15 Post-main sequence stellar evolution

Transcription

1 15 Post-main sequence stellar evolution After the MS, the core becomes structured in concentric shells in which H, He, C/O, etc., burn at various times. This happens together with readjustments involving the expansion or contraction of the core or envelope, and the presence of extended convection zones. In order to describe the evolution post-ms, a division based on the stellar mass is usually made: Low-mass stars (LMs) 0.8 M <M<2M stars develop a degenerate helium core, leading to a long-lived red giant branch phase. The ignition of He is unstable, it is the so-called helium flash. Intermediate-mass stars (IMs) 2M <M 8M stars develop a helium core that remains non-degenerate. They ignite helium in stable conditions. After the central He-burning phase these stars form a C/O core that becomes degenerate. High-mass stars (HMs) M 8M stars ignite carbon and oxygen in non degenerate conditions. M 11M stars also ignite heavier elements until an iron core is formed, which finally collapses. Stars can loss a large fraction of its mass during the post-ms phases, in intense and discontinuous episodes; e.g., The Asymptotic Giant Branch (AGB) phase for IMs and LMs; and The Luminous Blue Variable (LBV) phase and the Wolf-Rayet (WR) phase for HMs. Univ. Nebraska [astro.unl.edu/classaction/outlines/stellarevolution/stellar_evol.html]

2 In LMs, the AGB phase is a necessary ingredient to link the star mass with the mass of the white dwarf (WD) left at the end, because of the very small mass-loss rates during the MS phase. The evolution of HMs, is quite different from IMS evolution, both because the C ignition occurs under non-degenerate conditions and because of the strong mass losses that stars can suffer. C ignition in the core requires T > K, and a minimum mass of C/O (from the central He-burning) of ~1.06 M, with mass uncertainties related to mixing (e.g., convective overshooting). This C/O-core mass corresponds to an initial mass of ~8 M. For M > 15 M stars, mass loss by stellar winds becomes important during all evolution phases, including the MS. Particularly, if M > 30 M, the mass-loss rate, Ṁ, may be so large that the timescale for mass loss (M/Ṁ) becomes smaller than the nuclear timescale. Therefore, mass loss has a very significant effect on the HMs evolution. The stellar wind mechanisms involved are in many cases not well understood yet, so Ṁ often is a source of uncertainties in the evolution of HMs. For example, up to now, it is unclear whether the continuous stellar wind is enough to explain the final masses before the star collapses into a neutron star (NS) or a black hole (BH), or if violent eruptions, like those of LBVs, are needed to match the final mass of the objects. However, we will not go in deep in this issue [

3 15.1 Post MS: Low- and Intermediate-mass stars Stars reach the H-burning shell phase when the hydrogen in the core is almost over. An isothermal He-core slowly grows, and when it reaches the Schönberg-Chandrasekhar limit, it starts contracting. The core collapse leads to a nonzero T in the He-core and the H- burning shell contracts: T and ρ increase and, hence, the energy generation rapidly intensifies despite the shell narrows. The extra energy is mostly directed to stellar envelope expansion, and the stellar luminosity decreases. The envelope expansion and the enhanced opacity of H favours convection in the external layers. The surface material is mixed with that of the interior chemically modified, hence, surface abundances result modified ( dredged up process ). For example, the 3 He/ 4 He ratio grows, while the lithium (if there is still some) is dragged inwards and quickly burned. Combustió en capa prima de l hidrogen Embolcall d hidrogen Nucli isoterm d heli (inert) The Red Giant Branch (RGB) phase. Envelope expansion leads to cooling of the outermost layers, and to a reduction of the opacity, so convection softens. At the same time, the core is still contracting and the subsequent increase of the luminosity of the H-burning shell proceed further; and this implies an increase in luminosity and radius: the star enters in the RGB phase (from 3 to 6). At some point in this phase (6), the conditions in the core become suitable for the 3α-process to begin. No mass loss considered

4 The helium burning core and the horizontal branch The trigger of the 3α-process (T > K) is very sensitive to T and causes the core expansion. This reduces the (still dominant) H-burning shell luminosity and also triggers core convection. Then the envelope contracts, which results in the growth of T ef, followed by a compression and re-brightening of the H-burning shell. The larger T ef makes the deep convection zone to retreat outwards. LMs: The core is strongly degenerated when the 3α-process starts. It absorbs energy until the degeneration is removed and only then the energy released can make work on the C/O core material: there is an explosive degenerate to non-degenerate core transition called helium flash. As helium burning proceeds, μ grows leading to core contraction, expanding and cooling again the envelope, thereby reducing the luminosity. This He/H burning core/shell phase is the horizontal branch (HB) of the HR diagram. During this phase stars develop instabilities in their outer envelopes, leading to pulsations. The Asymptotic Giant Branch (AGB). With the contraction of the core T increases; and when T K (and ρ ~10 6 gcm 3 ), a helium burning shell develops. This shell narrows and strengthens, forcing the material above it to expand and cool, temporarily turning off the H-burning shell. The contraction of the He-exhausted core makes neutrino cooling important increasing ρ and decreasing T, and electron degeneracy pressure becomes significant in the C/O core. Much of the energy of the He-burning shell goes to envelope expansion. With the decrease of T, convection deepens reaching between the H-rich outer layers and the He-rich zone, leading to a dredged up phase. At some point (i.e., 8 9), the H-burning shell re-ignites and the Heburning shell begins to turn on and off periodically, due to the He falling from the above H-burning shell (He-flashes can result, depending on the mass accreted). During such flashes the H- burning shell shuts down and a convective zone forms between both layers; thereby, another dredge up episode occurs.

5 15.2 Post-MS of a 1M star After the helium flash, the resulting central He-core will be the future white dwarf. The expanding outer layers will result in a planetary nebulae. Evolution of a 1 M star (initial abundances: X = 0.7, Z = 0.02) Top panel: evolutionary track of the 1 M star. Bottom panel: internal structure as a function of mass coordinate m. Gray areas are convective; lighter-gray areas are semi-convective. Red hatched regions: areas of energy generation where ε nuc > 5 L/M (dark red) and ε nuc > L/M (light red). The letters indicate the corresponding points in the evolutio track in the HR diagram.

6 15.3 HMs with M > 4 M : beyond the C/O core The He-burning shell produces C and O that accumulate in the core. It contracts and begins to degenerate, while neutrino losses make T slightly decrease. Ignoring mass loss, it would follow: The contracting C/O-core of stars < 4 M does not reach a temperature high enough to ignite the carbon combustion. M > 4 M, the Pauli principle would not stop contraction and the core would eventually reach a point in which a C/O core flash would take place, which, unlike the He-core flash, is truly explosive, forming a supernova. However, the very strong AGB mass loss prevents such supernova scenario, instead what happens is: The external layers are expelled forming a planetary nebulae and it only remains a white dwarf formed by a C/O-core surrounded by a thin shell of helium and hydrogen. The planetary nebulae cools down rapidly, in a few thousand years, while the C/O WD very slowly cools by radiation losses, moving towards the left bottom corner of the HR diagram. The main outcome of stellar evolution modelling is this type of theoretical HR-diagram Detailed evolutionary track of a Pop. I star of 5 M in the HR diagram. L is given in L and T ef in K. The evolution time between the labeled points is years (from Iben 1967a).

7 Evolution tracks for 1, 2, 3, 5, 7 and 10 M stars (quasi-solar composition: X = 0.7, Z = 0.02). Left panel: HR diagram; right panel: T c ρ c plane. Dotted lines in both panels show the location of the ZAMS. Dashed lines in the right panel show the borderlines between equation-of-state regions. The 1 M model is characteristic of LMs: the central core becomes degenerate soon after leaving the MS and helium is ignited in an unstable flash at the top of the red giant branch. When the degeneracy is eventually lifted, He-burning becomes stable and the star moves to the zero-age horizontal branch in the HR diagram, at log L ~1.8. The 2 M model is a borderline case that just undergoes a helium flash. The He-flash itself is not computed in these models, hence a gap appears in the tracks. The 5 M model is representative of IMs, undergoing quiet helium ignition and He-burning in a loop in the HR diagram. The appearance of the 7 M and 10 M models is qualitatively similar. However, at the end of its evolution the 10 M star undergoes carbon burning in the centre, while the cores of LMs become strongly degenerate.

8 Evolution of a 5M star (initial abundances: X = 0.7, Z = 0.02) Left: Evolution in the HR diagram. The letters indicate the corresponding points in both diagrams. Right: internal evolution; the panels show various quantities; top to bottom: (a) Contributions to the L from H-burning (red line), Heburning (blue) and gravitational energy release (orange; dashed parts show net absorption of gravitational energy). The black line is the surface luminosity. (b) Central mass fractions of various elements ( 1 H, 4 He, 12 C, 14 N and 16 O) as indicated. (c) Internal structure as a function of mass, m. A vertical line (green) corresponds to a model at a given time. Gray areas are convective; lighter-gray areas are semi-convective. The red hatched regions show areas of nuclear energy generation, where ε nuc > 10 L/M (dark red) and ε nuc > 2 L/M (light red) Post-MS evolution of high mass stars The surface properties of HMs cannot be precisely described without considering the mass losses they suffer in thepost-ms. It is not correct (M > 15 M?) to say that, for the H and He burning phases, such evolution can be qualitatively described in a similar way that of most massive IMs. Compare, for example, the track of a 15 M star shown in the upper diagram with the corresponding track in the bottom one, or with those ones shown in Fig of (OP, from Maeder and Meynet, 1987). Tracks based on the computations of Schaller et al. (A&A Suppl. 96, 269) of Geneva Observatory for 1, 5, 15 and 40 M stars. They include the effects of mass loss due to stellar wind. Blue line: MS; red lines: post-ms evolutionary tracks; points C mark the end of the MS and points F are during the He-core burning phase. (Portions of tracks between marks D and E are on the Hayashi line.)

9 We skip studying in more detail the effects of mass loss to address a more relevant issue in HMs evolution: what happens beyond the C/O-burning. Fortunately, for HMs there is a decoupling between the evolution of the stellar core and that of the external regions (which some how justifies the shortcut). Evolution beyond the carbon burning. Once an enough massive C/O-core is formed (> 1.06 M ), carbon ignition starts and the subsequent evolution of the core is a series of alternating nuclear burning and core contraction cycles in rapid succession. The overall evolutionary trend is an increase of T c and ρ c ; roughly, T c ρ 1/3 c (red arrowed line; as expected from hc). If T c K, the tracks slope down, towards higher ρ c and lower T c since the core cools due to strong neutrino emission. An onion-like shell structure develops within the star as heavier and heavier elements ignite consecutively. 25 M pre-supernova structure Evolution stages of a 25 M star (typical for M ) Exercise Compare the values in this table with those in table 12.1 from (OP). What do you think about the differences? Can you justify any of them in some way? Silicon burning produces a host of nuclei centered near the 56 Fe 26 peak. Nuclei reactions below iron are exothermic, but the heavier ones are endothermic. When T c ~ K, photo-disintegration of the heavy nuclei proceeds efficiently and, as is a strongly endothermic process, it reduces the pressure that holds the core against gravity.

10 Kippenhahn diagram of the evolution of a 15 M star. Cross-hatching: convective regions and nuclear burning intensity (blue shading) during central H and He burning (top panel) and during the late stages in the central 5 M (bottom panel). A series of convective burning cores and shells appear, due to C-burning (log t ~3), Ne-burning (log t ~0.6), O-burning (log t ~0) and Si-burning (log t ~ 2) (from Woosley et al 2002) Final composition profiles of a 15 M star, just before core collapse. Iron refers to the sum of neutron-rich nuclei of the iron group, especially 56 Fe. (from Woosley et al 2002)

11 By derivative work: G.A.SStellar_evolutionary_tracks.gif: Jesusmaiz- Stellar_evolutionary_tracks-en.PNG The HMs core collapse While protons start to capture free electrons, reducing the degeneracy pressure and producing neutrons and neutrinos. Finally, the Fe-core collapses inwards on; initially homologously, v r, near the center (ρ gcm -3 ) and supersonically from a certain radius later on), leaving the external layers outer than the silicon burning one suspended on vacuum. The collapse is halted because neutrons become degenerate, which trigger sound waves that turn into shocks further up. If the core is not too massive (? M ), the shock propagates up to the external layers of the star (prompt hydrodynamic explosion). For large iron cores, the shock stalls due to photo-disintegration cooling, forming an accretion shock. Photo-disintegration and electron capture yield vast amounts of neutrinos that deposit part of their energy behind the shock (delayed explosion process). The shock propagates through the star encountering the remainder shells with processed matter and the H-rich envelope.

12 The neutrino energy liberated in the explosion is E ν erg (L ν >> L opt ) The kinetic energy is E ex erg (<< E ν ). The explosive event becomes visible when the material becomes optically thin ( cm from the center of the star) The energy radiated in the visible range is E opt erg (1%), with a peak luminosity L opt erg s L. This is a Type II supernova (with hydrogen lines). The neutron-dominated core may become again stable due to neutron degeneracy, forming a neutron star (NS) if M NS < M TOV. M TOV is an upper bound to the mass of neutron stars (the Tolman- Oppenheimer-Volkoff limit), 1.5 M <M TOV 3.0 M. The uncertainty in the value reflects the fact that the equations of state for extremely dense matter are not well known. The mass of the pulsar PSR J , 2.00 ± 0.05 M, puts a lower bound on the TOV limit. M NS corresponds to original stars with masses, 15 M <M<20M. Exercise Watch (a) (b) (c) (d) (e) (f) Representation of the evolutionary stages from stellar core collapse through the onset of the supernova explosion to the neutrino-driven wind during the neutrino-cooling phase of the proto-neutron star. The panels display the dynamical conditions in their upper half, arrows represent velocity vectors. The nuclear composition as well as the nuclear and weak processes are indicated in the lower half of each panel. [Janka et al, Phys. Rep. 442, 38 (2007)]

13 For the most massive stars the core continues the collapse forming a black hole (BH). It is thought that lighter BH form after some matter falls back to the neutron-rich core of the star (proto-ns) pushing its collapse further. heavier BHs form at once, as the central region of the star involved in the explosion mechanism quickly collapses. Initial to final mass relation for stars (solar composition). Blue line: stellar mass after He-core burning, reduced by mass loss. For M > 30 M, the He-core is exposed as a WR star (two possibilities depending on the uncertain mass-loss rates). Red line: mass of the stellar remnant, resulting from AGB mass loss in the case of IMs, and ejection of the envelope in a core-collapse SN for HMs. Green areas: amount of mass ejected that has been processed by Heburning and further nuclear burning stages. (from Woosley et al. 2002).

14 15.5 Calibrating HR evolution tracks with clusters Theoretical isochrones from the Geneva model. Left: Theoretical HR diagram for the ensemble of the calculated models (Z = , with overshooting). No post-agb phases computed. Hatched areas indicate the slow phases of nuclear burning [Geneva stellar evolution tracks and isochrones database; Lejeune & Schaerer, A&A 366, 358 (2001)]. Right: Same as left but in the Mv (B-V)o plane. Examples of young and old stellar clusters Young open cluster (M45) Old globular cluster Pleiades. Age ~ yr. Formed by 1000 stars (brightest stars: hot B) m M = 5.60 (136 pc). Radius ~20 ly 47 Tucanae. Age yr. m M = (4.5 kpc). Radius ~60 ly

15 Colour magnitude diagram of the cluster NGC 1651 (z = ; LMC, 50 kpc). Comparison of the number of stars at selected evolutionary stages. Blue points: typical turn-off stars, the basis for the comparison. Red points: comparison sample of subgiant-branch stars. Isochrones for different ages are also shown [Li et al. Nature 516, 367 (2014)]. [see