Download Fluorine nucleosynthesis

The sites and mechanisms of galactic
fluorine synthesis
Claudio Ugalde
Does she know where fluorine
come from in the galaxy??
(Neither we!)
Facts about fluorine
Our teeth are made of
hydroxylapatite (Ca10(PO4)6(OH)2)
By adding MFP (sodium monofluorophosphate)
to toothpaste, a chemical reaction may occur on
contact with teeth.
Fluorapatite = H2O + MFP + hydroxylapatite
(MFP is toxic, though!)
By adding MFP (sodium monofluorophosphate)
to toothpaste, a chemical reaction may occur on
contact with teeth.
Fluorapatite = H2O + MFP + hydroxylapatite
(MFP is toxic, though!)
Fluorapatite (Ca5(PO4)3F) is a very
hard material. However, it has a yellowbrownish hue.
Excess fluoride
may cause mottled enamel in teeth
Fluorine is important for making nuclear weapons. Isotopes of
uranium are separated by diffusion of UF6.
Fluoro-chloro hydrocarbons used in refrigeration destroy the
ozone layer in the atmosphere
Fluorine is highly toxic
Fluorine stands out as the most reactive of all elements. It
forms molecules even with noble gases. It behaves as the
predator element in the periodic table.
Facts about fluorine II
It has only one stable isotope: 19F.
Fluorine is by far the least abundant of all
11<A<33 stable nuclei in the solar system.
It is difficult to identify fluorine in stars, as HF does not show
strong absorption lines in the visible. Infrared-equiped
telescopes are needed.
Either fluorine is very fragile when inside stars,
or it is very hard to make.
Trivia
Where is fluorine synthesized?
a) Type II supernovae
b) Wolf-Rayet stars
c) AGB stars
d) All of the above
Trivia
Where is fluorine synthesized?
a) Type II supernovae
b) Wolf-Rayet stars
c) AGB stars
d) All of the above
Answer: d
Fluorine production in Type II supernovae
Proposed by Woosley & Haxton in 1988
Fluorine could be produced in the neon-rich
shell of exploding massive stars (SNII) via
weak interactions: the ν-process.
So far, fluorine has not been observed in SN
remnants.
SN1987a, HST
Thermonuclear burning
As temperature increases, heavier nuclei can
overcome the strong Coulomb repulsion between
them.
H →4He →12C →20Ne→16O→28Si
→56Ni (β+) 56Co (β+) 56Fe
1
Beyond 56Fe, thermonuclear
burning is not energy efficient.
Q<0
The silicon shell
The burning time is very short (~1 day) compared to other
burning stages.
Hydrostatic silicon burning occurs at T = 3.5 GK
At this temperature, photodisintegrations become as likely as
heavy-ion induced reactions. Nuclear statistical equilibrium
(NSE) is established for electromagnetic and strong
interaction-induced reactions. Light ions are abundant too.
The heaviest nuclei produced by NSE are deposited on the
core.
Surface
O, Mg, Si
Si, S + light ions
Fe, Ni, Cr
Si shell (NSE)
Core
The Fe core (before collapse)
Energy can not be released from the core by thermonuclear
reactions (as they are endothermic). Therefore, the core starts
contracting and the density is increased.
The Fe core (before collapse)
Energy can not be released from the core by thermonuclear
reactions (as they are endothermic). Therefore, the core starts
contracting and the density is increased.
The core becomes electron degenerate. As long as the mass
of the core remains below the Chandrasekhar limit, electron
pressure can support contraction by gravity.
The Fe core (before collapse)
Energy can not be released from the core by thermonuclear
reactions (as they are endothermic). Therefore, the core starts
contracting and the density is increased.
The core becomes electron degenerate. As long as the mass
of the core remains below the Chandrasekhar limit, electron
pressure can support contraction by gravity.
Any cooling process in the core would lead to further
contraction. Contraction increases density, so the Fermi
energy is increased as a result. Other cooling starts.
p + e- −−> n + νe
The Fe core (before collapse)
Energy can not be released from the core by thermonuclear
reactions (as they are endothermic). Therefore, the core starts
contracting and the density is increased.
The core becomes electron degenerate. As long as the mass
of the core remains below the Chandrasekhar limit, electron
pressure can support contraction by gravity.
Any cooling process in the core would lead to further
contraction. Contraction increases density, so the Fermi
energy is increased as a result. Other cooling starts.
p + e- −−> n + νe
Meanwhile, iron keeps accumulating from Si
burning. The Chandrasekhar mass is reached.
The collapse of the core
The core decouples from the rest of the star. It also splits itself into an
“homologous” and outer cores.
Matter falls while neutronization is ongoing in the homologous region. The
environment gets rich with neutrons and electron neutrinos.
Collapse stops when matter reaches nuclear density. Matter bounces off
the stiff neutron star and an outgoing shock wave carries out some energy.
To surface
Si burning shell
outer core
neutrino photosphere
Neutronization
ν trapping zone
Inner core
The collapse of the core
The core decouples from the rest of the star. It also splits itself into an
“homologous” and outer cores.
Matter falls while neutronization is ongoing in the homologous region. The
environment gets rich with neutrons and electron neutrinos.
Collapse stops when matter reaches nuclear density. Matter bounces off
the stiff neutron star and an outgoing shock wave carries out some energy.
To surface
Si burning shell
outer core
neutrino photosphere
Neutronization
ν trapping zone
Inner core
In the neutrino trapping zone, the diffusion
time exceeds the collapse time, so neutrinos
fall with matter. Cooling is stopped and ν's are
thermalized.
The ν-sphere corresponds to one mean free
path length.
The inner core contains neutrinos in
equilibrium with matter
The neutrino flux
Neutrinos will carry away most of the energy released in the
collapse (3 x1053 erg). Only 1051 erg come as photons or KE.
There are two main mechanisms of neutrino production:
The neutrino flux
Neutrinos will carry away most of the energy released in the
collapse (3 x1053 erg). Only 1051 erg come as photons or KE.
There are two main mechanisms of neutrino production:
a) neutronization (occurs during collapse in the homologous
core)
p + e- −−> νe + n
5%
b) deleptonization (during the cooling of the neutron star)
e+ + e- −−> νi + νi
95%
The neutrino flux
Neutrinos will carry away most of the energy released in the
collapse (3 x1053 erg). Only 1051 erg come as photons or KE.
There are two main mechanisms of neutrino production:
a) neutronization (occurs during collapse in the homologous
core)
p + e- −−> νe + n
5%
b) deleptonization (during the cooling of the neutron star)
e+ + e- −−> νi + νi
95%
Three flavors are produced (νe, νµ, ντ). Electron neutrinos
interact with matter via charged or neutral currents. Other
flavors scatter only via neutral currents. This leads to
thermalization of νe's
The neutrino flux (cont.)
Electron neutrinos with reduced energy (from thermalization
during the collapse) will escape from the core as their mean
free path is longer:
σν ~ E2
Bruenn 1987
The neutrino flux (cont.)
Electron neutrinos with reduced energy (from thermalization
during the collapse) will escape from the core as their mean
free path is longer:
σν ~ E2
Bruenn 1987
Other flavors remain
trapped until the neutron
star slowly contracts,
heats, and loses energy.
The neutrino interactions
The thermonuclear burning shells in the star first learn that
something wrong has happened with the star when neutrinos
from neutronization irradiate them in hydrostatic conditions
(pre-processing).
The neutrino interactions
The thermonuclear burning shells in the star first learn that
something wrong has happened with the star when neutrinos
from neutronization irradiate them in hydrostatic conditions
(pre-processing).
Soon after, the bounced shock front pushes the shells and the
hotter µ and τ−neutrinos start irradiating them
hydrodynamically while they explode.
Mean ν energies:
νe =15 MeV, νµ,τ=30 MeV
The neutrino interactions
The thermonuclear burning shells in the star first learn that
something wrong has happened with the star when neutrinos
from neutronization irradiate them in hydrostatic conditions
(pre-processing).
Soon after, the bounced shock front pushes the shells and the
hotter µ and τ−neutrinos start irradiating them
hydrodynamically while they explode.
Mean ν energies:
νe =15 MeV, νµ,τ=30 MeV
At these energies, µ and τ neutrinos can excite giant
resonance transitions in nuclei via neutral current weak
interactions. σ~10-42 cm2
The ν process
(aka Inelastic neutral current neutrino scattering)
(Z,A) + ν −−−> (Z,A)* + ν' −−−> (Z,A-1) + n + ν'
−−−> (Z-1,A-1) + p + ν'
−−−> (Z-2,A-4) + α + ν'
(neutrinos are mainly µ and τ )
The ν process
(aka Inelastic neutral current neutrino scattering)
(Z,A) + ν −−−> (Z,A)* + ν' −−−> (Z,A-1) + n + ν'
−−−> (Z-1,A-1) + p + ν'
−−−> (Z-2,A-4) + α + ν'
(neutrinos are mainly µ and τ )
The conditions:
a) The neutrino flux must be high enough so a
significant number of nuclei are excited.
b) The neutrino flux should be small enough so new
nuclei are not destroyed before they can escape via the
explosion mechanism.
Fluorine and ν's
The neon shell fulfills these conditions!!!
20Ne+
ν −−−> 20Ne* + ν' −−−> 19Ne + n + ν'
30%
−−−> 19F + p + ν'
66%
−−−> 16O + α + ν'
4%
Fluorine and ν's
The neon shell fulfills these conditions!!!
20Ne+
ν −−−> 20Ne* + ν' −−−> 19Ne + n + ν'
30%
−−−> 19F + p + ν'
66%
−−−> 16O + α + ν'
4%
The destruction mechanisms are:
19F(p,α)16O
19F(γ,α)15N
However, fluorine
from supernovae
remnants has never
been observed
Fluorine from Wolf-Rayet stars
WR124,
WR HST
124,
HST
(by Meynet and Arnould in 1993)
Wolf-Rayet stars
WR's are hot blue giant stars (>25 solar masses)
that lose mass via strong radiation-driven winds.
19
22
During the main sequence,
F  these
α ,p  Nestars may have
had masses ~ 80 Msol
Their surface composition is very exotic, as mass loss
is able to uncover the nucleosynthetic products. Helium
(instead of H) dominates their spectra.
Mass loss is more efficient for stars
with solar-like abundance than those
metal-poor.
WR18
Fluorine from Wolf-Rayet stars
19F(α,p)22Ne
18F(βν)18O(p,α)15N(α,γ)19F
14N(α,γ)18F
19F(n,γ)20F
19F(p,α)16O
18F(α,p)21Ne
WR 124,
HST
Fluorine synthesis takes place during
the early He-burning core phase
in massive stars. However, the
reaction would destroy most of the
newly produced fluorine. Therefore,
an efficient way of saving fluorine
before the end of the He-burning phase
is needed. Mass ejection
by Wolf-Rayets may do the job.
In AGB stars
Proposed by Goriely, Jorrisen, and Arnould in 1989
Fluorine is produced mainly in
the He-burning phase
of AGB stars in the same
site where the s process
occurs.
Fluorine in AGB stars has been
observed by Jorrisen, Smith and
Lambert in 1992.
NGC 6543,HST
AGB stars
M3, NOAO
Karakas, Ph.D. Thesis 2003
The AGB Star
He-burning shell
Jorrisen et al. 1992
C-O core
He intershell
H-burning shell
Convective envelope
Jorrisen, Smith & Lambert, 1992
* Observations include K, M, S,
MS, SC, N, J, and Ba stars.
* Fluorine was identified through
vibration-rotation lines of HF
* F vs. C abundances show a
correlation. (See F/O vs C/O figure)
* To disentangle the effects of
galactic chemical evolution on F
from stellar evolutionary effects,
F and C abundances should be
normalized by a species sensitive
to galactic chemical evolution only
and not to nucleosynthesis during
the AGB. Iron-group species or
oxygen are good choices.
* Finally, JS&L proved that the
correlation between F and C is not
a product of normalization by O.
(See O vs C/O figure)
13C(α,n)16O
14N(n,p)14C
19F(α,p)22Ne
18F(βν)18O(p,α)15N(α,γ)19F
14N(α,γ)18F
19F(n,γ)20F
19F(p,α)16O
18F(α,p)21Ne
Convective
envelope H burning He burning
C-O core
H mixing
Convective Pulse
pocket
He intershell
19
F(α,p)22Ne
Current status of the 19F(α,p)22Ne rate
The only experimental work available for this
reaction at low energies dates back to 1965.
(Kuperus, J., Physica 31 (1965) 1603.)
The error bars in the rate span 15 ORDERS OF
MAGNITUDE at temperatures of relevance to
fluorine synthesis in AGB stars !
(No wonder why this rate needed to be measured)
Data by Kuperus, Physica 31,1603-1616 (1965).
The Experiments
4 MV KN van de Graaff accelerator at Notre Dame
Photon detection experiment
● * We used a Ge detector and a BGO scintillator
placed at 55 degrees with respect to the beam for
detecting gammas.
● * We calibrated detectors’ efficiency with a Cobalt-60
source.
● * Room-background shielding of the detector was
provided by lead bricks.
● * Calcium fluoride targets were evaporated on
tantalum backings. The thickness of the targets was about
30 keV.
● * The observed yield was corrected to thick-target
yields to calculate the resonance strength.
● * Strong target deterioration was observed. We
corrected the yields by scanning through the
strongest resonance several times.
Excitation Curve
Ge
Yield (gammas/alpha)
BGO
1871
1660
1552
1490
1452
1396
1363
1321
1260
Alpha energy (keV)
Charged-particle experiments
Charged-particle experiment
● * The α-beam was produced with the University of Notre Dame KN
Van de Graaff accelerator.
● * CaF2 targets were evaporated on 20 µg/cm2 carbon backings.
● * We used the Ortec scattering chamber for the experiment.
● * Three Si detectors were placed very close to the reaction place. Ni
foils in front of the detectors helped to prevent elastic-scattered alphas
from being detected.
● * A collimated Si detector at 160o was used to monitor the target
content at all times.
● * We measured the excitation functions.
● * We got angular distributions.
● * The energy range explored in the experiment was limited by the
target stability.
Eα(keV)
ωγ (keV)
Eα(keV)
ωγ (keV)
And now, the nasty side of the story...
FLUORINE TARGETS
Remember fluorine is the
predator of the elements?
Fluorine reacts with everything around it. So it is of no
surprise that fluorine in targets is easily lost.
Trying to measure cross sections at lower energies means
increasing beam intensity and exposing the targets
to beam for a longer time. :(
We needed to find some way of preparing stable fluorine
targets!
Implanted targets
●
●
●
●
An implantation station was improvised from a pumping cross at the
entrance of the Tandem accelerator.
A fluorine beam was produced by the SNICS ion source at Notre Dame.
Negative beam currents of up to 20 microAmps were obtained from a
calcium fluoride cathode.
The beam was separated with an analyzing magnet coupled to the ion
source.
FN Tandem
Implantation
station
SNICS II
Target implantation (first try)
Due to the success of the charged particle
experiment we decided to implant fluorine on a carbon foil
substrate.
Failure!rgets?
UGLY
(Foils were literally chewed by heavy fluorine beam)
Target implantation (second try)
We concluded that target cooling would have helped in our
previous method for making a target.
Cooling had to be implemented at the implantation station.
This required water to flow inside the beam line just
outside the FN accelerator and above a turbo pump.
Water and turbo pumps do not happen to be a good
match. People got a little nervous about it.
Failure 2
(And we had an improved version this time!)
HORRIBLE failure
Target implantation (third try)
We had enough with carbon foils, so we redesigned the experiment by
moving on to solid targets.
The natural choice for substrate was tantalum. We implanted a 1.5” by 1.5”
Ta piece in our improvised implantation station.
Profile of implanted targets
We learned we had to wobble the beam and keep the
beam energy fixed. (Thanks to J. Kaiser for the wobbler!)
Saturation curve for implanted targets
We implanted tantalum substrates with a 70 keV wobbled beam. The
energy was kept constant during the implantation process. After we got
profiles of the targets we jumped to the real experiment.
Charged-particle experiment (low
energy)
FAILURE!
Photo by Rick Roberts
Electron microscope scan of a tantalum substrate after
6 hours of 7uA alpha beam (yes, it was water-cooled)
Target implantation (fourth try)
Tantalum failed in the sense that we were not going to be
able to put the required amount of charge on the targets for
our experiment. Besides, the targets were extremely thick to
get useful information from the experiment.
However, we saw a great improvement from previous targets.
We thought we were on the right track, so we started looking
for a better substrate.
We tried nickel, molybdenum, chromium, iron, gold... even
aluminum!
We baked materials, and polished and sand-papered surfaces.
We cleaned with alcohol, acetone, water, soap, HNO3, etc.
(This time we had a beam wobbler
and a fixed beam energy.)
Trivia question
Did any of these methods work?
YES!
Implantation of targets
●
●
●
We selected Fe as a substrate for our targets.
Targets were tilted at 75 degrees with respect to the fluorine beam. This reduced the
depth of implanted fluorine in the substrate.
The beam was wobbled in order to get a uniform implanted surface.
The implantation energy was 36 keV. The energy of the beam produced by SNICS
could not be reduced further as the maximum beam intensity was proportional to the
bias voltage. Each target took from 12 to 24 hours to make.
1000
fluorine content (arbitrary units)
●
100
10
1
1
2
3
4
5
6
7
8
target
9
10
11
12
13
Evap
Last experiment
We used the γ-ray beam line coupled to the KN accelerator.
●We designed a new solid-target scattering chamber.
●The chamber could hold up to two silicon detectors at a time in close
geometry to the target.
●Detectors were shielded with Al foils to prevent elastic scattered
alphas from being registered.
●Two MCAs were used to acquire spectra.
●The implanted targets were water-cooled and isolated from the
chamber.
●Yields were obtained by normalizing detector counts to integrated
beam current on targets.
●The chamber was isolated from the beam line to prevent electronic
noise from appearing in spectra.
●A copper plate in the chamber was cooled down to liquid nitrogen
temperature. temperature to prevent carbon build-up on the target.
●
1360 keV
p0
p1
1100 keV
p1
p0
792 keV
F(α,p0)22Ne
19
F(α,p1)22Ne
19
The reaction rate
〈συ 〉=

8
πμ  kT 
∞
3
∫0
 
E
σ  E  E exp −
dE
kT
T is the temperature of the plasma
E is the energy of the particle pair
σ(E) is the integrated cross section
But...
Problems : Coulomb barrier prevents us from
measuring the reaction cross section at small energies.
Therefore, the main goal here becomes to extrapolate the
cross section into the Gamow window.
Are there more resonances inside the Gamow window?
(We may get an idea if we look into the nuclear structure
of the compound) What are their properties?
Are there non-resonant contributions to the cross section?
Also, sometimes the number of parameters (energies of
resonances + reduced width amplitudes) is huge.
The theory
The 2-step model for low energy nuclear reactions
19
Compound
F
Exit channel
α
Entrance channel Step 1
V(r)
Potential
23
Step 2
Na
p
22
Ne
Coulomb+centrifugal
Point
r
Nuclear
22
The compound
Ne
Compound
?
26
In fact, we don't know what happens to the
nucleons during the formation of the compound.
The energy of the system is distributed among all
the nucleons.
Mg
The compound “looses memory” of the way in which it was
formed.
Basic rules still apply: conservation of energy,
angular momentum, charge, etc. Whatever happens
to the compound forward in time needs to follow the
rules.
Most interesting is that the process of formation of the
compound is time reversal symmetric !
Formation
Destruction
The Wigner hypersurface
Compound
26
Mg
The surface splits space in two:
R
a) Inside- where ALL nuclear
reactions between the pair of
nuclei take place
b) Outside-everything else
R can have any size as long as all reactions take place inside the surface.
The model restricts R to be finite. A very large R (say the size of a “finite” universe) is possible but
computations get extremely complex. In practice R < 10 fm.
Wigner chose a truncated octahedron to describe the
boundary (for historical reasons, irrelevant to the theory).
In general, the boundary is an hypersurface in a 3A
dimensional space, such that A is the number of
nucleons in the projectile+target system.
Each dimension corresponds to a spatial coordinate.
Each face of the hypersurface is called a channel.
A channel is one of the many ways the compound can be formed (or destroyed).
A channel c is defined by c = c{α(I1I2)sνlm}
α is the particle pair
I1 and I2 are the spins of the 2 particles
s is the channel spin s=I1+I2 and ν its projection
l is the orbital angular momentum of the 2 particles
and m its projection
Finding an initial set of R-matrix parameters
(needs to be done by hand)
1) Try to restrict the N space as much as possible. (Basically, answer the question
“How much we know about the compound?”)
2) Select the levels that should have a strong influence in the measured curves.
3) Set by hand the energies of these levels. Get peaks at the right position.
4) Turn off all resonances but the ones for a single Jπ.
5) Within a single Jπ, work in pairs trying to figure out how one resonance affects the others in the
group. Try to figure out what are the strongest conditions in the group (signs of reduced width
amplitudes + their absolute value) governing a “reasonable trend”
6) Once the signs of the reduced width amplitudes are set, turn on 2 groups
of Jπ's. Work for all possible pairs of Jπ's.
7) Turn on all Jπ's, changing one of the N parameters + signs, one at a time.
8) A small variation in one of the N parameters affects all the curves at the same time
(this is independent of the method).
9) The method is iterative and therefore very time-consuming. This means that all steps
in the fitting process need to retraced over and over again (3 to 5 times, as average).
Extrapolation to lower energies
From proton scattering experiments we got information about the compound nucleus
structure and proton widths.
But, what about α-widths?
γ α(J,π) = 10
2
2
<log(γ α)>
Interference between resonances
In the future, probably the most important sources of uncertainty in reaction
rates important to hydrogen and helium burning will be:
a) Fast, one step
processes (such
as direct captures)
b) Interference between
resonances
The effects of this
kind of uncertainty
needs to be simulated
with Monte Carlo
Reaction rate
Summary of fluorine destruction results from the Mount
Stromlo Stellar Structure Program (MSSSP).
1) In all models, the contribution to fluorine destruction stands as follows:
with CF88
(a,p) 50%, (n,g) 50%
with New rate (a,p) 10% (n,g) 90%
Maximum destruction by (a,p) is 17% in the m5z02
model
2) From CF88 to new rates F19 yields change by
m3z02
+24%
m3z008 +40%
<---- The largest producers
m2z0001 +43%
<---- of F19
m5z02
x3.7
3) By removing the (n,g) reaction artificially from the network we get
F19 yield changes by
m3z02
+18%
m3z008 +35%
~10% of the (n,g) destruction
m2z0001 +28%
occurs in the pocket
m5z02
x2.2
So where does fluorine come from?
So far fluorine has been observed in our solar system, in the LMC and omega Centauri
(Cunha et al. 2004), and in the Milky Way (Jorrisen et al. 1992). All observations are from AGB stars
and post-AGB stars (including their nebula).
Cunha et al. concluded that the main mechanism of fluorine production in the LMC is not the
AGB scenario. However, they were not able to include in their models the contribution from
Wolf-Rayet stars.
Renda et al., 2004 tested all three possible mechanisms of fluorine enrichment in the Milky
Way. By using our estimate of yields from AGB stars (Lugaro 2004), they concluded that
fluorine abundances in the Milky Way can be reproduced only if all three possible
mechanisms are considered at the same time.
Wolf-Rayet stars are very rare. However, it is worth to look for fluorine in their spectra.
It is puzzling that fluorine has not been observed in SNII. To be continued...
Thanks!
R. Azuma
A. Couture
J. Goerres
H. Y. Lee
E. Stech
E. Strandberg
W. Tan
M. Wiescher
A. Karakas
M. Lugaro
R. Stancliffe