CNO cycle

The CNO cycle (for carbonnitrogenoxygen) is one of the two known sets of fusion reactions by which stars convert hydrogen to helium, the other being the proton–proton chain reaction (pp-chain reaction). Unlike the latter, the CNO cycle is a catalytic cycle. It is dominant in stars that are more than 1.3 times as massive as the Sun.[1]

In the CNO cycle, four protons fuse, using carbon, nitrogen, and oxygen isotopes as catalysts, to produce one alpha particle, two positrons and two electron neutrinos. Although there are various paths and catalysts involved in the CNO cycles, all these cycles have the same net result:

4 1
1
H
 +  2
e
 →  4
2
He
 +  2
e+
 +  2
e
 +  2
ν
e
 +  3
γ
 +  24.7 MeV  →  4
2
He
 +  2
ν
e
 +  3
γ
 +  26.7 MeV

The positrons will almost instantly annihilate with electrons, releasing energy in the form of gamma rays. The neutrinos escape from the star carrying away some energy. One nucleus goes on to become carbon, nitrogen, and oxygen isotopes through a number of transformations in an endless loop.

The proton–proton chain is more prominent in stars the mass of the Sun or less. This difference stems from temperature dependency differences between the two reactions; pp-chain reaction starts at temperatures around 4×106 K[2] (4 megakelvin), making it the dominant energy source in smaller stars. A self-maintaining CNO chain starts at approximately 15×106 K, but its energy output rises much more rapidly with increasing temperatures[1] so that it becomes the dominant source of energy at approximately 17×106 K.[3] The Sun has a core temperature of around 15.7×106 K, and only 1.7% of 4
He
nuclei produced in the Sun are born in the CNO cycle. The CNO-I process was independently proposed by Carl von Weizsäcker[4][5] and Hans Bethe[6][7] in the late 1930s.

CNO Cycle
Overview of the CNO-I Cycle
CNO cycle
Carbon-Nitrogen-Oxygen Cycle-1

Cold CNO cycles

Under typical conditions found in stars, catalytic hydrogen burning by the CNO cycles is limited by proton captures. Specifically, the timescale for beta decay of the radioactive nuclei produced is faster than the timescale for fusion. Because of the long timescales involved, the cold CNO cycles convert hydrogen to helium slowly, allowing them to power stars in quiescent equilibrium for many years.

CNO-I

The first proposed catalytic cycle for the conversion of hydrogen into helium was initially called the carbon–nitrogen cycle (CN-cycle), also referred to as the Bethe–Weizsäcker cycle in honor of the independent work of Carl von Weizsäcker in 1937-38[4][5] and Hans Bethe. Bethe's 1939 papers on the CN-cycle[6][7] drew on three earlier papers written in collaboration with Robert Bacher and Milton Stanley Livingston[8][9][10] and which came to be known informally as "Bethe's Bible." It was considered the standard work on nuclear physics for many years and was a significant factor in his being awarded the 1967 Nobel Prize in Physics.[11] Bethe's original calculations suggested the CN-cycle was the Sun's primary source of energy.[6][7] This conclusion arose from what is now-known as a mistaken belief: that the abundance of nitrogen in the sun is approximately 10%, when it is actually less than half a percent.[12] The CN-cycle, named as it contains no stable isotope of oxygen involves the following cycle of transformations: 12
6
C
 → 13
7
N
 → 13
6
C
 → 14
7
N
 → 15
8
O
 → 15
7
N
 → 12
6
C
.[12] This cycle is now understood as being the first part of a larger process, the CNO-cycle, and the main reactions in this part of the cycle (CNO-I) are:[12]

12
6
C
 
1
1
H
 
→  13
7
N
 

γ
 
    1.95 MeV
13
7
N
 
    →  13
6
C
 

e+
 

ν
e
 
1.20 MeV (half-life of 9.965 minutes[13])
13
6
C
 
1
1
H
 
→  14
7
N
 

γ
 
    7.54 MeV
14
7
N
 
1
1
H
 
→  15
8
O
 

γ
 
    7.35 MeV
15
8
O
 
    →  15
7
N
 

e+
 

ν
e
 
1.73 MeV (half-life of 122.24 seconds[13])
15
7
N
 
1
1
H
 
→  12
6
C
 
4
2
He
 
    4.96 MeV

where the carbon-12 nucleus used in the first reaction is regenerated in the last reaction. After the two positrons emitted annihilate with two ambient electrons producing an additional 2.04 MeV, the total energy released in one cycle is 26.73 MeV; in some texts, authors are erroneously including the positron annihilation energy in with the beta-decay Q-value and then neglecting the equal amount of energy released by annihilation, leading to possible confusion. All values are calculated with reference to the Atomic Mass Evaluation 2003.[14]

The limiting (slowest) reaction in the CNO-I cycle is the proton capture on 14
7
N
. In 2006 it was experimentally measured down to stellar energies, revising the calculated age of globular clusters by around 1 billion years.[15]

The neutrinos emitted in beta decay will have a spectrum of energy ranges, because although momentum is conserved, the momentum can be shared in any way between the positron and neutrino, with either emitted at rest and the other taking away the full energy, or anything in between, so long as all the energy from the Q-value is used. The total momentum received by the electron and the neutrino is not great enough to cause a significant recoil of the much heavier daughter nucleus and hence, its contribution to kinetic energy of the products, for the precision of values given here, can be neglected. Thus the neutrino emitted during the decay of nitrogen-13 can have an energy from zero up to 1.20 MeV, and the neutrino emitted during the decay of oxygen-15 can have an energy from zero up to 1.73 MeV. On average, about 1.7 MeV of the total energy output is taken away by neutrinos for each loop of the cycle, leaving about 25 MeV available for producing luminosity.[16]

CNO-II

In a minor branch of the above reaction, occurring in the Sun's core 0.04% of the time, the final reaction involving 15
7
N
shown above does not produce carbon-12 and an alpha particle, but instead produces oxygen-16 and a photon and continues 15
7
N
16
8
O
17
9
F
17
8
O
14
7
N
15
8
O
15
7
N
:

15
7
N
 
1
1
H
 
→  16
8
O
 

γ
 
    12.13 MeV
16
8
O
 
1
1
H
 
→  17
9
F
 

γ
 
    0.60 MeV
17
9
F
 
    →  17
8
O
 

e+
 

ν
e
 
2.76 MeV (half-life of 64.49 seconds)
17
8
O
 
1
1
H
 
→  14
7
N
 
4
2
He
 
    1.19 MeV
14
7
N
 
1
1
H
 
→  15
8
O
 

γ
 
    7.35 MeV
15
8
O
 
    →  15
7
N
 

e+
 

ν
e
 
2.75 MeV (half-life of 122.24 seconds)

Like the carbon, nitrogen, and oxygen involved in the main branch, the fluorine produced in the minor branch is merely an intermediate product and at steady state, does not accumulate in the star.

CNO-III

This subdominant branch is significant only for massive stars. The reactions are started when one of the reactions in CNO-II results in fluorine-18 and gamma instead of nitrogen-14 and alpha, and continues 17
8
O
18
9
F
18
8
O
15
7
N
16
8
O
17
9
F
17
8
O
:

17
8
O
 
1
1
H
 
→  18
9
F
 

γ
 
    5.61 MeV
18
9
F
 
    →  18
8
O
 

e+
 

ν
e
 
1.656 MeV (half-life of 109.771 minutes)
18
8
O
 
1
1
H
 
→  15
7
N
 
4
2
He
 
    3.98 MeV
15
7
N
 
1
1
H
 
→  16
8
O
 

γ
 
    12.13 MeV
16
8
O
 
1
1
H
 
→  17
9
F
 

γ
 
    0.60 MeV
17
9
F
 
    →  17
8
O
 

e+
 

ν
e
 
2.76 MeV (half-life of 64.49 seconds)

CNO-IV

NuclearReaction
A proton reacts with a nucleus causing release of an alpha particle.

Like the CNO-III, this branch is also only significant in massive stars. The reactions are started when one of the reactions in CNO-III results in fluorine-19 and gamma instead of nitrogen-15 and alpha, and continues 18
8
O
19
9
F
16
8
O
17
9
F
17
8
O
18
9
F
18
8
O
:

18
8
O
 
1
1
H
 
→  19
9
F
 

γ
 
    7.994 MeV
19
9
F
 
1
1
H
 
→  16
8
O
 
4
2
He
 
    8.114 MeV
16
8
O
 
1
1
H
 
→  17
9
F
 

γ
 
    0.60 MeV
17
9
F
 
    →  17
8
O
 

e+
 

ν
e
 
2.76 MeV (half-life of 64.49 seconds)
17
8
O
 
1
1
H
 
→  18
9
F
 

γ
 
    5.61 MeV
18
9
F
 
    →  18
8
O
 

e+
 

ν
e
 
1.656 MeV (half-life of 109.771 minutes)

Hot CNO cycles

Under conditions of higher temperature and pressure, such as those found in novae and x-ray bursts, the rate of proton captures exceeds the rate of beta-decay, pushing the burning to the proton drip line. The essential idea is that a radioactive species will capture a proton before it can beta decay, opening new nuclear burning pathways that are otherwise inaccessible. Because of the higher temperatures involved, these catalytic cycles are typically referred to as the hot CNO cycles; because the timescales are limited by beta decays instead of proton captures, they are also called the beta-limited CNO cycles.

HCNO-I

The difference between the CNO-I cycle and the HCNO-I cycle is that 13
7
N
captures a proton instead of decaying, leading to the total sequence 12
6
C
13
7
N
14
8
O
14
7
N
15
8
O
15
7
N
12
6
C
:

12
6
C
 
1
1
H
 
→  13
7
N
 

γ
 
    1.95 MeV
13
7
N
 
1
1
H
 
→  14
8
O
 

γ
 
    4.63 MeV
14
8
O
 
    →  14
7
N
 

e+
 

ν
e
 
5.14 MeV (half-life of 70.641 seconds)
14
7
N
 
1
1
H
 
→  15
8
O
 

γ
 
    7.35 MeV
15
8
O
 
    →  15
7
N
 

e+
 

ν
e
 
2.75 MeV (half-life of 122.24 seconds)
15
7
N
 
1
1
H
 
→  12
6
C
 
4
2
He
 
    4.96 MeV

HCNO-II

The notable difference between the CNO-II cycle and the HCNO-II cycle is that 17
9
F
captures a proton instead of decaying, and neon is produced in a subsequent reaction on 18
9
F
, leading to the total sequence 15
7
N
16
8
O
17
9
F
18
10
Ne
18
9
F
15
8
O
15
7
N
:

15
7
N
 
1
1
H
 
→  16
8
O
 

γ
 
    12.13 MeV
16
8
O
 
1
1
H
 
→  17
9
F
 

γ
 
    0.60 MeV
17
9
F
 
1
1
H
 
→  18
10
Ne
 

γ
 
    3.92 MeV
18
10
Ne
 
    →  18
9
F
 

e+
 

ν
e
 
4.44 MeV (half-life of 1.672 seconds)
18
9
F
 
1
1
H
 
→  15
8
O
 
4
2
He
 
    2.88 MeV
15
8
O
 
    →  15
7
N
 

e+
 

ν
e
 
2.75 MeV (half-life of 122.24 seconds)

HCNO-III

An alternative to the HCNO-II cycle is that 18
9
F
captures a proton moving towards higher mass and using the same helium production mechanism as the CNO-IV cycle as 18
9
F
19
10
Ne
19
9
F
16
8
O
17
9
F
18
10
Ne
18
9
F
:

18
9
F
 
1
1
H
 
→  19
10
Ne
 

γ
 
    6.41 MeV
19
10
Ne
 
    →  19
9
F
 

e+
 

ν
e
 
3.32 MeV (half-life of 17.22 seconds)
19
9
F
 
1
1
H
 
→  16
8
O
 
4
2
He
 
    8.11 MeV
16
8
O
 
1
1
H
 
→  17
9
F
 

γ
 
    0.60 MeV
17
9
F
 
1
1
H
 
→  18
10
Ne
 

γ
 
    3.92 MeV
18
10
Ne
 
    →  18
9
F
 

e+
 

ν
e
 
4.44 MeV (half-life of 1.672 seconds)

Use in astronomy

While the total number of "catalytic" nuclei are conserved in the cycle, in stellar evolution the relative proportions of the nuclei are altered. When the cycle is run to equilibrium, the ratio of the carbon-12/carbon-13 nuclei is driven to 3.5, and nitrogen-14 becomes the most numerous nucleus, regardless of initial composition. During a star's evolution, convective mixing episodes moves material, within which the CNO cycle has operated, from the star's interior to the surface, altering the observed composition of the star. Red giant stars are observed to have lower carbon-12/carbon-13 and carbon-12/nitrogen-14 ratios than do main sequence stars, which is considered to be convincing evidence for the operation of the CNO cycle.

See also

References

  1. ^ a b Salaris, Maurizio; Cassisi, Santi (2005). Evolution of Stars and Stellar Populations. John Wiley and Sons. pp. 119–121. ISBN 0-470-09220-3.
  2. ^ Reid, I. Neill; Hawley, Suzanne L. (2005). "The Structure, Formation and Evolution of Low-Mass Stars and Brown Dwarfs – Energy Generation". New Light on Dark Stars: Red Dwarfs, Low-Mass Stars, Brown Dwarfs. Springer-Praxis Books in Astrophysics and Astronomy (2nd ed.). Springer Science & Business Media. p. 108-111. ISBN 3-540-25124-3.
  3. ^ Schuler, S. C.; King, J. R.; The, L.-S. (2009). "Stellar Nucleosynthesis in the Hyades Open Cluster". The Astrophysical Journal. 701 (1): 837–849. arXiv:0906.4812. Bibcode:2009ApJ...701..837S. doi:10.1088/0004-637X/701/1/837.
  4. ^ a b von Weizsäcker, Carl F. (1937). "Über Elementumwandlungen in Innern der Sterne I" [On Transformations of Elements in the Interiors of Stars I]. Physikalische Zeitschrift. 38: 176–191.
  5. ^ a b von Weizsäcker, Carl F. (1938). "Über Elementumwandlungen in Innern der Sterne II" [On Transformations of Elements in the Interiors of Stars II]. Physikalische Zeitschrift. 39: 633–646.
  6. ^ a b c Bethe, Hans A. (1939). "Energy Production in Stars". Physical Review. 55 (1): 103. doi:10.1103/PhysRev.55.103.
  7. ^ a b c Bethe, Hans A. (1939). "Energy Production in Stars". Physical Review. 55 (5): 434-456. doi:10.1103/PhysRev.55.434.
  8. ^ Bethe, Hans A.; Bacher, Robert (1936). "Nuclear Physics, A: Stationary States of Nuclei". Reviews of Modern Physics. 8 (2): 82–229. doi:10.1103/RevModPhys.8.82.
  9. ^ Bethe, Hans A. (1937). "Nuclear Physics, B: Nuclear Dynamics, Theoretical". Reviews of Modern Physics. 9 (2): 69–244. doi:10.1103/RevModPhys.9.69.
  10. ^ Bethe, Hans A.; Livingston, Milton S. (1937). "Nuclear Physics, C: Nuclear Dynamics, Experimental". Reviews of Modern Physics. 9 (2): 245–390. doi:10.1103/RevModPhys.9.245.
  11. ^ Bardi, Jason Socrates (23 January 2008). "Landmarks: What Makes the Stars Shine?". Physical Review Focus. 21 (3). Retrieved 26 November 2018.
  12. ^ a b c Krane, Kenneth S. (1988). Introductory Nuclear Physics. John Wiley & Sons. p. 537. ISBN 0-471-80553-X.
  13. ^ a b Ray, Alak (2010). "Massive Stars as Thermonuclear Reactors and their Explosions Following Core Collapse". In Goswami, Aruna; Reddy, B. Eswar (eds.). Principles and Perspectives in Cosmochemistry. Springer Science & Business Media. p. 233. ISBN 9783642103681.
  14. ^ Wapstra, Aaldert; Audi, Georges (18 November 2003). "The 2003 Atomic Mass Evaluation". Atomic Mass Data Center. Retrieved 25 October 2011.
  15. ^ LUNA Collaboration; Lemut, A.; Bemmerer, D.; Confortola, F.; Bonetti, R.; Broggini, C.; Corvisiero, P.; Costantini, H.; Cruz, J.; Formicola, A.; Fülöp, Zs.; Gervino, G.; et al. (2006). "First measurement of the 14N(p,γ)15O cross section down to 70 keV". Physics Letters B. 634: 483–487. arXiv:nucl-ex/0602012. Bibcode:2006PhLB..634..483L. doi:10.1016/j.physletb.2006.02.021.
  16. ^ Scheffler, Helmut; Elsässer, Hans (1990). Die Physik der Sterne und der Sonne [The Physics of the Stars and the Sun]. Bibliographisches Institut (Mannheim, Wien, Zürich). ISBN 3-411-14172-7.

Further reading

4 Lacertae

4 Lacertae is a single star in the northern constellation Lacerta, located about 1,900 light years away. This object visible to the naked eye as a white-hued star with an apparent visual magnitude of 4.55. It is moving closer to the Earth with a heliocentric radial velocity of −26 km/s. This star is a suspected member of the Lac OB1 association.This is a supergiant star with a stellar classification of A0 Ib. The surface abundances show evidence of material that has been processed via the CNO cycle at the core. It has ten times the mass of the Sun and has expanded to about 59 times the Sun's radius. The star is around 25 million years old and is spinning with a projected rotational velocity of 28 km/s.

B-type main-sequence star

A B-type main-sequence star (B V) is a main-sequence (hydrogen-burning) star of spectral type B and luminosity class V. These stars have from 2 to 16 times the mass of the Sun and surface temperatures between 10,000 and 30,000 K. B-type stars are extremely luminous and blue. Their spectra have neutral helium, which are most prominent at the B2 subclass, and moderate hydrogen lines. Examples include Regulus and Algol A.This class of stars was introduced with the Harvard sequence of stellar spectra and published in the Revised Harvard photometry catalogue. The definition of type B-type stars was the presence of non-ionized helium lines with the absence of singly ionized helium in the blue-violet portion of the spectrum. All of the spectral classes, including the B type, were subdivided with a numerical suffix that indicated the degree to which they approached the next classification. Thus B2 is 1/5 of the way from type B (or B0) to type A.Later, however, more refined spectra showed lines of ionized helium for stars of type B0. Likewise, A0 stars also show weak lines of non-ionized helium. Subsequent catalogues of stellar spectra classified the stars based on the strengths of absorption lines at specific frequencies, or by comparing the strengths of different lines. Thus, in the MK Classification system, the spectral class B0 has the line at wavelength 439 nm being stronger than the line at 420 nm. The Balmer series of hydrogen lines grows stronger through the B class, then peak at type A2. The lines of ionized silicon are used to determine the sub-class of the B-type stars, while magnesium lines are used to distinguish between the temperature classes.Type-B stars don't have a corona and lack a convection zone in their outer atmosphere. They have a higher mass loss rate than smaller stars such as the Sun, and their stellar wind has velocities of about 3,000 km/s. The energy generation in main-sequence B-type stars comes from the CNO cycle of thermonuclear fusion. Because the CNO cycle is very temperature sensitive, the energy generation is heavily concentrated at the center of the star, which results in a convection zone about the core. This results in a steady mixing of the hydrogen fuel with the helium byproduct of the nuclear fusion. Many B-type stars have a rapid rate of rotation, with an equatorial rotation velocity of about 200 km/s.

Dredge-up

A dredge-up is a period in the evolution of a star where a surface convection zone extends down to the layers where material has undergone nuclear fusion. As a result, the fusion products are mixed into the outer layers of the stellar atmosphere where they can appear in the spectrum of the star.

The first dredge-up occurs when a main-sequence star enters the red-giant branch. As a result of the convective mixing, the outer atmosphere will display the spectral signature of hydrogen fusion: the 12C/13C and C/N ratios are lowered, and the surface abundances of lithium and beryllium may be reduced.

The second dredge-up occurs in stars with 4–8 solar masses. When helium fusion comes to an end at the core, convection mixes the products of the CNO cycle. This second dredge-up results in an increase in the surface abundance of 4He and 14N, whereas the amount of 12C and 16O decreases.The third dredge-up occurs after a star enters the asymptotic giant branch and a flash occurs along a helium-burning shell. This dredge-up causes helium, carbon and the s-process products to be brought to the surface. The result is an increase in the abundance of carbon relative to oxygen, which can create a carbon star.The names of the dredge-ups are set by the evolutionary and structural state of the star in which each occurs, not by the sequence experienced by the star. As a result, lower-mass stars experience the first and third dredge-ups in their evolution but not the second.

Georgeanne R. Caughlan

Georgeanne (Jan) Caughlan (née Robertson; 25 October 1916 – 3 January 1994) was an American astrophysicist known for her work on stellar energy generation. Her compilation of experimental data of the rates of nuclear reactions was instrumental in establishing the theory of nucleosynthesis that led to a Nobel Prize for William A. Fowler.

Horizontal branch

The horizontal branch (HB) is a stage of stellar evolution that immediately follows the red giant branch in stars whose masses are similar to the Sun's. Horizontal-branch stars are powered by helium fusion in the core (via the triple-alpha process) and by hydrogen fusion (via the CNO cycle) in a shell surrounding the core. The onset of core helium fusion at the tip of the red giant branch causes substantial changes in stellar structure, resulting in an overall reduction in luminosity, some contraction of the stellar envelope, and the surface reaching higher temperatures.

Main sequence

In astronomy, the main sequence is a continuous and distinctive band of stars that appears on plots of stellar color versus brightness. These color-magnitude plots are known as Hertzsprung–Russell diagrams after their co-developers, Ejnar Hertzsprung and Henry Norris Russell. Stars on this band are known as main-sequence stars or dwarf stars. These are the most numerous true stars in the universe, and include the Earth's Sun.

After condensation and ignition of a star, it generates thermal energy in its dense core region through nuclear fusion of hydrogen into helium. During this stage of the star's lifetime, it is located on the main sequence at a position determined primarily by its mass, but also based upon its chemical composition and age. The cores of main-sequence stars are in hydrostatic equilibrium, where outward thermal pressure from the hot core is balanced by the inward pressure of gravitational collapse from the overlying layers. The strong dependence of the rate of energy generation on temperature and pressure helps to sustain this balance. Energy generated at the core makes its way to the surface and is radiated away at the photosphere. The energy is carried by either radiation or convection, with the latter occurring in regions with steeper temperature gradients, higher opacity or both.

The main sequence is sometimes divided into upper and lower parts, based on the dominant process that a star uses to generate energy. Stars below about 1.5 times the mass of the Sun (1.5 M☉) primarily fuse hydrogen atoms together in a series of stages to form helium, a sequence called the proton–proton chain. Above this mass, in the upper main sequence, the nuclear fusion process mainly uses atoms of carbon, nitrogen and oxygen as intermediaries in the CNO cycle that produces helium from hydrogen atoms. Main-sequence stars with more than two solar masses undergo convection in their core regions, which acts to stir up the newly created helium and maintain the proportion of fuel needed for fusion to occur. Below this mass, stars have cores that are entirely radiative with convective zones near the surface. With decreasing stellar mass, the proportion of the star forming a convective envelope steadily increases. Main-sequence stars below 0.4 M☉ undergo convection throughout their mass. When core convection does not occur, a helium-rich core develops surrounded by an outer layer of hydrogen.

In general, the more massive a star is, the shorter its lifespan on the main sequence. After the hydrogen fuel at the core has been consumed, the star evolves away from the main sequence on the HR diagram, into a supergiant, red giant, or directly to a white dwarf.

NGC 3603-A1

NGC 3603-A1 (HD 97950A1) is a double-eclipsing binary star system located at the centre of the HD 97950 cluster in the NGC 3603 star-forming region, about 20,000 light years from Earth. Both stars are of spectral type WN6h and among the most luminous and most massive known.

HD 97950 was catalogued as a star, but was known to be a dense cluster or close multiple star. In 1926, the six brightest members were given letters from A to F, although several of them have since been resolved into more than one star. Star A was first resolved into three components using speckle interferometry, although they can now be directly imaged using space-based or adaptive optics. Component A1 was finally determined to be a spectroscopic binary.The two component stars of NGC 3603-A1 circle each other every 3.77 days, and show brightness variations of about 0.3 magnitudes due to eclipses. The stars orbit very close to each other, separated by barely their own diameters and at or near to filling their roche lobes.The masses of A1a and A1b determined from the orbital parameters are 116 ± 31 M☉ and 89 ± 16 M☉respectively. This makes them the two most massive stars directly measured, i.e. with their masses determined (using Keplerian orbits), and not estimated from models. The masses estimated from analysis of the physical properties are slightly higher at 120 M☉ and 92 M☉.

Each component is a Wolf-Rayet (WR) star, with spectra dominated by strong broadened emission lines. Type WN6 indicates that ionised nitrogen lines are strong in comparison to ionised carbon lines, and the suffix h indicates that hydrogen is also seen in the spectrum. This type of WR star is not the classical stripped helium-burning aged star, but a young highly luminous object with CNO cycle fusion products showing at the surface due to strong conventional and rotational mixing, and high mass loss rates from the atmosphere. The emission lines are generated in the stellar wind and the photosphere is completely hidden. The surface fraction of hydrogen is still estimated to be 60-70%.Although the stars are very young, around 1.5 million years old, they have already lost a considerable fraction of their initial masses. The initial masses are estimated to have been 148 M☉ and 106 M☉, meaning they have lost 28 M☉ and 14 M☉ respectively.

NGC 3603-B

NGC 3603-B (HD 97950B) is a Wolf-Rayet star located at the centre of the HD 97950 cluster in the NGC 3603 star-forming region, about 20,000 light years from Earth. It has the spectral type WN6h and is among the most luminous and most massive stars known.

HD 97950 was catalogued as a star, but was known to be a dense cluster or close multiple star. In 1926, the six brightest members were given letters from A to F, although several of them have since been resolved into more than one star. Star B turned out to be the brightest single star.HD 97950B is a Wolf-Rayet (WR) star, with spectra dominated by strong broadened emission lines. Type WN6 indicates that ionised nitrogen lines are strong in comparison to ionised carbon lines, and the suffix h indicates that hydrogen is also seen in the spectrum. This type of WR star is not the classical stripped helium-burning aged star, but a young highly luminous object with CNO cycle fusion products showing at the surface due to strong conventional and rotational mixing, and high mass loss rates from the atmosphere. The emission lines are generated in the stellar wind and the photosphere is completely hidden. The surface fraction of hydrogen is still estimated to be around 60%.HD 97950B is the most massive and most luminous star known in the NGC 3603 region, nearly three million times more luminous than the sun and 132 times more massive. Although the star is very young, around 1.5 million years old, it has already lost a considerable fraction of its initial masses. The initial mass is estimated to have been 166 M☉, meaning it has lost 34 M☉.

NGC 3603-C

NGC 3603-C (HD 97950C) is a single-lined spectroscopic binary star system located at the centre of the HD 97950 cluster in the NGC 3603 star-forming region, about 20,000 light years from Earth. The primary has spectral type WN6h and is among the most luminous and most massive known.

HD 97950 was catalogued as a star, but was known to be a dense cluster or close multiple star. In 1926, the six brightest members were given letters from A to F, although several of them have since been resolved into more than one star. Star C was shown to be a binary, but its companion has not been observed.HD 97950C is a Wolf-Rayet (WR) star, with spectra dominated by strong broadened emission lines. Type WN6 indicates that ionised nitrogen lines are strong in comparison to ionised carbon lines, and the suffix h indicates that hydrogen is also seen in the spectrum. This type of WR star is not the classical stripped helium-burning aged star, but a young highly luminous object with CNO cycle fusion products showing at the surface due to strong conventional and rotational mixing, and high mass loss rates from the atmosphere. The emission lines are generated in the stellar wind and the photosphere is completely hidden. The surface fraction of hydrogen is still estimated to be 70%.The two component stars of NGC 3603-C circle each other every nine days. It is assumed that the secondary is sufficiently smaller and fainter than the primary not to affect the calculation of its physical properties. The mass is estimated to be 113 M☉ and the luminosity over two million L☉. Although the star is very young, around 1.5 million years old, it has already lost a considerable fraction of its initial mass. The initial mass is estimated to have been 137 M☉, meaning it has lost 24 M☉.

Nitrogen-13

Nitrogen-13 is a radioisotope of nitrogen used in positron emission tomography (PET). It has a half-life of a little under ten minutes, so it must be made at the PET site. A cyclotron may be used for this purpose.

Nitrogen-13 is used to tag ammonia molecules for PET myocardial perfusion imaging.

R136a3

R136a3 is a Wolf–Rayet star in R136, a massive star cluster located in Dorado. It is located near R136a1, the most massive and luminous star known. R136a3 is itself one of the most massive and most luminous stars known at 180 times more massive and 3.8 million times more luminous than the Sun.

The formal name of the star is RMC 136a3, standing for Radcliffe observatory, Magellanic Clouds, 136a3. The RMC survey identified luminous objects in the Large Magellanic Cloud and one of the brightest was RMC 136. This is now commonly shortened to R136, which is now known to be an extremely young dense open cluster at the core of the NGC 2070 cluster in the Tarantula Nebula. R136 was eventually resolved and the brightest "star" at the centre was termed R136a. This was further resolved into multiple components, one of which is R136a3.

Although R136a3 has a Wolf-Rayet spectral type dominated by intense emission lines of helium and nitrogen, usually indicating a highly evolved star that has lost its outer layers, R136a3 is actually an extremely young star. The spectrum also includes hydrogen lines and analysis shows the star is still 40% hydrogen at the surface. The helium and nitrogen in the atmosphere of such a young star are caused by strong convection due to the massive core and intense CNO cycle fusion, enhanced further by rotational mixing. The emission lines in the spectrum indicate strong mass loss caused by the fusion products at the surface and the enormous luminosity.

R136b

R136b is a Wolf-Rayet star in the R136 cluster in the Large Magellanic Cloud. It is one of the most massive and most luminous stars known. It is found in the dense R136 open cluster at the centre of NGC 2070 in the Tarantula Nebula.

R136b has a transitional spectral type between an O class supergiant and a Wolf-Rayet star, with a mix of absorption and emission lines. Although it shows enhanced helium and nitrogen at its surface, it is still a very young star, still burning hydrogen in its core via the CNO cycle, and still effectively a main sequence object.

Red-giant branch

The red-giant branch (RGB), sometimes called the first giant branch, is the portion of the giant branch before helium ignition occurs in the course of stellar evolution. It is a stage that follows the main sequence for low- to intermediate-mass stars. Red-giant-branch stars have an inert helium core surrounded by a shell of hydrogen fusing via the CNO cycle. They are K- and M-class stars much larger and more luminous than main-sequence stars of the same temperature.

Solar core

The core of the Sun is considered to extend from the center to about 0.2 to 0.25 of

solar radius. It is the hottest part of the Sun and of the Solar System. It has a density of 150 g/cm3 (150 times the density of liquid water) at the center, and a temperature of 15 million kelvins (15 million degrees Celsius, 27 million degrees Fahrenheit). The core is made of hot, dense plasma (ions and electrons), at a pressure estimated at 265 billion bar (3.84 trillion psi or 26.5 petapascals (PPa)) at the center. Due to fusion, the composition of the solar plasma drops from 68-70% hydrogen by mass at the outer core, to 33% hydrogen at the core/Sun center.

The core inside 0.20 of the solar radius contains 34% of the Sun's mass, but only 0.8% of the Sun's volume. Inside 0.24 solar radius, the core generates 99% of the fusion power of the Sun. There are two distinct reactions in which four hydrogen nuclei may eventually result in one helium nucleus: the proton-proton chain reaction – which is responsible for most of the Sun's released energy – and the CNO cycle.

Standard solar model

The standard solar model (SSM) is a mathematical treatment of the Sun as a spherical ball of gas (in varying states of ionisation, with the hydrogen in the deep interior being a completely ionised plasma). This model, technically the spherically symmetric quasi-static model of a star, has stellar structure described by several differential equations derived from basic physical principles. The model is constrained by boundary conditions, namely the luminosity, radius, age and composition of the Sun, which are well determined. The age of the Sun cannot be measured directly; one way to estimate it is from the age of the oldest meteorites, and models of the evolution of the Solar System. The composition in the photosphere of the modern-day Sun, by mass, is 74.9% hydrogen and 23.8% helium. All heavier elements, called metals in astronomy, account for less than 2 percent of the mass. The SSM is used to test the validity of stellar evolution theory. In fact, the only way to determine the two free parameters of the stellar evolution model, the helium abundance and the mixing length parameter (used to model convection in the Sun), are to adjust the SSM to "fit" the observed Sun.

Stellar core

A stellar core is the extremely hot, dense region at the center of a star. For an ordinary main sequence star, the core region is the volume where the temperature and pressure conditions allow for energy production through thermonuclear fusion of hydrogen into helium. This energy in turn counterbalances the mass of the star pressing inward; a process that self-maintains the conditions in thermal and hydrostatic equilibrium. The minimum temperature required for stellar hydrogen fusion exceeds 107 K (10 MK), while the density at the core of the Sun is over 100 g cm−3. The core is surrounded by the stellar envelope, which transports energy from the core to the stellar atmosphere where it is radiated away into space.

Stellar nucleosynthesis

Stellar nucleosynthesis is the theory explaining the creation (nucleosynthesis) of chemical elements by nuclear fusion reactions between atoms within stars. Stellar nucleosynthesis has occurred continuously since the original creation of hydrogen, helium and lithium during the Big Bang. It is a highly predictive theory that today yields excellent agreement between calculations based upon it and the observed abundances of the elements. It explains why the observed abundances of elements in the universe grow over time and why some elements and their isotopes are much more abundant than others. The theory was initially proposed by Fred Hoyle in 1946, who later refined it in 1954. Further advances were made, especially to nucleosynthesis by neutron capture of the elements heavier than iron, by Margaret Burbidge, Geoffrey Burbidge, William Alfred Fowler and Hoyle in their famous 1957 B2FH paper, which became one of the most heavily cited papers in astrophysics history.

Stars evolve because of changes in their composition (the abundance of their constituent elements) over their lifespans, first by burning hydrogen (main sequence star), then helium (red giant star), and progressively burning higher elements. However, this does not by itself significantly alter the abundances of elements in the universe as the elements are contained within the star. Later in its life, a low-mass star will slowly eject its atmosphere via stellar wind, forming a planetary nebula, while a higher–mass star will eject mass via a sudden catastrophic event called a supernova. The term supernova nucleosynthesis is used to describe the creation of elements during the evolution and explosion of a pre-supernova massive star (12–35 times the mass of the sun). Those massive stars are the most prolific source of new isotopes from carbon (Z = 6) to nickel (Z = 28).

The advanced sequence of burning fuels is driven by gravitational collapse and its associated heating, resulting in the subsequent burning of carbon, oxygen and silicon. However, most of the nucleosynthesis in the mass range A = 28–56 (from silicon to nickel) is actually caused by the upper layers of the star collapsing onto the core, creating a compressional shock wave rebounding outward. The shock front briefly raises temperatures by roughly 50%, thereby causing furious burning for about a second. This final burning in massive stars, called explosive nucleosynthesis or supernova nucleosynthesis, is the final epoch of stellar nucleosynthesis.

A stimulus to the development of the theory of nucleosynthesis was the discovery of variations in the abundances of elements found in the universe. The need for a physical description was already inspired by the relative abundances of isotopes of the chemical elements in the solar system. Those abundances, when plotted on a graph as a function of atomic number of the element, have a jagged sawtooth shape that varies by factors of tens of millions (see history of nucleosynthesis theory). This suggested a natural process that is not random. A second stimulus to understanding the processes of stellar nucleosynthesis occurred during the 20th century, when it was realized that the energy released from nuclear fusion reactions accounted for the longevity of the Sun as a source of heat and light.

Timeline of stellar astronomy

Timeline of stellar astronomy

2300 BC — First great period of star naming in China.

134 BC — Hipparchus creates the magnitude scale of stellar apparent luminosities

185 AD — Chinese astronomers become the first to observe a supernova, the SN 185

964 — Abd al-Rahman al-Sufi (Azophi) writes the Book of Fixed Stars, in which he makes the first recorded observations of the Andromeda Galaxy and the Large Magellanic Cloud, and lists numerous stars with their positions, magnitudes, brightness, and colour, and gives drawings for each constellation

1000s (decade) — The Persian astronomer, Abū Rayhān al-Bīrūnī, describes the Milky Way galaxy as a collection of numerous nebulous stars

1006 — Ali ibn Ridwan and Chinese astronomers observe the SN 1006, the brightest stellar event ever recorded

1054 — Chinese and Arab astronomers observe the SN 1054, responsible for the creation of the Crab Nebula, the only nebula whose creation was observed

1181 — Chinese astronomers observe the SN 1181 supernova

1580 — Taqi al-Din measures the right ascension of the stars at the Constantinople Observatory of Taqi ad-Din using an "observational clock" he invented and which he described as "a mechanical clock with three dials which show the hours, the minutes, and the seconds"

1596 — David Fabricius notices that Mira's brightness varies

1672 — Geminiano Montanari notices that Algol's brightness varies

1686 — Gottfried Kirch notices that Chi Cygni's brightness varies

1718 — Edmund Halley discovers stellar proper motions by comparing his astrometric measurements with those of the Greeks

1782 — John Goodricke notices that the brightness variations of Algol are periodic and proposes that it is partially eclipsed by a body moving around it

1784 — Edward Pigott discovers the first Cepheid variable star

1838 — Thomas Henderson, Friedrich Struve, and Friedrich Bessel measure stellar parallaxes

1844 — Friedrich Bessel explains the wobbling motions of Sirius and Procyon by suggesting that these stars have dark companions

1906 — Arthur Eddington begins his statistical study of stellar motions

1908 — Henrietta Leavitt discovers the Cepheid period-luminosity relation

1910 — Ejnar Hertzsprung and Henry Norris Russell study the relation between magnitudes and spectral types of stars

1924 — Arthur Eddington develops the main sequence mass-luminosity relationship

1929 — George Gamow proposes hydrogen fusion as the energy source for stars

1938 — Hans Bethe and Carl von Weizsäcker detail the proton-proton chain and CNO cycle in stars

1939 — Rupert Wildt realizes the importance of the negative hydrogen ion for stellar opacity

1952 — Walter Baade distinguishes between Cepheid I and Cepheid II variable stars

1953 — Fred Hoyle predicts a carbon-12 resonance to allow stellar triple alpha reactions at reasonable stellar interior temperatures

1961 — Chūshirō Hayashi publishes his work on the Hayashi track of fully convective stars

1963 — Fred Hoyle and William A. Fowler conceive the idea of supermassive stars

1964 — Subrahmanyan Chandrasekhar and Richard Feynman develop a general relativistic theory of stellar pulsations and show that supermassive stars are subject to a general relativistic instability

1967 — Eric Becklin and Gerry Neugebauer discover the Becklin-Neugebauer Object at 10 micrometres

1977 — (May 25) The Star Wars film is released and became a worldwide phenomenon, boosting interests in stellar systems.

2012 — (May 2) First visual proof of existence of black-holes. Suvi Gezari's team in Johns Hopkins University, using the Hawaiian telescope Pan-STARRS 1, publish images of a supermassive black hole 2.7 million light-years away swallowing a red giant.

WR 24

WR 24 (HD 93131) is a Wolf-Rayet star in the constellation Carina. It is one of the most luminous stars known. At the edge of naked eye visibility it is also one of the brightest Wolf Rayet stars in the sky.

The spectrum of WR 24 has the characteristic strong nitrogen and helium emission lines of a WN star, but also lines of hydrogen that show Doppler-displaced absorption components. The lowest ionisation nitrogen emission lines are strongest, with NV lines being very weak. The HeI lines are weaker than the HeII lines, leading to a WN6ha spectral class. The spectral type is annotated with a letter w, indicating weaker emission than for a typical WN6 star.WR 24 is thought to be a member of the open cluster Collinder 228, sometimes considered to be just an extension of the rich cluster Trumpler 16. It lies on the southwestern side of the Carina Nebula. Collinder 228 and the Carina Nebula are approximately 2.2 kpc away. However, the Gaia Data Release 2 parallax gives a distance around 4200 for WR 24.WR 24 has been reported to vary in brightness by about 0.02 magnitudes. Analysis of Hipparcos photometry shows an amplitude of 0.082 magnitudes and a primary period of 4.76 days. It has not yet been assigned a variable star designation in the General Catalogue of Variable Stars and is still formally listed as a suspected variable.The hydrogen-rich WN stars have been referred to as WNL stars or as WNH stars since they do not necessarily have late nitrogen-sequence spectra. They are systematically more massive and more luminous than stars with similar spectra but lacking nitrogen. WR 24 has a mass of 54 M☉ and is over two million times as luminous as the sun. These stars are proposed to be young hydrogen-burning stars, effectively main sequence objects, rather than post-supergiant stars. WR 24 is calculated to have 44% hydrogen in its atmosphere. The cluster Collinder 228 is thought to be around 6.78 million years old. The WR-type spectra are caused because helium and nitrogen and convected to the surface by the extreme temperature gradients caused by the CNO cycle in the core, and then expelled by powerful stellar winds. WR 24 has a wind reducing its mass by 40×10−6 M☉ per year, at a velocity of 2,160 km/s.

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