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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