quark matter or plasma neutrinos. 1 Introduction During the last few years white dwarfs have been the topic of several thorough review papers focused on rather different aspects. The interested reader is referred to [85] and to [86] for a comprehensive discussion concerning the use of white dwarfs as stellar tracers of Galactic stellar populations and the physics of cool white dwarfs. The advanced evolutionary phases and their impact on the dynamical evolution of open and globular clusters have been reviewed by [116], while [4] provide a comprehensive discussion of the use of arXiv:0806.4456v3 [astro-ph] 30 Jun 2011 white dwarfs to constrain stellar and cosmological parameters together with a detailed analysis of the physical mechanisms driving their evolutionary and pulsation properties.
    [Show full text]
  • (NASA/Chandra X-Ray Image) Type Ia Supernova Remnant – Thermonuclear Explosion of a White Dwarf
    Stellar Evolution Card Set Description and Links 1. Tycho’s SNR (NASA/Chandra X-ray image) Type Ia supernova remnant – thermonuclear explosion of a white dwarf http://chandra.harvard.edu/photo/2011/tycho2/ 2. Protostar formation (NASA/JPL/Caltech/Spitzer/R. Hurt illustration) A young star/protostar forming within a cloud of gas and dust http://www.spitzer.caltech.edu/images/1852-ssc2007-14d-Planet-Forming-Disk- Around-a-Baby-Star 3. The Crab Nebula (NASA/Chandra X-ray/Hubble optical/Spitzer IR composite image) A type II supernova remnant with a millisecond pulsar stellar core http://chandra.harvard.edu/photo/2009/crab/ 4. Cygnus X-1 (NASA/Chandra/M Weiss illustration) A stellar mass black hole in an X-ray binary system with a main sequence companion star http://chandra.harvard.edu/photo/2011/cygx1/ 5. White dwarf with red giant companion star (ESO/M. Kornmesser illustration/video) A white dwarf accreting material from a red giant companion could result in a Type Ia supernova http://www.eso.org/public/videos/eso0943b/ 6. Eight Burst Nebula (NASA/Hubble optical image) A planetary nebula with a white dwarf and companion star binary system in its center http://apod.nasa.gov/apod/ap150607.html 7. The Carina Nebula star-formation complex (NASA/Hubble optical image) A massive and active star formation region with newly forming protostars and stars http://www.spacetelescope.org/images/heic0707b/ 8. NGC 6826 (Chandra X-ray/Hubble optical composite image) A planetary nebula with a white dwarf stellar core in its center http://chandra.harvard.edu/photo/2012/pne/ 9.
    [Show full text]
  • Announcements
    Announcements • Next Session – Stellar evolution • Low-mass stars • Binaries • High-mass stars – Supernovae – Synthesis of the elements • Note: Thursday Nov 11 is a campus holiday Red Giant 8 100Ro 10 years L 10 3Ro, 10 years Temperature Red Giant Hydrogen fusion shell Contracting helium core Electron Degeneracy • Pauli Exclusion Principle says that you can only have two electrons per unit 6-D phase- space volume in a gas. DxDyDzDpxDpyDpz † Red Giants • RG Helium core is support against gravity by electron degeneracy • Electron-degenerate gases do not expand with increasing temperature (no thermostat) • As the Temperature gets to 100 x 106K the “triple-alpha” process (Helium fusion to Carbon) can happen. Helium fusion/flash Helium fusion requires two steps: He4 + He4 -> Be8 Be8 + He4 -> C12 The Berylium falls apart in 10-6 seconds so you need not only high enough T to overcome the electric forces, you also need very high density. Helium Flash • The Temp and Density get high enough for the triple-alpha reaction as a star approaches the tip of the RGB. • Because the core is supported by electron degeneracy (with no temperature dependence) when the triple-alpha starts, there is no corresponding expansion of the core. So the temperature skyrockets and the fusion rate grows tremendously in the `helium flash’. Helium Flash • The big increase in the core temperature adds momentum phase space and within a couple of hours of the onset of the helium flash, the electrons gas is no longer degenerate and the core settles down into `normal’ helium fusion. • There is little outward sign of the helium flash, but the rearrangment of the core stops the trip up the RGB and the star settles onto the horizontal branch.
    [Show full text]
  • Evolution of Stars
    Evolution of Stars (Part I: Solar-type stars) 1 Death of the Sun Parts I and II 2 Learning goals: Be able to …. ! summarize the future of the Sun on a rough timescale; ! apply the basics of the conservation of energy and the battle between gravity and outward pressure to what “drives” a star to evolve at each major stage of evolution; ! explain what is meant by main sequence, subgiant, red giant branch, electron degeneracy, helium flash, horizontal branch, asymptotic giant branch, planetary nebula, white dwarf. 3 Be able to summarize the future of the Sun in a rough timescale. Apply the basics of the conservation of energy and the battle between gravely and pressure to what “drives” a star to evolve at each major stage of evolution. Explain what is meant by red giant branch, electron degeneracy, helium flash, horizontal branch, planetary nebula, white dwarf 4 5 http://chandra.harvard.edu/xray_sources/browndwarf_fg.html Brown Dwarf "In between a star and a planet "Jupiter < BD mass < 0.08 Sun "Radiates in infrared due to low temperature. "Mass too low to start fusion in core. "Slowly cools over trillions of years. 6 7 Core producing energy--Fusing H to He 8 A review of what we know the Sun is doing as a main-sequence star 9 Number of particles in core is decreasing. NET RESULT IN: 6 H OUT: 1 He + 2 H FUSION RATE MUST INCREASE TO OFFSET DECREASE IN PARTICLE PRESSURE! Sun is slowly becoming more and more luminous. 10 Explain how the Sun maintains a constant balance in its interior (solar thermostat) Part of the pressure depends on the number of particles Pressure decreases, core shrinks slightly, pressure evens out, fusion rate must increase to offset the smaller pressure provided by the number of particles present.
    [Show full text]
  • Accretion Regimes in the X-Ray Pulsar 4U 1901+03 P
    A&A 594, A45 (2016) Astronomy DOI: 10.1051/0004-6361/201629200 & c ESO 2016 Astrophysics Accretion regimes in the X-ray pulsar 4U 1901+03 P. Reig1; 2 and F. Milonaki2 1 IESL, Foundation for Reseach and Technology-Hellas, 71110 Heraklion, Greece e-mail: [email protected] 2 Physics Department, University of Crete, 71003 Heraklion, Greece Received 28 June 2016 / Accepted 17 August 2016 ABSTRACT Context. The source 4U 1901+03 is a high-mass X-ray pulsar than went into outburst in 2003. Observation performed with the Rossi X-ray Timing Explorer showed spectral and timing variability, including the detection of flares, quasi-periodic oscillations, complex changes in the pulse profiles, and pulse phase dependent spectral variability. Aims. We re-analysed the data covering the 2003 X-ray outburst and focused on several aspects of the variability that have not been discussed so far. These are the 10 keV feature and the X-ray spectral states and their association with accretion regimes, including the transit to the propeller state at the end of the outburst. Methods. We extracted light curves and spectra using data from the Rossi X-ray Timing Explorer. Low time resolution light curves were used to create hardness-intensity diagrams and study daily changes in flux. High time resolution light curves were used to create pulse profiles. An average spectrum per observation allowed us to investigate the evolution of the spectral parameters with time. Results. We find that 4U 1901+03 went through three accretion regimes over the course of the X-ray outburst.
    [Show full text]
  • Stellar Evolution: Evolution Off the Main Sequence
    Evolution of a Low-Mass Star Stellar Evolution: (< 8 M , focus on 1 M case) Evolution off the Main Sequence sun sun - All H converted to He in core. - Core too cool for He burning. Contracts. Main Sequence Lifetimes Heats up. Most massive (O and B stars): millions of years - H burns in shell around core: "H-shell burning phase". Stars like the Sun (G stars): billions of years - Tremendous energy produced. Star must Low mass stars (K and M stars): a trillion years! expand. While on Main Sequence, stellar core has H -> He fusion, by p-p - Star now a "Red Giant". Diameter ~ 1 AU! chain in stars like Sun or less massive. In more massive stars, 9 Red Giant “CNO cycle” becomes more important. - Phase lasts ~ 10 years for 1 MSun star. - Example: Arcturus Red Giant Star on H-R Diagram Eventually: Core Helium Fusion - Core shrinks and heats up to 108 K, helium can now burn into carbon. "Triple-alpha process" 4He + 4He -> 8Be + energy 8Be + 4He -> 12C + energy - First occurs in a runaway process: "the helium flash". Energy from fusion goes into re-expanding and cooling the core. Takes only a few seconds! This slows fusion, so star gets dimmer again. - Then stable He -> C burning. Still have H -> He shell burning surrounding it. - Now star on "Horizontal Branch" of H-R diagram. Lasts ~108 years for 1 MSun star. More massive less massive Helium Runs out in Core Horizontal branch star structure -All He -> C. Not hot enough -for C fusion. - Core shrinks and heats up.
    [Show full text]
  • Today's Outline
    Today's outline Review High and low mass Low-mass Stars Summary of evolution Main sequence Hydrogen Exhaustion Review: Clusters, Birth of Stars Giant phase Helium flash Horizontal Branch Evolution of low mass stars Helium Exhaustion Planetary Nebula I Low and high mass stars Summary again I Interior of a giant star I Phases of burning I White dwarf formation Review High and low mass Low-mass Stars Summary of evolution Main sequence Evolution of low mass stars Hydrogen Exhaustion Giant phase I Low and high mass stars Helium flash Horizontal Branch Helium Exhaustion I Interior of a giant star Planetary Nebula Summary again I Phases of burning I White dwarf formation Reference stars Low Mass: Review I Sun - low-mass "dwarf" High and low mass Low-mass Stars I Vega - low-mass "dwarf" Summary of evolution Main sequence I Sirius - low-mass "dwarf" Hydrogen Exhaustion Giant phase Helium flash I Arcturus - low-mass giant Horizontal Branch Helium Exhaustion I Sirius B - white dwarf (very small) Planetary Nebula Summary again High Mass: I Rigil - high-mass (blue) giant I Betelgeuse - high-mass (red) giant High and low mass stars Black hole Review "high−mass" High and low mass Low-mass Stars (hydrogen burning) Neutron Star Summary of evolution Main Sequence Main sequence 8 Hydrogen Exhaustion Giants Giant phase Protostar Helium flash Horizontal Branch "low−mass" Helium Exhaustion Planetary Nebula 2 White dwarf Summary again Birth Mass Time Stars generally classified by their end-of-life I Low mass { form white dwarf stars, no supernova I High mass { form
    [Show full text]
  • 37 — the End of Low-Mass Stars [Revision : 1.2]
    37 — The End of Low-Mass Stars [Revision : 1.2] • Burning to a degenerate core – Stars with masses M & 9 M do not burn as far as iron (although some may burn further than helium) – Recall: after helium flash, star settles onto horizontal branch (HB), with (convective) core helium burning and shell hydrogen burning – As helium burns to carbon and also oxygen, molecular weight increases, core contracts – Contraction accompanied by envelope expansion; star moves redward in HRD, toward Hayashi line – Eventually, core runs out of helium, burning stops; contraction continues until electron degeneracy pressure halts it – Helium burning resumes in shell • The Asymptotic Giant Branch – When star reaches Hayashi line, turns upwards with luminosity increase; this asymptotic giant branch is similar to red giant branch, but star now has He-burning shell, and is more luminous – Initially, star is on lower part of AGB — the early AGB (E-AGB) – On E-AGB, helium shell source is dominant source of energy; hydrogen shell is almost dormant – Because on Hayashi line, deep convection zone – When convection zone dips through hydrogen shell, mixes helium and nitrogen (produced by incomplete CNO-cycle burning) up to surface — second dredge-up (first dredge up occured on RGB) – As star moves to upper AGB, hydrogen shell reignites and becomes dormant – Helium burning shell become unstable, switching on and off periodically — thermal pulses. This part of AGB known as thermal pulse AGB (TP-AGB) – When helium shell switched off, helium accumulates from above due to outer
    [Show full text]