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- - - - -possible solution to G-dwarf problem is that gas from which disc was made already had some metals before it arrived in solar vicinity
        -heavy elements produced by earliest stars could have mixed with gas that eventually formed disc, to "pre-enrich" it
        -in that case, should expect all stars to have metal abundance above some min value
        -also possible that SF near Sun started before gaseous raw material had been fully assembled
        -in that case, 1st stars would enrich only small amount of gas to moderately high metal abundance
        -subsequent inflow of fresh metal-deficient gas would dilute that material, preventing abundance from rising as fast as closed-box model predicts
        -incompletel mixing could explain large dispersion in stellar abundance at given age
        -since gas in outer parts of galaxies is poorer in heavy elements, a slow inward flow, perhaps caused by energy loss in passing through shocks in spiral arms, would dilute metals in local disc
        -long-lived stars formed at early times should also return metal-poor gas as they age
        -if enough gas released, fraction of metals in newly made stars might even decline with time
    - - -to predict metal abundance of stars in galaxy, first need to estimate birthrate of stars at each mass
        -can then calculate how much of each element is relased back to interstellar gas over time
        -in MW, long-lived stars born soon after disc formed are now returning metal-poor gas, so abundance of heavy elements in gas is growing only slowly
        -chemical composition of new stars also depends on how gas moves within galaxy, since clock of metal enrichment runs more slowly in outer regions
        -require chemodynamical models, under development
- - - - -AGB stars known to lose mass at a rapid rate and Teffs quite cool; as result, dust grains form in the expelled matter
        -since silicate grains tend to form in environment rich in O and graphite grains will form in C-rich environment, composition of ISM may be related to relative numbers of C- and O-rich stars
        -observations of UV extinction curves in MW and LMC/SMC supports idea that mass loss from these stars does help enrich ISM
        -as stars with initial masses below 8Msun continue to evolve up the AGB, He-burning shell converts more and more of the He into C and then into O, increasing mass of C-O core
        -at the same time, core continues to contract slowly, causing central density to increase
        -depending on star's mass, neutrino energy losses may decrease central temperature somewhat during this phase
        -densities in core become large enough that electron degeneracy pressure begins to dominate
        -mass loss prevents catastrophic core collapse, and stars experience additional nucloesynthesis in cores, leading to core compositions of O, Ne, Mg and masses below Chandrasekhar limit
        -in latest stages of evolution on AGB, superwind develops, stars shrouded in optically thick dust clouds that radiate energy primarily in IR part of EM spectrum
        -during final phase of mass loss, remainder of envelope is expelled
  - - - -in each case, star ends life in SN explosion, see scheme for different masses on page 537
        -most massive stars never evolve to red supergiant portion of HR diagram
        -although very massive stars are extremely rare (only 1 100Msun str exists for every 1 million 1Msun star), they play a major role in dynamics and chemical evolution of ISM
        -tremendous amont of KE deposited in ISM through stellar winds of massive stars has significant impact on kinematics of ISM
        -when very massive stars form, they have ability to quench SF in their regions
        -UV light from massive stars also ionizes gas clouds in their region
        -highly enriched gases of massive stellar winds increase the metal content of the ISm, resulting in formation of increasingly metal-rich stars
      - -Type I SNe are those that don't exhibit any H lines in spectra while TYpe II have strong H lines
        -Type Ia are those that show strong Si II line at 615nm, Type Ib or Ic based on presence or absence of strong He lines
        -lack of H lines in Type I SNe indicates that stars involved have been stripped of their H envelopes
        -differences in spectral signatures between Ia, Ib, Ic indicate that different physical mechanisms at work
        -reflected in different environments observed for these outbursts: Ia found in all types of galaxeis, including ellipticals that show vry little evidence of recent SF; Ib and Ic only seen in spirals near sites of recent SF (HII regions), implying that short-lived massive stars are probably involved with Types Ib and Ic, but not with Type Ia
        -Type II SNE are characterized by rapid rise in luminosity, reaching max brightness typically 1.5 mag dimmer than Type Ia's
        -peak output is followed by steady decrease, dropping 6-8mags in a year
        -spectra exhibit lines associated with H and heavier elements
        -light curves of Type II can be classified as Type II-P (plateau) or Type II-L (linear)
        -Type II's are related to Type Ib's and Ic's
        
        -sheer amount of energy released in SN event is staggering: typical Type II releases 10^46J of energy, with 1% of that appearing as KE of ejected material and less than 0.01% being released as photons; remainder of energy radiated in form of neutrinos; similar values for Ib and Ic
        -H is converted into He on MS, followed by He burning leading to a CO core, but the very high temperature in core of massive star means C and O can burn as well; end result is that rather than star ending life through formation of planetary nebula, catastrophic SN explosure occurs instead
        -II, Ib, and Ic known as CCSNe
        -inner core rebounds after collapse (outer layers suspended), sneding pressure waves outward into infalling material from outer core
        -when velocity of pressure waves reach sound speed, they build into a shock wave that begins to move outward
        -as shock wave encounters infalling outer Fe core, high temperatures that result cause further photodisintegration, robbing shock of much energy
        -shock stalls, infalling material accretes, but neutrinosphere develops below shock and since overlying material is so dense that neutrinos cannot easily penetrate, some neutrino energy will be deposited in matter just behind shock
        -additional energy allows shock to resume
        -shock will drive envelope and remainder of nuclear-processed matter in front of it
        -when material becomes optically thin, tremendous optical display results, releasing 10^42J of energy in form of photons
        -if initial mass not too large, remnant will become a NS, but if large produces BH
        
        -believed that H is primordial, having been synthesized immediately following Big Bang
        -much of present-day He was also produced directly from Big Bang, while remainder was generatred from H burning in stellar interiors
        -relative to H and He, Li, Be, and B are very under-abundant because not prominent end products of nuclear reaction changes and can be destroyed by collisions with protons
        -peaks occu for C, N, O, Ne, etc. because they are created as consequence of star's evolutionary trek toward Fe peak and because they are relatively stable alpha-particle-rich nuclei
        -CCSNe are also responsible for generation of significant quantities of O, and Type Ia SNe are responsible for creation of most of the Fe observed
        -s-p[rocess reactions tend to occur in normal phases of stellar evolution, whereas r-processes can occur during a SN when large flux of neutrinos exists
  - - - -WD in binary can also explode as Type Ia SN
        -such SNe lack H lines in spectra; result from explosive burning of C and O
        -if WD takes enough matter from binary companion, can be pushed above Chandrasekhar limit
        -no WD can be heavier than this; if it gains more mass, forced to collapse, like Fe core in most massive stars
        -unlike core, WD still has nuclear fuel: C and O burn to heavier elements, releasing energy which blows it apart
        -no remnant: Fe and other elements scattered to ISM
        -much Fe we now find has been produced in these SNe
- - - - High-mass slope: appears to everywhere obey a PL dN/dm ~ m^-alpha, alpha~2.3=2.4
        -this PL is scale-free and thus a scale-free physical phenomenon is necessary to understand its origin
        -turbulence-regulated models result in estimates for high-mass slope that are generally consistent with observations, as long as parameters carefully chosen (Nam+2021)
        -turbulence produces a PL velocity spectrum and a nearly PL spectrum of logarithmic density, so reasonable to think it could result in M* distributions that are also PLs
        -models that discuss turbulence as origin for high-mass slope make use of theory of turbulent fragmentation, which can be divided into 3 steps (Dib+2010, Krumholz 2014): CMF, bound cores (provides explanation for universality of high-mass slope, but at low-masses, CMF can change potentially providing avenue for variation), and SF from bound cores
        
        Peak: peak is scale-dependent and more challenging to describe
        -this is because physical assumptions of isothermal ideal MHD and gravity that are commonly used to describe high-mass slope are scale-free
        -these theories can't explain origin of peak mass without adding additional pysics
        -for turbulence, models derive characteristic mass scale by setting outer mass scale for MCs and their turbulence by appealing to galaxy-scale physics
        -simplest hypothesis is that origin of peak of IMF is related to mea-density MJ in SF cloud, such that fragment masses trace mean-density MJ
        -this leads to great deal of variation in IMF between 2 different clouds,much more than observed extragalactically and inconsistent with local universal observations
        -to make hypothesis more realistic, can add turbulence
        -if turbulence initially saturated or driven to be saturated by stellar feedback, simulated clouds fragment in a way that is very different than in simple case
        -this affects both peak and PL tail
        -this is method PN model uses to predict peak, where random density is drawn from lognormal PDF and corresponiding M is calculated and compared to mass fragment, similar to criterion for determining whether cores are bound
        -procedure serves to set lognormal-shaped cutoff ofr masses below MJ, evaluated at peak of density PDF (PN, Krumholz)
        -derive peak mass
        -model can explain why IMF looks universal locally, but also variability