DOCSLIB.ORG

Star Formation and Protostars 1 Star Formation and Protostars

Total Page：16

File Type：pdf, Size：1020Kb

Download full-text PDF Read full-text
Star Formation and Protostars 1 Star Formation and Protostars Stellar Objects: Star Formation and Protostars 1 Star Formation and Protostars 1 Preliminaries Objects on the way to become stars, but extract energy primarily from grav- itational contraction are called Young stellar objects here. Young stellar objects (YSO): the entire stellar system throughout all pre- main sequence (MS) evolutionary phases. • Protostar: the optically thick stellar core that forms during the adi- abatic contraction phase and grows during the accretion phase. The YSO is fully convective and evolves along the so-called Hayash track in the HR diagram. • PMS (pre-main sequence) star: the premature star that becomes visible once the natal envelope has been fully accreted and which still contracts towards the main sequence (MS), where the hydrogen burning domi- nates. In this evolutionary stage, the YSO has formed a radiative core, though still growing with time. Theoretically, the formation of a star may be divided into four stages (as illustrated in Fig. 2.1: (a) proto-star core formation; (b) the proto-star star builds up from the inside out, forming a disk around (core still contracts and is optically thick); (c) bipolar outflows; (d) the surrounding nebula swept away. The energy source is gravitational potential energy. While the total luminosity is comparable to the solar value, the proto-star stages have the 7 2 −1 −1 KH time scale [∼ 2 × 10 (M/M⊙) (L/L⊙) (R/R⊙) ]. They are also fully convective and is thus homogeneous chemically. Observationally, the proto-stars are classified into classes 0, 1, and 2, ac- cording to the ratio of infrared to optical, amount of molecular gas around, inflow/outflow, etc. Class 0 proto-stars are highly obscured and have short time scales (corresponding to the stage b?); a few are known. The class 1 and 2 proto-stars are already living partly on nuclear energy (c and d); but the total luminosity is still dominated by gravitational energy. The low-mass Stellar Objects: Star Formation and Protostars 2 YSO prototype is T Tauri. We still know little about high-mass YSOs, which evolve very fast and interact strongly with their environments. Signatures of YSOs: • variability on hours and days due to temperature irregularities on both the stellar surface and disk • emission lines from the disk and/or outflow • more inferred luminosity due to dust emission • high level of magnetic field triggered activities (flares, spots, corona ejection, etc) due to fast rotation and convection • Strong X-ray emission from hot corona. 2 The birth of stars From the virial theorem, 2E = −Ω, we have 3kT M M GM = r dM (1) µm r r A Z0 for the hydrostatic equilibrium of a gas sphere with a total mass M. Assum- ing that the density is constant, the right side of the equation is 3/5(GM 2/R). If the left side is smaller than the right side, the cloud would collapse. For the given chemical composition, ρ and T , this criterion gives the minimum mass (called Jeans mass) of the cloud to undergo a gravitational collapse: 3 1/2 5kT 3/2 M > M ≡ . (2) J 4πρ Gµm A For typical temperatures and densities of large molecular clouds, MJ ∼ 5 −1/2 10 M⊙ with a collapse time scale of tff ≈ (Gρ) . Such mass clouds may be formed in spiral density waves and other density perturbations (e.g., caused by the expansion of a supernova remnant or superbubble). Stellar Objects: Star Formation and Protostars 3 What exactly happens during the collapse depends very much on the temper- ature evolution of the cloud. Initially, the cooling processes (due to molecular and dust radiation) are very efficient. If the cooling time scale tcool is much −1/2 shorter than tff , the collapse is approximately isothermal. As MJ ∝ ρ decreases, inhomogeneities with mass larger than the actual MJ will collapse by themselves with their local tff , different from the initial tff of the whole cloud. This fragmentation process will continue as long as the local tcol is shorter than the local tff , producing increasingly smaller collapsing subunits. Eventually the density of subunits becomes so large that they become opti- cally thick and the evolution becomes adiabatic (i.e., T ∝ ρ2/3 for an ideal 1/2 gas), then MJ ∝ ρ . As the density has to increase, the evolution will al- ways reach a point when M = MJ , when we assume that a star is born. From this moment on the cloud would start to evolve in hydrostatic equilibrium. This way a giant molecular cloud can form a group of stars with their mass distribution being determined by the fragmentation process. The process depends on the physical and chemical properties of the cloud (ambient pres- sure, magnetic field, rotation, composition, dust fraction, stellar feedback, etc.). Much of the process is yet to be understood. We cannot yet theoret- ically determine the distribution of the initial mass function (IMF) of stars. The classic expression for the IMF, determined empirically is the so-called Salpeter’s law dn/dM = CM −x (3) where x = 2.35 for M/M⊙ ≥ 0.5 and x = 1.3 for 0.1 ≤ M/M⊙ ≤ 0.5 in the solar neighborhood. 3 Evolution of YSOs Still incomplete!!! The actually formation and evolution of a YSO is much more complicated, but roughly can be divided into the following stages: During the collapse the density increases inwards. The optically thick phase is reached first in the central region, which leads to the formation of a more-or- less hydrostatic core with free falling gas surrounding it. The energy released Stellar Objects: Star Formation and Protostars 4 by the core (now obeying the virial theorem) is absorbed by the envelope and radiated away as infrared radiation. Because of the heavy obscuration by the surrounding dusty gas, stars in this stage cannot be directly observed in optical and probably in near-IR. When the accretion is essentially completed, the YSO increases the temper- ature as the core continue to contract. The steady increase of the central temperature causes the dissociation of the H2, then the ionization of H, and the first and second ionization of He. The sum of the energy involved in all these processes has to be at most equal to the energy available to the star through the virial theorem. The simple estimate gives the maximum initial i radius Rmax of a YSO has to be Ri M max ≈ 50 (4) R⊙ M⊙ Therefore, the luminosity can be very large. 3.1 Evolution along the Hayashi Track Starting from the surface inward, we have encountered the simple atmosphere model in the heating transfer section. The temperature and pressure at the g ≈ 2 s 2 photosphere are T (τP )= Teff and P (τP ) 3 , where gs = GM/R is the κP local gravity and κP is the opacity κ at the photosphere. Note that we have assumed that convection plays no role in hear transport between the true and photosphere surface, consistent with our notion of a radiating, static, and visible surface (although this is not true for the sun). Next we consider the stellar “envelope”, which consists of the portion of a star that starts at the photosphere, and continues inward, but contains negligible mass, has no thermonuclear or gravitational energy sources, and is in hydrostatic equilibrium. We want to see how deep the radiative envelope may be until the convection takes over for heat transfer. In such a radiative envelope, ∇ = ∇rad with dlnT 3κL P ∇ = = . (5) dlnP 16πacGM T 4 Stellar Objects: Star Formation and Protostars 5 Have the opacity written in the interpolation form for an idea gas, κ = n −n−s κgP T , the above equation can be re-written as 16πacGM P ndP = T n+s+3dT. (6) 3κgL n+1 n+1 ′ n+1 n+s+4 n+s+4 P − PP =(K ) (T − TP ) (7) where 1 16πacGM 1/(n+1) K′ = (8) 1+ n 3κ L eff g where neff =(s + 3)/(n + 1). It is easy to prove that this can be written as 1 T n+s+4 1 ∇(r)= + eff ∇ − (9) 1+ n T (r) P 1+ n eff eff where ∇P is ∇ evaluated at the photosphere: 3L P ∇ P P = 4 = 1/8 (10) 16πacGMκP TP 2 2 4 where Pp = 2gs/3κP , gs = GM/R , and L = 4πR σReff are used. − For H opacity, n = 1/2 and s = −9, we have neff = −4 and 1 11 T (r) 9/2 ∇(r)= − + (11) 3 24 T eff Note that since temperature increases with depth, so does ∇. Eventually, when ∇ > ∇ad, the stellar material becomes convective. For simplicity, we assume that the convection is adiabatic. Thus at depths deeper than the critical depth, ∇ = ∇ad = 0.4 for ideal gas. The transition to convection 2/9 occurs at Tf = (8/5) Teff = 1.11Teff . The corresponding pressure Pf at the top of the convective interiors can be found P n+1 1 1 T (r) n+s+4 =1+ − 1 (12) P 1+ n ∇ T P eff P " eff # 2/3 which yield Pf = 2 PP . Stellar Objects: Star Formation and Protostars 6 For the assumed convection, the implying polytrope of index is 3/2 and ′ 5/2 P = Kn=3/2T (13) (K′ here is different from the previously defined??). ′ For a completely convective star, Kn=3/2 can be related to the the mass and radius of the star as defined by the E-solution polytrope 1 K′ ∝ (14) n=3/2 µ5/2M 1/2R3/2 ′ 5/2 2 4 Using Pf = Kn=3/2Tf and L = 4πσR Teff to replace R, we get 13/51 7/51 1/102 Teff ∝ µ (M/M⊙) (L/L⊙) (15) Therefore, such a completely convective star has a nearly constant effective temperature, independent of the luminosity.
Load more

Recommended publications
  • Revisiting the Pre-Main-Sequence Evolution of Stars I. Importance of Accretion Efficiency and Deuterium Abundance ?
    Astronomy & Astrophysics manuscript no. Kunitomo_etal c ESO 2018 March 22, 2018 Revisiting the pre-main-sequence evolution of stars I. Importance of accretion efficiency and deuterium abundance ? Masanobu Kunitomo1, Tristan Guillot2, Taku Takeuchi,3,?? and Shigeru Ida4 1 Department of Physics, Nagoya University, Furo-cho, Chikusa-ku, Nagoya, Aichi 464-8602, Japan e-mail: [email protected] 2 Université de Nice-Sophia Antipolis, Observatoire de la Côte d’Azur, CNRS UMR 7293, 06304 Nice CEDEX 04, France 3 Department of Earth and Planetary Sciences, Tokyo Institute of Technology, 2-12-1 Ookayama, Meguro-ku, Tokyo 152-8551, Japan 4 Earth-Life Science Institute, Tokyo Institute of Technology, 2-12-1 Ookayama, Meguro-ku, Tokyo 152-8551, Japan Received 5 February 2016 / Accepted 6 December 2016 ABSTRACT Context. Protostars grow from the first formation of a small seed and subsequent accretion of material. Recent theoretical work has shown that the pre-main-sequence (PMS) evolution of stars is much more complex than previously envisioned. Instead of the traditional steady, one-dimensional solution, accretion may be episodic and not necessarily symmetrical, thereby affecting the energy deposited inside the star and its interior structure. Aims. Given this new framework, we want to understand what controls the evolution of accreting stars. Methods. We use the MESA stellar evolution code with various sets of conditions. In particular, we account for the (unknown) efficiency of accretion in burying gravitational energy into the protostar through a parameter, ξ, and we vary the amount of deuterium present. Results. We confirm the findings of previous works that, in terms of evolutionary tracks on the Hertzsprung-Russell (H-R) diagram, the evolution changes significantly with the amount of energy that is lost during accretion.
    [Show full text]
  • Stellar Evolution
    Stellar Astrophysics: Stellar Evolution 1 Stellar Evolution Update date: December 14, 2010 With the understanding of the basic physical processes in stars, we now proceed to study their evolution. In particular, we will focus on discussing how such processes are related to key characteristics seen in the HRD. 1 Star Formation From the virial theorem, 2E = −Ω, we have Z M 3kT M GMr = dMr (1) µmA 0 r for the hydrostatic equilibrium of a gas sphere with a total mass M. Assuming that the density is constant, the right side of the equation is 3=5(GM 2=R). If the left side is smaller than the right side, the cloud would collapse. For the given chemical composition, ρ and T , this criterion gives the minimum mass (called Jeans mass) of the cloud to undergo a gravitational collapse: 3 1=2 5kT 3=2 M > MJ ≡ : (2) 4πρ GµmA 5 For typical temperatures and densities of large molecular clouds, MJ ∼ 10 M with −1=2 a collapse time scale of tff ≈ (Gρ) . Such mass clouds may be formed in spiral density waves and other density perturbations (e.g., caused by the expansion of a supernova remnant or superbubble). What exactly happens during the collapse depends very much on the temperature evolution of the cloud. Initially, the cooling processes (due to molecular and dust radiation) are very efficient. If the cooling time scale tcool is much shorter than tff , −1=2 the collapse is approximately isothermal. As MJ / ρ decreases, inhomogeneities with mass larger than the actual MJ will collapse by themselves with their local tff , different from the initial tff of the whole cloud.
    [Show full text]
  • Phys 321: Lecture 7 Stellar EvoluOn
    Phys 321: Lecture 7 Stellar Evolu>on Prof. Bin Chen, Tiernan Hall 101, [email protected] Stellar Evoluon • Formaon of protostars (covered in Phys 320; briefly reviewed here) • Pre-main-sequence evolu>on (this lecture) • Evolu>on on the main sequence (this lecture) • Post-main-sequence evolu>on (this lecture) • Stellar death (next lecture) The Interstellar Medium and Star Formation the cloud’s internal kinetic energy, given by The Interstellar Medium and Star Formation the cloud’s internal kinetic energy, given by 3 K NkT, 3 = 2 The Interstellar MediumK andNkT, Star Formation = 2 where N is the total number2 ofTHE particles. FORMATIONwhere ButNNisis the just total OF number PROTOSTARS of particles. But N is just M Mc N c , Our understandingN , of stellar evolution has= µm developedH significantly since the 1960s, reaching = µmH the pointwhere whereµ is the much mean of molecular the life weight. history Now, of by a the star virial is theorem, well determined. the condition for This collapse success has been where µ is the mean molecular weight.due to advances Now,(2K< byU the in) becomes observational virial theorem, techniques, the condition improvements for collapse in our knowledge of the physical | | (2K< U ) becomes processes important in stars, and increases in computational2 power. In the remainder of this | | 3MckT 3 GMc chapter, we will present an overview of< the lives. of stars, leaving de(12)tailed discussions 2 µmH 5 Rc of s3oMmeckTspecia3l pGMhasces of evolution until later, specifically stellar pulsation, supernovae, The radius< may be replaced. by using the initial mass density(12) of the cloud, ρ , assumed here µmH 5 Rc 0 and comtop beac constantt objec throughoutts (stellar theco cloud,rpses).
    [Show full text]
  • Astronomy General Information
    ASTRONOMY GENERAL INFORMATION HERTZSPRUNG-RUSSELL (H-R) DIAGRAMS -A scatter graph of stars showing the relationship between the stars’ absolute magnitude or luminosities versus their spectral types or classifications and effective temperatures. -Can be used to measure distance to a star cluster by comparing apparent magnitude of stars with abs. magnitudes of stars with known distances (AKA model stars). Observed group plotted and then overlapped via shift in vertical direction. Difference in magnitude bridge equals distance modulus. Known as Spectroscopic Parallax. SPECTRA HARVARD SPECTRAL CLASSIFICATION (1-D) -Groups stars by surface atmospheric temp. Used in H-R diag. vs. Luminosity/Abs. Mag. Class* Color Descr. Actual Color Mass (M☉) Radius(R☉) Lumin.(L☉) O Blue Blue B Blue-white Deep B-W 2.1-16 1.8-6.6 25-30,000 A White Blue-white 1.4-2.1 1.4-1.8 5-25 F Yellow-white White 1.04-1.4 1.15-1.4 1.5-5 G Yellow Yellowish-W 0.8-1.04 0.96-1.15 0.6-1.5 K Orange Pale Y-O 0.45-0.8 0.7-0.96 0.08-0.6 M Red Lt. Orange-Red 0.08-0.45 *Very weak stars of classes L, T, and Y are not included. -Classes are further divided by Arabic numerals (0-9), and then even further by half subtypes. The lower the number, the hotter (e.g. A0 is hotter than an A7 star) YERKES/MK SPECTRAL CLASSIFICATION (2-D!) -Groups stars based on both temperature and luminosity based on spectral lines.
    [Show full text]
  • Understanding the Luminosity and Surface Temperature of Stars
    Understanding the Luminosity and Surface Temperature of Stars 1. The Hertsprung-Russell diagram Astronomers like to plot the evolution of stars on what is known as the Hertsprung- Russell diagram. Modern versions of this diagram plot the luminosity of the star against its surface temperature, so to understand how a star moves on this diagram as it ages, we need to understand the physics that sets the luminosity and surface temperature of stars. Surprisingly, these can be understood at an approximate level using only basic physics, over most of the evolution of both low-mass and high-mass stars. 2. The role of heat loss Stellar evolution is fundamentally the story of what happens to a star by virtue of it losing an enormous amount of heat. By the virial theorem, the total kinetic energy in a star equals half its total gravitational potential energy (in absolute value), as long as the gas in the star remains nonrelativistic. This implies that when a gravitationally bound object loses heat by shining starlight, it always contracts to a state of higher internal kinetic energy. For a nonrelativistic gas, the pressure P is related to the kinetic energy density KE/V by 2 KE P = , (1) 3 V we can apply the nonrelativistic virial theorem KE 1 GMm = i (2) N 2 R where mi is the average ion mass (about the mass of a proton, usually) to assert that the characteristic pressure over the interior (i.e., the core pressure Pc) approximates GMρ 2 P ∼ ∝ , (3) c R R4 where ρ ∝ M/R3 is the mass density.
    [Show full text]
  • Proto and Pre-Main Sequence Stellar Evolution in a Molecular Cloud Environment
    MNRAS 000,1{20 (2017) Preprint 1 February 2018 Compiled using MNRAS LATEX style file v3.0 Explaining the luminosity spread in young clusters: proto and pre-main sequence stellar evolution in a molecular cloud environment Sigurd S. Jensen? & Troels Haugbølley Centre for Star and Planet Formation, Niels Bohr Institute and Natural History Museum of Denmark, University of Copenhagen, Øster Voldgade 5-7, DK-1350 Copenhagen K, Denmark Accepted 2017 October 31. Received October 30; in original form 2017 June 27 ABSTRACT Hertzsprung-Russell diagrams of star forming regions show a large luminosity spread. This is incompatible with well-defined isochrones based on classic non-accreting pro- tostellar evolution models. Protostars do not evolve in isolation of their environment, but grow through accretion of gas. In addition, while an age can be defined for a star forming region, the ages of individual stars in the region will vary. We show how the combined effect of a protostellar age spread, a consequence of sustained star formation in the molecular cloud, and time-varying protostellar accretion for individual proto- stars can explain the observed luminosity spread. We use a global MHD simulation including a sub-scale sink particle model of a star forming region to follow the accre- tion process of each star. The accretion profiles are used to compute stellar evolution models for each star, incorporating a model of how the accretion energy is distributed to the disk, radiated away at the accretion shock, or incorporated into the outer lay- ers of the protostar. Using a modelled cluster age of 5 Myr we naturally reproduce the luminosity spread and find good agreement with observations of the Collinder 69 cluster, and the Orion Nebular Cluster.
    [Show full text]
  • 8.901 Lecture Notes Astrophysics I, Spring 2019
    8.901 Lecture Notes Astrophysics I, Spring 2019 I.J.M. Crossfield (with S. Hughes and E. Mills)* MIT 6th February, 2019 – 15th May, 2019 Contents 1 Introduction to Astronomy and Astrophysics 6 2 The Two-Body Problem and Kepler’s Laws 10 3 The Two-Body Problem, Continued 14 4 Binary Systems 21 4.1 Empirical Facts about binaries................... 21 4.2 Parameterization of Binary Orbits................. 21 4.3 Binary Observations......................... 22 5 Gravitational Waves 25 5.1 Gravitational Radiation........................ 27 5.2 Practical Effects............................ 28 6 Radiation 30 6.1 Radiation from Space......................... 30 6.2 Conservation of Specific Intensity................. 33 6.3 Blackbody Radiation......................... 36 6.4 Radiation, Luminosity, and Temperature............. 37 7 Radiative Transfer 38 7.1 The Equation of Radiative Transfer................. 38 7.2 Solutions to the Radiative Transfer Equation........... 40 7.3 Kirchhoff’s Laws........................... 41 8 Stellar Classification, Spectra, and Some Thermodynamics 44 8.1 Classification.............................. 44 8.2 Thermodynamic Equilibrium.................... 46 8.3 Local Thermodynamic Equilibrium................ 47 8.4 Stellar Lines and Atomic Populations............... 48 *[email protected] 1 Contents 8.5 The Saha Equation.......................... 48 9 Stellar Atmospheres 54 9.1 The Plane-parallel Approximation................. 54 9.2 Gray Atmosphere........................... 56 9.3 The Eddington Approximation................... 59 9.4 Frequency-Dependent Quantities.................. 61 9.5 Opacities................................ 62 10 Timescales in Stellar Interiors 67 10.1 Photon collisions with matter.................... 67 10.2 Gravity and the free-fall timescale................. 68 10.3 The sound-crossing time....................... 71 10.4 Radiation transport.......................... 72 10.5 Thermal (Kelvin-Helmholtz) timescale............... 72 10.6 Nuclear timescale..........................
    [Show full text]
  • 21 Stellar Evolution: the Core
    21StellarEvolution:TheCore 21.1 Useful References • Prialnik,2nd ed., Ch.7 21.2 Introduction It’sfinally time to combine much of what we’ve introduced in the past weeks to address the full narrative of stellar evolution. This sub-field of astrophysics traces the changes to stellar composition and structure on nuclear burning timescales (which can range from Gyr to seconds). We’ll start with a fairly schematic overview –first of the core, where the action is, then zoom out to the view from the surface, where the physics in the core manifest themselves as observables via the equations of stellar structure (Sec. 15.5). Then in Sec. 22 we’ll see how stellar evolution behaves in the rest of the star, and its observa- tional consequences. Critical in our discussion will be the density-temperature plane. For a given composition (which anyway doesn’t vary too widely for many stars), a star’s evolutionary state can be entirely determined solely by the conditions in the core,T c andρ c. Fig. 39 introduces this plane, in which (as we will see) each particular stellar mass traces out a characteristic, parametric curve.8 21.3 The Core There are several key ingredients for our “core view.” These include: 8We will see that it costs us very little to populate Fig. 39; that is, there’s no need to stress over theT-ρ price. Figure 42: The plane of stellar density and temperature. See text for details. 143 21.StellarEvolution:TheCore • Equation of State. This could be an ideal gas (P∝ρT), a non-relativistic degenerate gas (Eq.
    [Show full text]
  • Miljenko Čemeljić, CAMK, Warsaw & ASIAA, Taiwan
    arXiv:2012.08550 Miljenko Čemeljić, CAMK, Warsaw & ASIAA, Taiwan Miljenko Čemeljić, CAMK Journal Club, March, 2021, Warsaw Outline -Hypergiants? -VY Canis Majoris -Possible implications on the explanation of the recent dimming of Betelgeuse. Miljenko Čemeljić, CAMK Journal Club, March, 2021, Warsaw Hypergiants -Hypergiants: how giant are they? -At the beginning of the article is written that if VY Cma, with radius about 1420 R_sun, would be in place of our Sun, it would reach to Jupiter. Jupiter? I remembered “giants, e.g. Antares, would reach to Mars”, this one ... Jupiter?! Other supergiants are Westerlund 1-26, WOH G64, with about 1200 and 1500 R_sun, and Stephenson 2-18 with about 2500 R_sun-it would reach to Saturn! Light travelling around the equator of VY CMa star would take 6 hours, compared to 14.5 seconds for the Sun. It is not at all that “obvious” for me that gaseous spheres can be THAT big. I did not remember it from the HR diagrams I saw. “Space is big. You just won't believe how vastly, hugely, mind-bogglingly big it is. I mean, you may think it's a long way down the road to the chemist's, but that's just peanuts to space.” ― Douglas Adams, “The Hitchhiker's Guide to the Galaxy” Supergiants in HR diagram Radius of the Sun is about 695,700 km --> diameter is of the order of 1 million km. Distance Earth-Moon is 384 000 km. Hypergiants? Supergiants in HR diagram About 90 percent of the stars in the universe exist along Main sequence line at one time in their lives, when they still fuse hydrogen to helium.
    [Show full text]
  • 18 Stellar Evolution
    18StellarEvolution:TheRest of thePicture We’ll now consider stellar evolution from the external perspective. All the interior physics is now translated through millions of km of starstuff; Plato’s cave doesn’t have anything on this. As in Fig. 36 we will still observe stars move through a2D space as their evolution goes on. But now, instead ofT c and ρc (which we can only infer and never directly measure) our new coordinates will be the external observables luminosity and effective temperature,L and Teff. In truth even these quantities rely on inference at some level; so while the astrophysicist thinks of (L,T eff) the observing astronomer will often think in terms of absolute magnitude (a proxy forL) and photometric colors (a proxy forT eff). Yes, we’ve returned once again to the Hertzprung-Russell diagram first introduced in Sec.8. 18.1 Stages of Protostellar Evolution: The Narrative The process of star formation is typically divided into a number of separate stages (which may or may not have distinct boundaries). Here, we will con- sider eight stages, beginning with the initial collapse (which we have already touched on) and ending when a star reaches what is known as the ‘main se- Figure 37: 131 18.StellarEvolution:TheRest of thePicture quence’, or a stable state of central nuclear burning of Hydrogen, in which a star will remain for the majority of its life. Note that understanding all of the physics that go into these stages of evolution requires some knowledge of top- ics we have previously discussed: adiabatic processes (Sec.
    [Show full text]
  • The Red Supergiants in the Supermassive Stellar Cluster Westerlund 1 (As Supergigantes Vermelhas No Aglomerado Estelar Supermassivo Westerlund 1)
    Universidade de S~aoPaulo Instituto de Astronomia, Geof´ısicae Ci^enciasAtmosf´ericas Departamento de Astronomia Aura de las Estrellas Ram´ırezAr´evalo The Red Supergiants in the Supermassive Stellar Cluster Westerlund 1 (As Supergigantes Vermelhas no Aglomerado Estelar Supermassivo Westerlund 1) S~aoPaulo 2018 Aura de las Estrellas Ram´ırezAr´evalo The Red Supergiants in the Supermassive Stellar Cluster Westerlund 1 (As Supergigantes Vermelhas no Aglomerado Estelar Supermassivo Westerlund 1) Disserta¸c~aoapresentada ao Departamento de Astronomia do Instituto de Astronomia, Geof´ısica e Ci^enciasAtmosf´ericasda Universidade de S~aoPaulo como requisito parcial para a ob- ten¸c~aodo t´ıtulode Mestre em Ci^encias. Area´ de Concentra¸c~ao:Astronomia Orientador: Prof. Dr. Augusto Damineli Neto Vers~aoCorrigida. O original encontra- se dispon´ıvel na Unidade. S~aoPaulo 2018 To my parents, Eustorgio and Esmeralda, to whom I owe everything. Acknowledgements This dissertation was possible thanks to the support of many people, to whom I owe not only this manuscript and the work behind it, but many years of personal and professional growth. In the first place I thank my parents, Eustorgio and Esmeralda, for all the unconditional love and the sacrifices they have made to make me the woman I am today, for never leaving me alone no matter how far we are, and for all the support they have given me throughout the years, that has allowed me to fulfill my dreams. This accomplishment is theirs. Pap´a y mam´a,este logro es de ustedes. I thank Alex, my companion during this phase, for his love, patience and priceless support especially in those difficult times, when he would cheer me up.
    [Show full text]
  • From the Hayashi Track to the Main Sequence: During Collapse, The
    12 From the Hayashi Track to the Main Sequence: During collapse, the Virial Theorem says that half the energy goes into radiation and the other half into heat (if n fixed): 1 dE n 3 GM 2 dR L = = − ; − 2 dt 5 n 2R2 dt − dL 1 d2E n 3 GM 2 2 dR 2 d2R = = − + : dt − 2 dt2 5 n 2R2 −R dt dt2 − " # 2 4 Constancy of Teff with L = 4πσR Teff implies dL=L = 2dR=R and d2R=dt2 = 4(dR=dt)2=R. This leads to a solution 1=3 6LoRot − R = Ro 1 + ; GM 2 where Ro and Lo refer to the initial radius and luminosity. Continued collapse to the main sequence involves a change in the polytropic structure from n = 3=2 to n = 3 (the standard model). If we neglect this, however, and use the hydrostatic condition d2I=dt2 = 0 where I MR2 is the moment of iner- tia, one finds that Rd2R=dt2 =/ (dR=dt)2. This implies that dL=L = 3dR=R, which is now p−ositive, and to a solution − 4RoLot R = Ro 1 2 : r − GM This also implies that d ln Teff =d ln R = 5=4 and d ln L=d ln Teff = 12=5. This phase of collapse is known−as the Henyey phase, and results in an increase in L and Teff with stellar shrink- age. The stellar radius, for a one solar mass star, must shrink from Rmin to R , the effective temperature rises from about 12=5 3500 K to 5500 K, so Lmin (3500=5500) 3 L and 4=5 ' ' Rmin (5500=3500) 0:7 R .
    [Show full text]