Star Clusters & Stellar Evolution

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1 Star Clusters & Stellar Evolution The final stages of stellar evolution Open and globular stellar clusters 22 February 2018 University of Rochester

2 Stellar Clusters & Stellar Evolution Late stellar evolution: the giant branch, horizontal branch, and asymptotic giant branch Evolution of high mass stars: the iron catastrophe Type II (core-collapse) supernovae Open and globular clusters as stellar clocks Reading: Kutner Ch & 13; Ryden Ch. 17.2, 18.2, & 18.4; Shu Ch. 8 & 9 Supernova progenitor simulation (Mosta et al. 2014). Colors indicate entropy. 22 February 2018 (UR) Astronomy 142 Spring / 38

3 Late stages of stellar evolution Recall from last time: in a late main sequence star hydrogen fusion in the center moves out in a concentric shell, leaving an isothermal He core. After the maximum mass of the core is exceeded, the star moves up and right in the H-R diagram, entering the red giant phase. The core collapses, causing a rise in ρ c and T c. The convection zone extends inward ( dredge-up ). Stellar radius dramatically increases due to the increase in radiation pressure from the interior, which now dominates support against gravitational collapse. The core temperature reaches 10 8 K and the triple-α process He 8 4 Be γ He 12 6 C γ begins helium fusion. The onset is very rapid in stars with M M, leading to a phenomenon called the helium flash. 22 February 2018 (UR) Astronomy 142 Spring / 38

4 The triple-α process The half-life of 8 Be is s and it decays back into two 4 He nuclei unless it interacts with a third 4 He nucleus to form 12 C. Some of the 12 C will then fuse with another 4 He to form 16 O. 22 February 2018 (UR) Astronomy 142 Spring / 38

5 Explanation of solar elemental abundances The triple-α process explains why C and O are so abundant (in a relative sense) compared to the other heavy elements. 22 February 2018 (UR) Astronomy 142 Spring / 38

6 A 1M star as a Red Giant Red giant structure: extremely dense He core and extremely tenuous and cool H envelope (ρ = 10 4 g cm 3 ). Images from ATNF. 22 February 2018 (UR) Astronomy 142 Spring / 38

7 Late stages of stellar evolution The horizontal branch is the phase after triple-α onset. Interior: core He burning, shell H burning. The core is on the He main sequence. Note the differences and similarities between a 1M, 5M, and 10M star. Diagram from ATNF. 22 February 2018 (UR) Astronomy 142 Spring / 38

8 A helium flash in the wild Sakurai s Object, a white dwarf that has undergone a He flash and returned to the red giant phase, shedding its outer envelope. From ESO/VLT/FORS. 22 February 2018 (UR) Astronomy 142 Spring / 38

9 Late stages of stellar evolution H-R diagram showing observations of stars in binary systems and globular clusters 22 February 2018 (UR) Astronomy 142 Spring / 38

10 Nomenclature: Spectral type The Harvard spectral type of a star is a classification based on the strength of absorption lines of several molecules, atoms, and ions (Cannon & Pickering 1901). Generally, types A M correspond to steady decreases in the strength of hydrogen lines. Type O is weaker still. Type P for planetary nebulae and Q for miscellany. N not used. In 1925 Cecilia Payne showed that the spectral type sequence OBAFGKM is a sequence of decreasing temperature from roughly 35000K to 3500K (Payne 1925). Numbers 0 9 add a further refinement of temperature within the spectral classes (0 = hottest, 9 = coolest). L, T, and Y recently added for cooler brown dwarfs. Examples: Vega is an A0 star; the Sun is a G2 star; Pollux is a K2 star; Betelgeuse is an M2 star. 22 February 2018 (UR) Astronomy 142 Spring / 38

11 Harvard spectral classification Class T e [K] M [M ] R [R ] L [L ] Frac. MS Pop. [%] O B A F G K M Table of spectral classes O through M for the Main Sequence. From wikipedia. See also Habets & Heintze (1981) and Ledrew (2001). 22 February 2018 (UR) Astronomy 142 Spring / 38

12 Nomenclature: Luminosity class At Yerkes Observatory in the 1940s, Morgan, Keenan, and Kellman added another dimension to classification with luminosity classes (Morgan et al. 1943). V main sequence or dwarf stars (the Sun) IV subgiants, brighter than V by 1 or 2 magnitudes for the same spectral type III normal giants, another few magnitudes brighter II, I bright giants and supergiants, even brighter (Antares, Betelgeuse) VI, VII subdwarfs and white dwarfs (Sirius B). Subdwarfs are metal-poor MS stars. White dwarfs are degenerate stellar remnants. Examples: Vega is an A0V star; the Sun is G2V, Pollux is K2III, and Betelgeuse is M2I. 22 February 2018 (UR) Astronomy 142 Spring / 38

13 Stellar classes on the observational H-R diagram Color-magnitude diagram showing luminosity classes (left) and spectral classes (right). Isochrones in log 10 (t/yr) are shown as dashed lines. 22 February 2018 (UR) Astronomy 142 Spring / 38

14 Spectral type, T e, and luminosity class Data from Schmidt-Kaler (1982) 22 February 2018 (UR) Astronomy 142 Spring / 38

15 After the horizontal branch: M < 2M Isothermal C-O core forms as He is exhausted. Not enough weight to overcome degeneracy pressure. Core cannot collapse and ignite C-O fusion (requires 500 MK!) H/He burning in outer core, ejection of rest of stellar envelope, forming a planetary nebula (few 1000 yr timescale). The core becomes a 0.6M C-O white dwarf with T K. It lasts forever unless it has a close stellar companion. 22 February 2018 (UR) Astronomy 142 Spring / 38

16 After the horizontal branch: M > 2M Asymptotic giant branch (ABG/supergiant) evolution Repeated core collapse after fuel exhaustion, up to Si fusion to produce Fe-peak elements. R and L steadily increase while T e decreases. Each successive fuel is exhausted faster than the last. For a 15M star (Woosley & Janka 2006): Fuel H He C Ne Time 10 7 yr 10 6 yr 10 3 yr 0.7 yr O, Mg Si,... Fe yr 18 dy 1 s 22 February 2018 (UR) Astronomy 142 Spring / 38

17 What happens after the nuclear fuel is exhausted? For most M > 2M stars: During burning of the heavier elements and radiative support of the stellar envelope, stars tend to be hydrodynamically unstable This leads to the loss of large fractions of the stellar mass. Oscillations: note that evolution takes stars across the instability strip, which is nearly vertical at T e 5000 K. Stellar winds can also remove significant amounts of material. These processes can keep a star s core mass below the SAC limit, so the final states of the star are like those of lower-mass objects: planetary nebula phase and white dwarf remnant. 22 February 2018 (UR) Astronomy 142 Spring / 38

18 What happens after the nuclear fuel is exhausted? For the most massive stars (M 8M ): Mass loss is insufficient to keep the core in the white dwarf phase; maximum mass of Fe white dwarf is 1.26M. Result: further collapse and neutronization. When the collapsing core reaches tens of km in size, neutron degeneracy pressure sets in, which can stop or slow the collapse. However, since the collapse has been from white-dwarf to neutron-star dimensions, infalling material from the stellar envelope is going very fast (a large fraction of c). Result: envelope rebounds off the stiff neutron-degenerate material and blows up the rest of the star. This is called a core collapse or Type II supernova. 22 February 2018 (UR) Astronomy 142 Spring / 38

19 Core-collapse (Type II) supernova Nomenclature: Type II or CCSN. Remnant is a NS or more rarely a BH, depending on core mass (M = M ). Image from NASA/CXC 22 February 2018 (UR) Astronomy 142 Spring / 38

20 Supernova types: Light curve classification Light curve: luminosity vs. time. Note the insane absolute magnitudes of these explosions! 22 February 2018 (UR) Astronomy 142 Spring / 38

21 Core collapse: t = 0 Si burning shell on core surface Core interior is Fe and Ni, which do not burn Start of collapse: electrons captured by Fe-peak nuclei, forming neutron-rich nuclei in the core Electron degeneracy pressure drops and the collapse accelerates Start of the core collapse (Janka et al. 2006). 22 February 2018 (UR) Astronomy 142 Spring / 38

22 Core collapse: t = 100 ms Capture of e, β decay, and photodisintegration of Fe-group nuclei to 4 2He occurs Electron density in the core goes down further and it loses energy; the collapse accelerates. A region 100 km across forms in the inner core with densities > g cm 3. Neutrinos become trapped and their density builds up in the inner core. Neutrino-trapping phase (Janka et al. 2006). 22 February 2018 (UR) Astronomy 142 Spring / 38

23 Core collapse: t = 110 ms The inner core collapses to 10 km and reaches nuclear densities (10 14 g cm 3 ). Neutron degeneracy pressure sets in and the inner core becomes extremely stiff. The outer core decelerates and bounces off the stiff inner core. A shock wave (discontinuity in velocity, density, etc.) is formed. The formation of the shock wave triggers the explosion. Shock wave formation (Janka et al. 2006). 22 February 2018 (UR) Astronomy 142 Spring / 38

24 Core collapse: t = 120 ms Shock moves into the outer core, losing energy to dissociation of nuclei. Material density decreases and neutrinos backed up behind the shock start escaping. A proto-neutron star (PNS) remnant forms. A huge pulse of neutrinos is emitted (break-out burst), carrying 99% of the gravitational binding energy of the remnant. Shock wave propagation (Janka et al. 2006). 22 February 2018 (UR) Astronomy 142 Spring / 38

25 Core collapse: t = s Hypothesis: neutrinos escaping the proto-neutron star heat up the shock and restore its lost energy, driving the explosion (Janka et al. 2006). The outer layers of the star mix. Neutron-rich nuclei act as seeds for the creation of very heavy elements (r-process). 22 February 2018 (UR) Astronomy 142 Spring / 38

26 SN1987A: A recent nearby supernova The last naked eye supernova occurred in the Large Magellanic Cloud, a Milky Way satellite galaxy. SN1987A (Feb. 23, 1987): a Type II Output luminosity temporarily exceeded luminosity of the entire Galaxy! First supernova for which we knew the progenitor star. Neutrino burst observed by three detectors (Kamiokande II, IMB, Baksan). 22 February 2018 (UR) Astronomy 142 Spring / 38

27 SN1987A neutrino light curve Only 25 neutrinos observed in total. Still enough to test models of CCSN. Left: The neutrino signal expressed in terms of the energy of the electrons detected after ν e scattering (Burrows 1988). Total ν count estimated to be 10 58, carrying erg (Pagliaroli et al. 2008). First detection of neutrinos from outside the solar system. 22 February 2018 (UR) Astronomy 142 Spring / 38

28 SN1987A after a few decades X-ray: NASA/CXC/U. Colorado/S. Zhekov et al. Optical: NASA/STScI/CfA/P. Challis. 22 February 2018 (UR) Astronomy 142 Spring / 38

29 A Type II supernova after 1000 years NASA/HST mosaic. 22 February 2018 (UR) Astronomy 142 Spring / 38

30 Simulating a Type II supernova Modeling a CCSN is a difficult task that requires significant computing resources. While we wait for the next Galactic SN, we can conduct interesting numerical experiments. Curie supercomputer, GENCI/TGCC-CEA. Movie: Example of a Standing Accretion Shock Instability (SASI), (MPI-Garching). 22 February 2018 (UR) Astronomy 142 Spring / 38

31 Observation of stellar evolution: Star clusters Stars tend to form in clusters at about the same age (within 1 2 Myr of each other) and are nearly the same distance from us. H-R color-magnitude diagrams for stars in the cluster make it very easy to infer the stars location on the main sequence. If we make an H-R diagram of a cluster we tend to see a turnoff point. The turnoff is caused by stars exhausting their H supply and moving along the horizontal branch. Higher-mass stars move off the MS first, so the turn off starts on the high-mass/high-t e end of the H-R diagram (upper left) and moves down as lower-mass stars age and exit the MS. In other words, the turnoff point acts like a clock, telling us the age of stars in the cluster with very good accuracy. 22 February 2018 (UR) Astronomy 142 Spring / 38

32 The turnoff point Left: H-R diagram for two open clusters, from wikimedia commons. Right: Cartoon of turnoff points for several clusters. From P. Eskridge. 22 February 2018 (UR) Astronomy 142 Spring / 38

33 Open and globular clusters Stars are classified into open (young) or globular (old) clusters. Properties of open clusters: Low density, irregular, lots of blue stars Low random velocities (few km/s) Hundreds of thousands of stars, not always gravitationally bound Archetypes: Pleiades (M45), Hyades, h and χ Persei Properties of globular clusters: High density, spherical, few blue stars Higher random velocities (tens of km/s) Millions of stars, gravitationally bound Archetypes: ω Centauri, M3, M13, 47 Tucanae 22 February 2018 (UR) Astronomy 142 Spring / 38

34 Open clusters: the Pleiades (M45) The Pleiades are about 135 pc away and are roughly 110 Myr old. Left: Optical image of the Pleiades cluster with the brightest stars highlighted. Right: H-R diagram of X-ray sources in the cluster (Stauffer et al. 1994). 22 February 2018 (UR) Astronomy 142 Spring / 38

35 Open clusters: the Hyades The Hyades are the closest open cluster (43 pc) and are 625 Myr old. Left: the Hyades (lower left) and Pleiades (upper right). Right: H-R diagram of X-ray sources in the cluster (Stauffer et al. 1994). 22 February 2018 (UR) Astronomy 142 Spring / 38

36 Open clusters: M67 M67 is 900 pc away and is the oldest known open cluster: 4 Gyr, the same age as the Sun. Left: M67, from N. Sharp and M. Hanna, NOAO/AURA/NSF. Right: H-R diagram using CCD data (Montgomery et al. 1993). 22 February 2018 (UR) Astronomy 142 Spring / 38

37 Globular clusters: M3 M3 is about 12 Gyr old, like all Galactic globular clusters, and it is 10.4 kpc away. RR Lyrae stars are not plotted (note gap in HB). Left: M3, from P. Challis (APOD). Right: H-R diagram from Montgomery et al February 2018 (UR) Astronomy 142 Spring / 38

38 More globular clusters: M13 & 47 Tuc Left: Globular cluster M13 taken with an 8 Schmidt-Cassegrain telescope (wikimedia commons). Right: 47 Tucanae (NASA/ESA/HST). 22 February 2018 (UR) Astronomy 142 Spring / 38