Lecture 20: The Nature of Stars

Don't Panic

Sections

1. The Surface or Photosphere Temperature of Stars
2. Spectral Types
3. The Hertzsprung-Russell (HR) Diagram
4. Luminosity Classes
5. The Mass-Luminosity Relation
6. Binaries and Physical Star Groups
7. Population I, II, and III Stars

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The Surface or Photosphere Temperature of Stars

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As was discussed in IAWL Lecture 19: Some Star Basics, luminosity is a basic star parameter needed for understanding and modeling stars.

Another basic parameter is SURFACE TEMPERATURE.

As giant gas balls, stars actually do NOT have a distinct surface.

Their density just decreases going outward and there is a gradual transition to a stellar wind in most cases???.

But there is a layer of the star, where the density falls so low that outward going photons from that layer mostly escape to INFINITY.

This layer is called the photosphere and is loosely what we mean when we say the surface of a star.

[The photosphere is actually somewhat wavelength-dependent. When we say photosphere, we usually mean the visual photosphere: i.e., the photosphere of the visual band 400--700 nm.]

Stellar photospheres approximate BLACKBODY RADIATORS. (See IAWL Lecture 7: Spectra for blackbody spectra.)

They are NOT exactly that since there is a range of temperatures in the photosphere from which photons are emitted. But a fit of blackbody spectru to the spectrum of a photosphere gives a sort of average temperature for the photosphere.

There are also absorbing lines in the dilute stellar atmosphere above the photosphere: these give rise to the absorption line spectra of stars.

Recall blackbody spectra depend only on the absolute temperature of the BLACKBODY RADIATOR---which must be at a single temperature to be that---and on no other properties of the radiatior.

Here are some example blackbody spectra: the 6000 K spectrum approximates the spectrum of the Sun.

Blackbody spectra.

The temperature of a BLACKBODY RADIATOR can be determined by the shape of its blackbody spectrum.

The most direct means is by Wien's law:

Wien's law.

Astronomers in practice would use Wien's law only for a crude estimate of the photosphere temperature. This is because the maximum is often hard to pinpoint exactly in crude data and the temperature from Wien's law is probably not the best estimate of the average temperature of a source which is actually a mixture of blackbody radiatiors.

For a more definitive temperature, astronomers make use of photometry: they compare the ratios of fluxes from different broad-band filters (i.e., filters that accept light over a broad range of wavelengths): from these comparisons, one extracts sufficient information about the shape to determine a good average photosphere temperature.

But broad-band photometry can often be distorted by interstellar dust: this can sometimes be dealt with adequately, but not always.

Because of the problem of interstellar dust and other reasons, it usually better to determine temperature from spectroscopy rather than photometry if one has spectroscopy available.

Question: Spectroscopy is:
1. the study of spectroscopes.
2. the study of flux from broad-band filters.
3. the study of spectra.
4. a set of spectral observations.



Answer 3 and 4 are right, but answer 3 is the primary definition.

Answer 4 is a secondary definition. One often says something like ``the spectroscopy of star X shows ...''

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Spectral Types

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A spectrum is taken by studying fluxes in very narrow wavelength bands.

A spectrum therefore gives information about the spectral lines which are usually (but not always) very narrow in wavelength for stars.

The analysis of a stellar spectrum gives one the photosphere temperature of a star as discussed above in the section The Surface or Photosphere Temperature of Stars.

This analysis in general takes modeling, but stars fall into standard types as determined by spectra--- spectral types and once that type has been modeled accurately, then the modeling has been done once-for-all.

Of course, NO modeling is perfectly accurate and stars are actually all somewhat unique and only approximately conform to ideal spectral types.

But nowadays, star spectrum modeling is very advanced at least for the most common spectral types.

One obtains spectra by dispersing (i.e., spreading out) the light of different wavelengths with a prism or diffraction grating.
[A CD is incidentally a diffraction grating. That is not its function, but disperses light rather prettily.]

In order to understand stellar spectra, recall that the mean cosmic composition by mass (which most stars approximately conform to) is approximately the solar composition:

      H, hydrogen        70.7 +/- 2 %    ,

      He, helium         27.4 +/- 2 %    ,

      metals              1.89 +/- 0.1 %  

        (Cox-28),

        where metals
          in astro-jargon are everything not hydrogen and helium.

      The hydrogen and helium abundances are pretty standard, but
        the metal abundance can vary and, in particular, can much
        lower:  i.e., down to nearly zero????.

Because stars are mostly hydrogen, we expect to see hydrogen lines: in the visual these would be the hydrogen Balmer lines: in usual decreasing order of strength

          Halpha   656.80 nm  (red) ,
          Hbeta    486.131 nm (blue-green) ,
          Hgamma   434.046 nm (blue-violet) ,
          Hdelta   410.173 nm (violet)  , etc. 

The hydrogen Balmer lines can be seen in the color spectrum below.

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Atomic spectra for stellar spectroscopic analysis.

An Angstrom is 10**(-10) meters. It is a non-standard unit, but often used by spectroscopists.

The upper spectra are spectra for actual stars. The lower spectra are just the spectra of particular atoms. One uses the lower spectra to identify the atoms in the upper spectra.

Hydrogen offers the simplest emission line spectrum: a red line, a blue-green line, a blue line, and a violet line.

The hydrogen red line or Halpha at 6563 Angstroms (656 nm) is usually the strongest visible line of hydrogen. Hot interstellar gas often emits this line which gives true color pictures of such gas a reddish or pinkish color.

Halpha and the sodium doublet can both be identified in the star A spectrum.

Credit: NASA: Imagine the Universe.

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We also can expect helium lines, metal lines, and, in cooler stars, molecular lines (lines of bound atoms). Although metals and molecules are only traces in stars, they have many lines, some of them very strong.

[Typically above about 4000 K, molecules cannot exist because the collisions of the particles are sufficiently energetic to prevent formation or to break molecules apart (FK-423).]

We can look at examples of star spectra of various spectral types.

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Star spectra for the visual range 400--700 nm.

The 13 top spectra show regular spectral types from the OBAFGKM SEQUENCE: temperature decreases top to bottom.

The bottom three spectra are special types.

The hydrogen Balmer lines (i.e., the visual band hydrogen lines) are strongest in absorption for A0--A5 stars which have photosphere temperatures of order 9000 K (FK-424).

In the image the blue-green Hbeta (486.132 nm) and violet-blue Hgamma (434.046 nm) absorption lines are quite apparent on and about the A1 spectrum (FK-423).

The red Halpha (656.280 nm) is somewhat darkened out for the A1 region, but can still be made out (FK-423).

The Na I doublet (588.995, 589.592 nm) in the yellow is quite noticeable for the M0 and M5 stars (FK-423).

As we know from Wien's law, the maximum of a blackbody spectrum should shift to the red as temperature decreases.

This is apparently shown in this figure since bright zone forms a broad diagonal from upper right to lower left.

But since colors in images can be easily manipulated, it is hard to know if we are really seeing this effect.

Credit: NOAO/AURA/NSF.

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In the late 19th and early 20th centuries, the spectral types were fixed just empirically (i.e., based on observed characteristics alone) before modern spectral analysis was invented.

Each type was designated by a capital letter.

Originally, the spectral types went AB...P and represented decreasing strength of hydrogen Balmer lines (FK-422; CK-286).

But this ordering turned out NOT to be very physically significant: it was NOT a temperature ordering.

The strength of hydrogen Balmer lines is not monotonic with temperature: they don't just get stronger or weaker as temperature increases.

At low temperatures, the lines are weak because the energy state they depend on is not much EXCITED: i.e., not many atoms are in that state.

At high temperatures, hydrogen tends to be ionized: i.e., it has lost its only electron and become a bare proton. Bare protons are very simple objects and have no atomic transitions and NO lines.

In stellar photospheres the hydrogen Balmer lines tend to be strongest for temperatures of order 9000 K (FK-424).

Rather than rename the spectral types---that would have been the easy way---the types were re-ordered and a few were dropped (FK-422; CK-286).

The standard modern spectral types are OBAFGKM: these are in order of decreasing photosphere temperature.

The OBAFGKM spectral types can be remembered by the mnemonic ``O be a fine girl/guy kiss me.'

[Sometimes the only sensible thing to say.]

There are other less common spectral types.

The spectral types are divided into subtypes: each type has a ten subtypes: 0, 1, 2, 3, 4, 5, 6, 7, 8, 9 in order of decreasing temperature within the type. Thus, for example, one has stars of spectral types:

        O0, O1, ...  ,

        G0, G1, G2, G3, ...  ,  and

        M0, M1, M2, M3, M5, M6, ...

And ``you must remember this,'' the Sun is a G2 star.

Thus, it's a middle of the pack star---but not an ``average star'': spectral types are too diverse for average of all stars to be a useful concept.

Still we use the Sun as convenient standard and use solar mass M_Sun, solar luminosity L_Sun, and solar radius R_Sun as convenient units for star mass, luminosity, and radius.

To paraphrase Protagoras (circa 481--411 BCE): the Sun is the measure of all stars.

We can look at a Table of Spectral Types to see the characteristics of the main spectral types.

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The Hertzsprung-Russell (HR) Diagram

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Now we have considered two basic INTRINSIC stellar parameters: luminosity and photosphere temperature.

In understanding and modeling stars, it would be interesting to know if these parameters were related.

Question: If you plotted two quantities for some objects against each other for a lot of those objects and they were NOT related, you would get a/an:
1. line.
2. inverse-linear curve.
3. inverse-square curve.
4. scatter diagram.



Answer 4 is right.

The data points would be just randomly scattered over the plot.

Some simple function behaviors.

The plot of luminosity versus photosphere temperature (or nearly equivalently spectral type), for stars is called a Hertzsprung-Russell (HR) diagram.

Before we look at a HR diagram, we should point out that temperature increases to the left: just an old convention of astronomy---that no one's had the guts to change.

Cartoon of a Hertzsprung-Russell (HR) diagram with stars (not accurate) (FK-428).

As you can see, a Hertzsprung-Russell (HR) diagram is NOT a scatter diagram.

But it's NOT altogether simple either.

There are various groupings of stars in luminosity-temperature/spectral-type space:

1. Main Sequence: a smooth band of increasing luminosity with temperature.

   In the solar neighborhood and similar stellar neighborhoods---but not all stellar neighborhoods---about 90 % of stars lie on the main sequence (FK-428).

   The Sun is a main-sequence G2 star.

   Not to be coy, main sequence stars are those that burn hydrogen to helium in their cores: here we mean ``burning'' in a nuclear sense, not chemical burning.

   Typically a star spends about 90 % ????? of its nuclear-burning life on the main sequence---which partially explains why main sequence stars are abundant.

2. Giants and Supergiants: These are very luminous post-main-sequence stars.

   They are NOT burning hydrogen in their cores: they are burning it and, perhaps, other elements in concentric shells about a core which is perhaps burning some heavier element than those in the burning shells.

   They can have either hot or cool surfaces.

   The cool ones emit primarily red light, and are called red giants or red supergiants.

   In the solar neighborhood and similar stellar neighborhoods---but not all stellar neighborhoods---about 1 % of stars are giants or supergiants (FK-429).

   But because they are so luminous, they are more conspicuous that their numbers indicate.

3. White Dwarfs (WDs): These are the compact remnants of stars: they are what is left when the all nuclear burning has stopped.

   Though typically their masses are about 0.5 M_Sun, they are of order Earth-size: they are compact and super-dense.

   They shine by residual heat or heat from continuing contraction.

   [Gravitational potential energy is converted to heat energy as stars contract.]

   They just cool off forever.

   In the solar neighborhood and similar stellar neighborhoods---but not all stellar neighborhoods---about 9 % of stars are white dwarfs (FK-429).

   Actually, because they can be very dim, there might be more white dwarfs around than we notice.

We will discuss these kinds of stars further in subsequent lectures:
21. IAWL Lecture 21: Star Formation,
22. IAWL Lecture 22: Main Sequence Star Evolution,
23. IAWL Lecture 23: Late Star Evolution and Star Death, and
24. IAWL Lecture 24: Compact Remnants: White Dwarfs and Neutron Stars.

Given the stellar parameters luminosity and photosphere temperature, one can use Stefan-Boltzmann law of blackbody spectra to determine to good approximation the stellar radius: i.e., the radius of the photosphere treated as a blackbody radiator (FK-429).

We will NOT go into how this is done: it's a tiny bit beyond the scope of this class.

But contours of equal radius can be plotted on an HR diagram.

Cartoon of a Hertzsprung-Russell (HR) diagram with radius contours (not accurate) (FK-428).

The essential behavior is that for fixed radius, the luminosity increases with increasing temperature.

This is because the flux (which is energy per unit time per unit area) of blackbody radiator increases with increasing temperature.

Somewhat more complicated behaviors are:
1. For a fixed temperature, the luminosity. increases with the size of the radiator: i.e., with radius.

2. For a fixed luminosity, temperature decreases with increasing radius.

   As temperature decreases, the flux (which is energy per unit time per unit area) of blackbody radiator decreases.

   Area and therefore radius must rise to compensate if luminosity is held fixed.

We CANNOT usually measure stellar radii directly in most cases because mostly stars CANNOT be resolved.

Relatively close, very-large-radius stars can be resolved with special techniques or instruments: one still CANNOT see a lot of detail, but at least the finite size of the star in the image is NOT false.

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Betelgeuse imagined by the HST.

Betelgeuse is an M1 Iab red supergiant star. It is 131 pc from Earth.

It is the eastern shoulder of Orion: i.e., the left shoulder on the image.

Orion is, of course, a giant hunter of Greek mythology: he pursued the Pleiades and was slain by Artemis (Ba-855).

The lines joining the stars are NOT present on the sky, of course.

Orion is one of the three constellations anyone can recognize: the other two are the Big Dipper (officially an asterism in Ursa Major) and Cassiopeia (the big W): bother are in the northern sky and are all-year constellations.

Credit: NASA/HST

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Luminosity Classes

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Spectral types essentially classify stars by photosphere temperature.

There is also a luminosity classification.

Giving spectral type and luminosity class pretty uniquely specifies most kinds of star.

NOT completely uniquely since in both classification schemes the categories are only fairly narrow BINS, and not very fine BINS. Also star of the same spectral type and luminosity class can have vary from each other in other ways: these variations are are sometimes important???.

The luminosity classes are best illustrated by a HR diagram.

Cartoon of a Hertzsprung-Russell (HR) diagram with luminosity classes (not accurate) (FK-430).

White dwarfs do NOT have a luminosity class---but I think they ought to be class VI---just my wild and crazy idea.

Often one gives both classifications together and calls that full classification just spectral type for simplicity (FK-A-7): e.g., the Sun's spectral type is G2 V.

We can now understand, for example, the spectral types of the stars in the Table of Brightest Stars.

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Mass-Luminosity Relation

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Luminosity, photosphere temperature, and photosphere radius are basic stellar parameters needed in understanding and modeling stars.

Another one is stellar MASS.

In fact, for an isolated star there is no observational way of determining MASS.

But orbital parameters of gravitationally bound pairs of stars (binaries) allow MASSmass to determined using Newtonian physics.

The masses of all the spectral types can be determined by examples in binaries: we expect the mass of a spectral type (including luminosity class) to be an approximately fixed value usually.

Question: For main sequence stars, there is a temperature-luminosity relation.

This relation on the HR diagram with temperature is:

1. linear.
2. quadratic.
3. inverse-square.
4. main sequence itself.



Answer 4 is right.

There is also a mass-luminosity relation for the main sequence: the masses being determined by examples of each kind of main-sequence star found in a well observed binary.

Cartoon of the mass-luminosity relation for main sequence stars (not accurate).

For the math cognoscenti, luminosity increases approximately as the a power of 3.5 to 4 of the mass (Cl-40):

       L proportional to M**y , where y is in the range 3.5 to 4.

       This means if M doubles, L increases by a factor 11 to 16.

Important results to remember forever in this course:
As star mass increases:
1. Luminosity increases for each life phase???.

2. Total lifetime decreases and lifetime of each phase decreases.

   The more massive a star is the more nuclear fuel it has, but it burns the fuel more rapidly and that is the overriding consideration in the nuclear fuel burning phases.

3. The frequency of stars decreases at least for stars more massive than about 0.25 M_Sun (HI-331).

   Cartoon of star frequency in the solar neighborhood (not accurate) (HI-330--331,339).

   The upper limit to star mass is about 100 M_Sun and these stars are very rare. Protostars more massive than that expel mass through strong solar winds, and so never become more massive than about 100 M_Sun on the main sequence (CK-305--306).

   Below about 0.5 M_Sun, it is unclear if the frequency of stars of a given mass starts declining with smaller mass or keeps rising (HI-331).

   Below about 0.08 M_Sun, one has brown dwarfs rather than stars: see the discussion in IAWL Lecture 22: The Main Sequence Life of Stars.

The mass-luminosity relation for the main sequence can also be displayed on an HR diagram.

Cartoon of a Hertzsprung-Russell (HR) diagram with main sequence star masses (not accurate) (FK-433).

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Binaries and Physical Star Groups

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Binaries (i.e., gravitationally bound pairs of stars) are a whole massive subject in themselves---which we will mostly skip in this course.

The stars in a binary orbit their mutual center of mass in ellipses.

Binaries and binary orbits (FK-435).

We will just mention why binaries are important:

1. They are numerous.

   In the solar neighborhood---which may be representative of the observable universe as a whole---about 2/3 of all stars are in binaries.

   Thus, binaries are as common a singles.

2. Star masses can be obtained by analyzing the observed parameters of binaries using Newtonian physics.

   This allows the masses of all spectral types to fixed.

   Mass is a basic parameter in understanding and modeling stars.

3. Close binaries (i.e., binaries where the stars orbit relatively closely to each other) often interact: e.g., they can heat each other or exchange mass during certain phases of stellar evolution.

   Thus, close binaries can show behaviors singles never can.

   These behaviors are sometimes cosmically important: e.g., some kinds of supernovae happen only or mostly in close binary systems.

There are larger physical groups of stars: e.g., gravitationally bound or gravitationally interacting groups.

We will just briefly mention them.

1. Triples, etc.: These are systems of 3 or more gravitationally bound stars with complex orbits.

2. Open Clusters: These are loosely-bound, irregularly-shaped groups of stars consisting of order 100 to 1000 stars with a size scales of order 4 to 20 pc (HI-392--393).

   The stars formed at about the same time (i.e., same time to within a few million or tens of millions of years: HI-338) in a star formation region.

   In the Milk Way they are only in the Galactic disk.

   Both internal and external gravitational perturbations tend to break up open clusters and they probably only survive for a few hundred million years.

   Probably the best known example of an open cluster is the Pleiades in Taurus.

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   The Pleiades, an open star cluster, in Taurus centered about 127 pc away and spanning about 13 pc.

   The stars in this image are really unresolved point sources.

   The round and pointed stars are finite in size because of finite resolution of the imaging system.

   The Pleiades stars themselves have a glow around them because their light is being reflected from interstellar dust.

   Credit: NOAO/AURA/NSF.

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   The Pleiades are well known in many cultures because they form a distinct group on the sky.

   In Europe they have also been called the Seven Sisters and their Japanese name is Subaru like the car and the telescope. Chaucer alludes to them in his Chauntecleer tale.

   Usually 6 Pleiades at least can be seen with the unaided eye; 9 can be seen under good conditions; 14 were claimed visible by Kepler. Telescopically, the cluster has over 500 stars.

   To find the Pleiades, one can use the constellations as landmarks---well SKYMARKS

   First, locate Orion and Sirius (the brightest star in the sky) off to the lower left of Orion (south-east on the sky). A line from Sirius though the belt of Orion and then through the bright orangy Aldebaran (the eye of Taurus leads pretty much to the Pleiades---a distinct close little group of six or more unaided-eye stars---there are at least 500 stars altogether in the cluster recall.

   The northern constellations map below illustrates the method.

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   The northern constellations: a mid-winter night-time view judging from the position of old man Orion.

   Credit: Mount Wilson Observatory StarMap program by Bob Donahue. StarMap is fortran program, but it's been broke since 2000jan03. Download site: Univ. of Tennessee, Knoxville Astro course; more precisely here.

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3. Associations: These are structures of a few to a few hundred stars and span of order 10 to 100 pc (HI-393,395).

   They are generally gravitationally UNBOUND though gravitationally interacting (FK-456).

   [The division between open cluster and associations is probably not sharp. Open clusters are a bit more compact with of order 100 to 1000 stars with a size scales of order 4 to 20 pc (HI-392--393), and so are more definitely gravitationally bound.]

   The stars in an association formed at about the same time (i.e., same time to within a few million or tens of millions of years: HI-338) in a star formation region.

   In the Milk Way they are only in the Galactic disk.

   The most discussed kind are the OB associations: these are groups of hot (and therefore bluish), young OB stars.

   The stars' own kinetic energies and gravitational perturbations will break them up within a few tens to hundreds of millions of years.

4. Globular Clusters: These are compact, dense, spherical, gravitationally-bound systems of stars (HI-395).

   They can have from of order 20,000 to several million stars and their central concentrations have diameters of order to 5 to 25 pc.

   Near the center of a globular cluster, there could be 10,000 stars per cubic parsec: the night sky would be 10 times brighter than the full Moon.

   In the Milk Way, there are about 150 globular clusters and they are spread around in the Galactic halo.

   [As of 2000, there were 146 known globular clusters in the Galactic halo (Cox-554). There might be 15 to 30 more which are hidden by interstellar dust.]

   Globular clusters are very old at least in the Milky Way: i.e., their stars, which formed at about the same time, are very old.

   Globular cluster ages have been calculated to about 12.5 Gyr (FK-638).

   This age is an important lower limit on the age of the observable universe.

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   Globular cluster M15 in Pegasus at about 13 kpc.

   The image approximates true color.

   The small pink nebula in the upper left is a planetary nebula discovered in a globular cluster (Pease 1928): there are still only a few other globular clusters.

   Credit: NASA and The Hubble Heritage Team (STScI/AURA) Acknowledgment: H. Bond (STScI).

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5. Galaxies, clusters of galaxies, superclusters of galaxies, and large scale structure: These are all large scale groupings of stars that we will take up in IAWL Lecture 27: The Milky Way and IAWL Lecture 28: Galaxies.

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Population I, II, and III Stars

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There is another classification scheme for stars that divides them into Population I, II, and III stars based on their age or, nearly equivalently, their metallicity (i.e., abundance of metals [i.e., all elements heavier than H and He]) (HI-413--414):

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1. Population I stars: These are relatively young and metal-rich stars:
   ages typically from a few hundred million to a few billion years;

   metallicity typically 2 to 4 % by mass.

   The Sun is a Population I star.

   Population I stars are found mainly in the galactic disks and galactic bulges, but much not in galactic halos of spiral galaxies.

   They are not much found and in elliptical galaxies.

   We will discuss disks, bulges, and halos of galaxies later in IAWL Lecture 27: The Milky Way and IAWL Lecture 28: Galaxies, but here we just give an image as a preview.

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   M104, NGC 4594 (Sa): The Sombrero Galaxy in Virgo.

   The Sombrero Galaxy is Sa spiral galaxy seen nearly edge-on: it is tilted 6 degrees from its equatorial plane.

   It is in the Virgo cluster at the southern edge.

   It is about 9 Mpc from Earth and the disk is about 17 kpc in diameter.

   The disk's angular diameter is about 6 arcminutes which is about 1/5 of the Moon's angular diameter

   The large bulge is the giveaway for being an Sa galaxy even thought the arms are not easily discerned.

   Note the strong dust lane in the disk.

   There is a swarm of about 2000 globular clusters in the halo. This about 13 times more than has the Milky Way. These globular clusters, like the Milky Way's, are calculated to be about 10--13 Gyr old.

   It is hard to tell which of the star-like objects in the halo are foreground stars in the Milky Way and which are globular clusters: but we know some of them are globular clusters.

   Credit: NASA and The Hubble Heritage Team (STScI/AURA). For more information see the SEDS M104 page.

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2. Population II stars: These are relatively old and metal-poor stars:
   ages typically from a 10 to 12 or so Gyr (HI-415;

   metallicity typically of order 0.1 % by mass, but it can be a lot lower.

   The lowest metallicity stars on record have have iron at of order 10**-5 that of the Sun and carbon at of order 1/25 of the Sun (Cayrel, R. 2005, Nature, 434, 838).

   Population II stars are found mainly in the galactic halos and galactic bulges of spiral galaxies, and in small fraction in galactic disks of spiral galaxies and in elliptical galaxies.

3. Population III stars: These are hypothetical nearly zero-metallicity stars formed in the first age of star formation after the big bang.

   Such stars should be about 13 Gyr old.

   They have never been observed.

   In the modern universe they must either be very rare or non-existent.

   It may be that the first generation of stars were all massive stars and were not long-lived.

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Why is there varying metallicity with star age?

As we will discuss in subsequent lectures, stars synthesize metals in their nuclear-fuel burning lifetimes and eject those into interstellar medium (ISM) either through stellar winds or as supernovae.

Out of the ISM new generations of stars form as is discussed in IAWL Lecture 21: Star Formation.

Thus, every new generation of stars has higher metallicity than the generation before.

Because stars have vastly varying lifetimes, some old, metal-poor stars persist on and overlap with younger, metal-richer stars.

Thus, there are very old, very metal-poor stars today: e.g., stars in globular cluster have calculated ages of 12.5 Gyr (FK-638) as we mentioned above.

Question: Can there really be a sharp dichotomy between Population I stars and Population II stars?
1. Yes, of course.

2. Yes, but not of course: i.e., it takes an explanation.

3. No. There must be a continuum of metallicities since star formation has been more or less continuous since the first generation of stars.

   The terms Population I stars and Population II stars are just convenient labels for thinking and rough classification purposes.



Answer 3 is right (CM-373) although the above discussion doesn't make that obvious.