Classical-Nova Contribution to the Milky Way's 26 Al Abundance: Exit Channel of the Key 25 Al( p,γ ) 26 Si Resonance
M. B. Bennett, C. Wrede, K. A. Chipps, J. José, S. N. Liddick, M. Santia, A. Bowe, A. A. Chen, N. Cooper, D. Irvine, E. McNeice, F. Montes, F. Naqvi, R. Ortez, S. D. Pain, J. Pereira, C. Prokop, J. Quaglia, S. J. Quinn, S. B. Schwartz, S. Shanab, A. Simon, A. Spyrou, E. Thiagalingam
aa r X i v : . [ nu c l - e x ] D ec Classical-Nova Contribution to the Milky Way’s Al Abundance: Exit Channel of theKey Al( p, γ ) Si Resonance
M. B. Bennett,
1, 2, ∗ C. Wrede,
1, 2, 3, † K. A. Chipps, J. Jos´e, S. N. Liddick,
6, 2
M. Santia,
1, 2
A. Bowe,
1, 2, 7
A. A. Chen, N. Cooper, D. Irvine, E. McNeice, F. Montes,
2, 10
F. Naqvi, R. Ortez,
1, 2, 3
S. D. Pain, J. Pereira,
C. Prokop,
6, 2
J. Quaglia,
12, 10, 2
S. J. Quinn,
1, 2, 10
S. B. Schwartz,
1, 2, 13
S. Shanab,
1, 2
A. Simon,
2, 10
A. Spyrou,
1, 2, 10 and E. Thiagalingam Department of Physics and Astronomy, Michigan State University, East Lansing, Michigan 48824, USA National Superconducting Cyclotron Laboratory, Michigan State University, East Lansing, Michigan 48824, USA Department of Physics, University of Washington, Seattle, Washington 98195, USA Department of Physics, Colorado School of Mines, Golden, Colorado 08401, USA Departament F´ısica i Enginyeria Nuclear (UPC) and Institutd’Estudis Espacials de Catalunya (IEEC), E-08034 Barcelona, Spain Department of Chemistry, Michigan State University, East Lansing, Michigan 48824, USA Physics Department, Kalamazoo College, Kalamazoo, Michigan 49006, USA Department of Physics and Astronomy, McMaster University, Hamilton, ON L8S 4M1, Canada Department of Physics and Wright Nuclear Structure Laboratory,Yale University, New Haven, Connecticut 06520, USA Joint Institute for Nuclear Astrophysics, Michigan State University, East Lansing, Michigan 48824, USA Oak Ridge National Laboratory, Oak Ridge, Tennessee 37831, USA Department of Electrical Engineering, Michigan State University, East Lansing, Michigan 48824, USA Geology and Physics Department, University of Southern Indiana, Evansville, Indiana 47712, USA
Classical novae are expected to contribute to the 1809-keV Galactic γ -ray emission by producingits precursor Al, but the yield depends on the thermonuclear rate of the unmeasured Al( p, γ ) Sireaction. Using the β decay of P to populate the key J π = 3 + resonance in this reaction, we reportthe first evidence for the observation of its exit channel via a 1741 . ± . ± . γ ray, where the uncertainties are statistical and systematic, respectively. By combiningthe measured γ -ray energy and intensity with other experimental data on Si, we find the center-of-mass energy and strength of the resonance to be E r = 414 . ± . ± . ± . ωγ = 23 ± +11 − (lit.) meV, respectively, where the last uncertainties are from adoptedliterature data. We use hydrodynamic nova simulations to model Al production showing that thesemeasurements effectively eliminate the dominant experimental nuclear-physics uncertainty and weestimate that novae may contribute up to 30 % of the Galactic Al.
PACS numbers: 23.20.Lv, 25.40.Lw, 26.30.Ca, 27.30.+t
Gamma-ray telescopes pointed at the Milky Way havedetected a diffuse 1809-keV line that is characteristic of Al decay ( τ = 1 . Al mass of 2 . ± . M ⊙ [5, 6]. The abundance of Al in protoplanetary disksorbiting young stars may influence the formation of hab-itable planetary systems such as our own because, in suf-ficient quantities, the energy released by its in situ de-cay can heat planetesimals inducing differentiation andwater sublimation [7–9]. The inhomogeneous spatial dis-tribution of the 1809-keV emission across the Milky Waysuggests that the outflows of massive stars and their su-pernovae are the primary sites for Al production [10].In the limit where secondary sites such as classical no-vae and asymptotic giant branch stars are well under-stood, one can use the Al line to estimate the rate ofcore-collapse supernovae in the Milky Way [5] or comparewith the Fe gamma-ray line intensity [11] to benchmarksimulations of nucleosynthesis in models of massive-star evolution and death [12].Classical novae are thermonuclear explosions onhydrogen-accreting white-dwarf stars that have been es-timated to contribute up to 0.4 M ⊙ to the Galactic Alinventory [13]. This contribution needs to be quantifiedaccurately for the intrinsic study of classical novae andbecause it could present a significant background to themassive-star component. Fortunately, modeling of nu-cleosynthesis in novae is relatively advanced and now in-cludes experimental constraints on most of the essentialthermonuclear reaction rates [14], which are primarilyresonant radiative proton captures at peak temperaturesbetween 0.1 and 0.4 GK. For example, the direct produc-tion mechanism for Al, the Mg( p, γ ) Al reaction, iswell studied experimentally [15–18] because the reactantsare stable. A recent experiment using a Al rare isotopebeam has reduced the uncertainty in the rate of the di-rect destruction mechanism, the Al( p, γ ) Si reaction[19]. The dominant outstanding experimental nuclear-physics uncertainty lies in the thermonuclear rate of the Al( p, γ ) Si reaction ( Q = 5513 . ± . Al ground state [16]and, therefore, reduces the intensity of 1809-keV gammaray emission.The Al( p, γ ) Si rate depends on the center-of-massenergies and strengths of Si resonances (for a recentsummary see Ref. [23]). Direct measurements of the res-onance strengths using a Al ( τ = 10 . Si have included a variety of experimental nuclear-physics methods utilizing both stable and rare isotopebeams [20, 21, 24–35]. In addition, reaction-rate eval-uations employing available data and supplemented byshell-model calculations or information from the isospinmirror nucleus Mg have been conducted [16, 23, 36–38].Currently, there are three known Si states (spin andparity 1 + , 0 + , and 3 + ) that could potentially contributeto the Al( p, γ ) Si reaction rate as resonances at novatemperatures. The center-of-mass energy of the 1 + res-onance is 163 . ± . Q value.The 1 + resonance likely only contributes to the total re-action rate at temperatures below 0.2 GK, where the Al( p, γ ) Si route bypassing Al is not strongly acti-vated [38]. The energy of the 0 + resonance is not settled,but it appears to lie in the vicinity of 400 keV based on itspopulation in the Mg( He, n ) Si reaction [25, 33, 35].The strength of the 0 + resonance is expected [27] to bemuch lower than the nearby 3 + resonance and, therefore,the 3 + resonance likely dominates the reaction rate at thehighest nova temperatures, where Al( p, γ ) Si is mostactive, making experimental constraints on the energyand strength of the 3 + resonance essential. It has beenargued [23] that the 3 + resonance energy is 412 ± P beta-delayed proton decay [26] and otherexperimental data, corresponding to a Si excitation en-ergy of 5926 ± Q value; anexcitation-energy value of 5927 ± Si( p, t ) Si ∗ ( p ) Al reaction [32] is ingood agreement. The proton-decay partial width of the3 + resonance has been determined to be Γ p = 2 . ± . Al( d, n ) Si ∗ ( p ) Al reaction [30], provid-ing valuable further information on the Al( p, γ ) Si en-trance channel.The radiative exit channel of the key 3 + 25
Al( p, γ ) Siresonance sets the resonance strength, but it has notyet been observed due to the dominance of the proton-decay channel, which is generally expected to be about2 orders of magnitude stronger. Discovery of the exitchannel via the strongest expected primary gamma-raytransition ( E γ = 1739 ± ± + level at 4187 keV [28] could lead to anexperimental value for the small gamma-ray branchingratio Γ γ / Γ ≈ Γ γ / Γ p of the 3 + resonance. Together withΓ p [30], such a value would complete the experimental information on Si that is needed to calculate the res-onance strength without relying heavily on shell-modelcalculations or properties of the mirror nucleus. Evena sufficiently strong upper limit could prove that the Al( p, γ ) Si reaction channel bypassing Al produc-tion is effectively closed in novae.We have exploited the strong population of the third3 + (3 +3 ) Si level of interest in P beta decay [23, 26]to search for the radiative exit channel and measure thegamma-ray branching ratio. Fast ions of P were pro-duced at Michigan State University’s National Super-conducting Cyclotron Laboratory (NSCL) using projec-tile fragmentation of a 150 MeV/ u , 75 pnA Ar pri-mary beam from the Coupled Cyclotron Facility, inci-dent upon a 1.55 g/cm Be transmission target. The P ions were separated from other fragmentation prod-ucts by magnetic rigidity using the A1900 fragment sep-arator [39] (employing a 120 mg/cm Al wedge) and bytime of flight using a radio-frequency fragment separa-tor [40]. Up to 100 P ions s − were delivered to theexperimental setup. Clean ion identification was accom-plished using both the time of flight from a scintillatorat the focal plane of the A1900 to two 60 µ m-thick sili-con detectors located ≈
70 cm upstream of the countingstation, and the energy losses in the latter. On average,the composition of the beam delivered to the experimentwas found to be 74% P with 18% contamination by Al and small fractions of lighter ions. The P ionswere implanted in a 1-cm thick planar germanium detec-tor (GeDSSD) [41] that was divided electronically by 16segmented strips of 5 mm pitch on the front side and 16orthogonal ones on the back. The GeDSSD recorded theradioactive ion implantations and their subsequent betadecays using parallel low- and high-gain amplifications,respectively. The SeGA array of Ge detectors [42] sur-rounded the GeDSSD in two coaxial 13-cm radius ringsconsisting of eight germanium detectors apiece and wasused to detect gamma rays. The NSCL digital data ac-quisition [43] was employed.The SeGA spectra were gain matched to producecumulative spectra with 1-keV-wide bins using thestrong gamma-ray lines from room-background activityat 1460.8 keV (from K decay) and 2614.5 keV (from
Tl decay) as reference points, providing an in situ first-order relative-energy calibration. Efficiency calibra-tions were performed using standard sources of , Euand Co placed along the beam axis on the outside sur-face of the GeDSSD cryostat (5 cm downstream of the P implantation position). The calibration data wereused together with geant4
Monte Carlo simulations in-corporating the essential components of the experimentalgeometry to determine the efficiency at the P position300 µ m deep inside the GeDSSD crystal.In order to reduce the room-background contributionto the online gamma-ray spectra, a beta-delayed gamma-ray spectrum (see Fig. 1) was produced by requiring coin- FIG. 1. (color online). Cumulative P( βγ ) Si (left axis) and P( βγγ ) Si (right axis) spectra. The data points show the βγ spectrum with error bars spanning 1 standard deviation.The solid line is a fit to the data including known gamma-ray transitions (Doppler broadened for Al), a straight-linebackground, and a new peak at 1742 keV. The gamma-rayemitting nuclides contributing to each feature in the spec-trum are labeled, where two asterisks denote peaks producedby the escape of two 511-keV positron-annihilation gammarays. The open histogram shows βγγ coincidences with 1401-keV gamma rays. The hatched histogram shows coincidenceswith continuum background in a relatively broad energy re-gion just above 1401 keV in 9-keV-wide bins and normalizedto correspond to the expected background per keV in the1401-keV coincidence spectrum. cidences with high-gain events in the GeDSSD, which in-cluded beta particles, in a 1 . µ s software gate. Althoughthere were clear contributions from beam contaminants,room background, and daughter activities, this βγ spec-trum was dominated in the region of interest by Si and Al lines from the decay of P.The spectrum was fit (see Fig. 1) in the region of in-terest using an exponentially modified Gaussian effectiveresponse function whose shape was fixed based on peaksof similar energy in the high statistics gamma-ray singlesspectrum. A small extra peak was needed at 1801 keVto achieve a reasonable fit, but we considered this to bepart of the intense 1797-keV line, which could not be fitadequately with our simplified response model. Dopplerbroadening of peaks originating from the P( βpγ ) Aldecay sequence was incorporated to account for nuclearrecoil. Where relative energies or intensities of lines weresufficiently well known, they were constrained in the fit.The continuum background from Compton scattering ofhigher-energy gamma rays was modeled with a straightline spanning the range shown in Fig. 1; there was noevidence for significant curvature.We observed evidence for an unbroadened peak in theregion of interest that was 3.9 standard deviations abovethe expected background level. The energy of this peakwas found to be 1741 . ± . ± . +3 → +2 transitionto the 4187-keV state in Si. When additional narrowpeaks were added to the fit in the vicinities of 1735 and1754 keV (neither one included in the fit shown in Fig.1), they were found to be 2.6 standard deviations abovethe expected background level. The 1742-keV peak isthe most statistically significant new peak in the spec-trum and it is consistent with the expected energy forthe 3 +3 → +2 transition.In order to further test the hypothesis that the new1742-keV peak is from the 3 +3 → +2 26 Si gamma-raytransition following the beta decay of P, we searchedfor gamma-ray coincidences with the 1401-keV gammaray, which is known to be the strongest transition deex-citing the 3 +2 level at 4187 keV [28]. Supposing thatthe 1742-keV peak in the βγ spectrum was from the3 +3 level and using its measured intensity together withthe known branching [23, 28] and detection efficiencyfor the 1401-keV gamma ray, we would expect to ob-serve 3 . ± . ± . βγγ coincidence spectrum with nine candidateevents over a 10-keV range that could be reasonably at-tributed to the 1742-keV peak in the βγ spectrum. Weestimate the background to be 0.3 counts/keV using bothcoincidences in nearby bins and the spectrum observedin coincidence with a relatively broad background re-gion just above 1401 keV, suggesting an expected back-ground of three counts over the 10-keV range. Assuminga Poisson distribution, the probability of obtaining nineor more counts when three background counts are ex-pected is only 0.4 %. This 99.6 % confidence-level excessin the vicinity of the 1742-keV peak is the most statis-tically significant coincidence signal in the region of in-terest with the exception of the signal from the 1797-keV peak, for which the corresponding gamma ray isknown [26, 28] to be emitted in cascade with the 1401-keV gamma ray. Considering the expected coincidencebackground of three counts and the nine observed events,we find the observed number of real coincidences to be6 . +3 . − . , where the uncertainties are adopted from thetables of Ref. [44]. This βγγ coincidence result is consis-tent with the hypothesis that the 1742-keV peak in the βγ spectrum is produced by deexcitation of the 3 +3 26 Silevel, providing further evidence for such an identifica-tion.While no individual piece of evidence on the identi-fication of the 1742-keV gamma ray (energy, intensity,coincidences) is absolutely conclusive on its own, consid-eration of the evidence as a whole presents a relativelystrong case that this gamma ray is from the 3 +3 → +226 Si transition. The most significant excess in the regionof interest of both the βγ and βγγ spectra is at 1742keV, which is consistent with the expected transition en-ergy. The intensities of the signals at 1742 keV in thetwo spectra are mutually consistent. The βγ feature at1754 keV and, to a lesser extent the one at 1735 keV,are inconsistent with the transition energy. In addition,these less statistically significant βγ features do not havesignificant corresponding excesses in the βγγ spectrum.We will, therefore, consider the 1742-keV gamma ray tobe from the 3 +3 → +2 26 Si transition for subsequent cal-culations.Using the fits of the 1742- and 1797-keV peaks in the βγ spectrum shown in Fig. 1 [26, 28], the ratio of theirrespective areas was derived to be (3 . ± . × − .Using this value, the known intensity of the 1797-keVline (52% ±
11% [26]), and the ratio of efficiencies be-tween these two energies (effectively unity) we derivethe βγ intensity of the 1742-keV gamma ray to be[0 . ± . ± . +3 level is expectedto be 71 +13 − % based on the decay of the Mg mirror level[23] (the shell model also predicts 71% [45]), suggestingthat the total βγ intensity of all primary gamma raysfrom this level is [0 . ± . +0 . − . (lit.)]%. For com-parison, the beta delayed proton decay intensity throughthis level is 17 . ± .
90% [26]. Dividing the gamma in-tensity by the proton intensity yields the ratio of partialwidths, Γ γ / Γ p = 0 . ± . +0 . − . (lit.). Adopt-ing the experimentally determined value of Γ p = 2 . ± . γ = 40 ± +19 − (lit.) meV allow-ing us to calculate a Al( p, γ ) Si resonance strength ωγ = 23 ± +11 − (lit.) meV.The presently derived Al( p, γ ) Si resonancestrength is the first one based on measurements of Sipartial widths and branching ratios. The shell model hasbeen used elsewhere [37] to predict a resonance strengthof 68 meV, whereas estimates [23, 24] based on a lifetimemeasurement [46] of the Mg mirror level yield a valueof 18 +18 − meV. The resonance strength derived in thepresent work favors those based on the mirror level.We derive a Al( p, γ ) Si resonance energy for the3 +3 level using our measured primary gamma-ray energyby adding it to the excitation energy of the 3 +2 state(4187 . ± . . ± . ± . ± . Q value yields a resonance energy of414 . ± . ± . ± . ± ± p, t ) reaction TABLE I. Mass ejected from shiva [47] simulations of no-vae occurring on oxygen-neon white dwarfs of various masses. M tot is the total mass ejected in a single outburst; M ( , Al)are the masses of , Al ejected. The uncertainties showninclude only the effects of the standard deviation of the Al( p, γ ) Si reaction rate from the present work. The un-certainties in parentheses represent the results when the lowerlimit on the Al( p, γ ) Si reaction rate is calculated with the3 + resonance strength set equal to zero.White-dwarf mass 1.15 M ⊙ M ⊙ M ⊙ M tot (10 g) 4 . . . M ( Al)/ M tot (10 − ) 85 +1(+0) − +0(+0) − +1(+20) − M ( Al)/ M tot (10 − ) 9 . +0 . . − . . +0 . . − . . +0 . . − . [32].We calculated a new thermonuclear Al( p, γ ) Si re-action rate using the 3 + resonance energy and strength(and corresponding uncertainties) from the present work.For the 1 + and 0 + resonances and the direct-capturecomponent, we adopted the values and uncertainties fromRef. [23]. We simulated the production of Al in novaeon oxygen-neon (ONe) white dwarfs with masses of 1.15,1.25, and 1.35 M ⊙ using the spherically symmetric, La-grangian, hydrodynamic code shiva [47] and a nuclearreaction network that includes the rates from Ref. [14]and our new Al( p, γ ) Si rate. The simulations wererepeated with the standard-deviation limits on our rate.The results are summarized in Table I, which shows thatthe uncertainties related to the Al( p, γ ) Si reactionare typically ≪ Al( p, γ ) Si reaction-rate uncertainty was the last substantial experimentalnuclear-physics uncertainty associated directly with theexplosion, the reported Al yields represent model pre-dictions that are effectively independent of these experi-mental nuclear-physics uncertainties — a significant mile-stone.Following the estimate of Ref. [13] (based on Ref. [48])and changing only the amount of Al produced in novaeon 1.15 M ⊙ ONe white dwarfs (see Table I), we find anincrease from 20% to 30% in the maximum contributionof classical novae to the Al observed [5, 6] in the MilkyWay. In order to deduce the nova contribution more ac-curately, the number of ONe novae per year in the Galaxyneeds to be determined more accurately and multidimen-sional aspects of nova modeling, including mixing at thecore-envelope interface, need to be integrated with thenucleosynthesis.In conclusion, we have observed evidence for a new Pbeta-delayed gamma ray at 1742 keV. The gamma-rayenergy and intensity are consistent with those expectedfor the strongest primary transition deexciting the 3 +326
Si level. Coincidences with the secondary gamma rayat 1401 keV provide further evidence for such an identi-fication. This is the first experimental evidence for theexit channel of the key 3 + resonance in the thermonuclear Al( p, γ ) Si reaction rate, which influences the produc-tion of Al in classical-nova models. Using the energyand intensity of the observed Si gamma-ray line, wehave derived the resonance energy and strength, allow-ing us to estimate Al production in novae in a mannerthat is effectively independent of nuclear-physics uncer-tainties. We checked the sensitivity of Al productionto our lower limit by running our nova simulations withthe 3 + resonance strength set equal to zero and foundthat that only the simulation employing a 1 . M ⊙ whitedwarf displayed any change (see Table I). Therefore, onecould also interpret our experimental result as an upperlimit and reach essentially the same astrophysical con-clusions: our experiment is sufficiently sensitive to provefor the first time that the Al( p, γ ) Si reaction rate isvery slow so that the path bypassing Al productionis closed in novae hosted by typical oxygen-neon whitedwarfs with masses below 1 . M ⊙ and only open for ashort period of time near peak temperature for higherwhite-dwarf masses, which are expected to be scarce ac-cording to stellar evolution models.We encourage future measurements to further reducethe uncertainties in the 3 + 25 Al( p, γ ) Si rate, which aredominated by the resonance strength uncertainty. For ex-ample, a 5 σ detection of the 1742-keV gamma ray wouldbe an improvement and higher-statistics data on γγ co-incidences are desirable. First evidence for the weakerprimary gamma-decay branches from the 5929-keV Silevel could provide direct experimental constraints on the3 +3 → +2 branch we adopted from Ref. [23]. Indepen-dent measurements of Γ p could be conducted to confirmthe existing value [30] and improve the uncertainty. Di-rect measurements of the 3 + resonance with intense low-energy Al beams will hopefully be feasible at the nextgeneration of rare-isotope facilities; the present resultsprovide essential information for the planning of such ex-periments.This work was supported by the U.S. National Sci-ence Foundation under Grants No. PHY-1102511 andNo. PHY 08-22648, the U.S. Department of Energyunder Contract No. DE-FG02-97ER41020, the U.S.National Nuclear Security Agency under Contract No.DE-NA0000979, MEC Grant No. AYA2010-15685, andthe ESF EUROCORES Program EuroGENESIS throughGrant No. EUI2009-04167. We gratefully acknowledgeA. Garc´ıa for advice during the preparation of our ex-perimental proposal and the NSCL Operations staff fordelivering the beam. ∗ [email protected] † [email protected][1] W. A. Mahoney, J. C. Ling, W. A. Wheaton, and A. S. Jacobson, Astrophys. J. , 578 (1984).[2] G. H. Share, R. L. Kinzer, J. D. Kurfess, D. J. Forrest,E. L. Chupp, and E. Rieger, Astrophys. J. , L61(1985).[3] C. J. MacCallum, A. F. Huters, P. D. Stang, and M. Lev-enthal, Astrophys. J. , 877 (1987).[4] R. Diehl, C. Dupraz, K. Bennett, H. Bloemen,W. Hermsen, J. Knoedlseder, G. Lichti, D. Morris,J. Ryan, V. Schoenfelder, H. Steinle, A. Strong, B. Swa-nenburg, M. Varendorff, and C. Winkler, Astron. Astro-phys. , 445 (1995).[5] R. Diehl, H. Halloin, K. Kretschmer, G. G. Lichti,V. Sch¨onfelder, A. W. Strong, A. von Kienlin, W. Wang,P. Jean, J. Kn¨odlseder, J.-P. Roques, G. Weidenspoint-ner, S. Schanne, D. H. Hartmann, C. Winkler, andC. Wunderer, Nature (London) , 45 (2006).[6] W. Wang, M. G. Lang, R. Diehl, H. Halloin, P. Jean,J. Kn¨odlseder, K. Kretschmer, P. Martin, J. P. Roques,A. W. Strong, C. Winkler, and X. L. Zhang, Astron.Astrophys. , 713 (2009).[7] H. C. Urey, Proc. Natl. Acad. Sci. USA , 127 (1955).[8] G. Srinivasan, J. N. Goswami, and N. Bhandari, Science , 1348 (1999).[9] F. Timmes, “Connecting Nuclear Astrophysics with Ex-oplanets and Astrobiology,” 2012 Nuclear AstrophysicsTown Meeting, Detroit, MI, USA.[10] N. Prantzos and R. Diehl, Phys. Rep. , 1 (1996).[11] W. Wang, M. J. Harris, R. Diehl, H. Halloin, B. Cordier,A. W. Strong, K. Kretschmer, J. Kn¨odlseder, P. Jean,G. G. Lichti, J. P. Roques, S. Schanne, A. von Kienlin,G. Weidenspointner, and C. Wunderer, Astron. Astro-phys. , 1005 (2007).[12] M. Limongi and A. Chieffi, Astrophys. J. , 483 (2006).[13] J. Jos´e, M. Hernanz, and A. Coc, Astrophys. J. Lett. , L55 (1997).[14] C. Iliadis, R. Longland, A. E. Champagne, A. Coc, andR. Fitzgerald, Nucl. Phys. A841 , 31 (2010).[15] C. Iliadis, T. Schange, C. Rolfs, U. Schr¨oder, E. Somorjai,H. P. Trautvetter, K. Wolke, P. M. Endt, S. W. Kikstra,A. E. Champagne, M. Arnould, and G. Paulus, Nucl.Phys.
A512 , 509 (1990).[16] C. Iliadis, L. Buchmann, P. M. Endt, H. Herndl, andM. Wiescher, Phys. Rev. C , 475 (1996).[17] C. Iliadis, R. Longland, A. E. Champagne, and A. Coc,Nucl. Phys. A841 , 251 (2010).[18] F. Strieder, B. Limata, A. Formicola, G. Imbriani,M. Junker, D. Bemmerer, A. Best, C. Broggini, A. Ca-ciolli, P. Corvisiero, H. Costantini, A. DiLeva, Z. Elekes,Z. F¨ul¨op, G. Gervino, A. Guglielmetti, C. Gustavino,G. Gy¨urky, A. Lemut, M. Marta, C. Mazzocchi,R. Menegazzo, P. Prati, V. Roca, C. Rolfs, C. Rossi Al-varez, E. Somorjai, O. Straniero, F. Terrasi, and H. P.Trautvetter, Phys. Lett. B , 60 (2012).[19] C. Ruiz, A. Parikh, J. Jos´e, L. Buchmann, J. A.Caggiano, A. A. Chen, J. A. Clark, H. Crawford,B. Davids, J. M. D’Auria, C. Davis, C. Deibel, L. Erik-son, L. Fogarty, D. Frekers, U. Greife, A. Hussein, D. A.Hutcheon, M. Huyse, C. Jewett, A. M. Laird, R. Lewis,P. Mumby-Croft, A. Olin, D. F. Ottewell, C. V. Ouel-let, P. Parker, J. Pearson, G. Ruprecht, M. Trinczek,C. Vockenhuber, and C. Wrede, Phys. Rev. Lett. ,252501 (2006).[20] A. Parikh, J. A. Caggiano, C. Deibel, J. P. Greene,R. Lewis, P. D. Parker, and C. Wrede, Phys. Rev. C , 055804 (2005).[21] T. Eronen, V.-V. Elomaa, U. Hager, J. Hakala, A. Joki-nen, A. Kankainen, T. Kessler, I. D. Moore, S. Rahaman,J. Rissanen, C. Weber, and J. ¨Ayst¨o, Phys. Rev. C ,032802(R) (2009).[22] G. Audi, F. Kondev, M. Wang, B. Pfeiffer, X. Sun, J. Bla-chot, and M. MacCormick, Chin. Phys. C , 1157(2012).[23] C. Wrede, Phys. Rev. C , 035803 (2009).[24] D. W. Bardayan, J. C. Blackmon, A. E. Champagne,A. K. Dummer, T. Davinson, U. Greife, D. Hill, C. Iliadis,B. A. Johnson, R. L. Kozub, C. S. Lee, M. S. Smith, andP. J. Woods, Phys. Rev. C , 032801(R) (2002).[25] Y. Parpottas, S. M. Grimes, S. Al-Quraishi, C. R. Brune,T. N. Massey, J. E. Oldendick, A. Salas, and R. T.Wheeler, Phys. Rev. C , 065805 (2004).[26] J.-C. Thomas, L. Achouri, J. ¨Ayst¨o, R. B´eraud, B. Blank,G. Canchel, S. Czajkowski, P. Dendooven, A. Ensallem,J. Giovinazzo, N. Guillet, J. Honkanen, A. Jokinen,A. Laird, M. Lewitowicz, C. Longour, F. Oliveira San-tos, K. Per¨aj¨arvi, and M. Stanoiu, Eur. Phys. J. A ,419 (2004).[27] D. W. Bardayan, J. A. Howard, J. C. Blackmon, C. R.Brune, K. Y. Chae, W. R. Hix, M. S. Johnson, K. L.Jones, R. L. Kozub, J. F. Liang, E. J. Lingerfelt, R. J.Livesay, S. D. Pain, J. P. Scott, M. S. Smith, J. S.Thomas, and D. W. Visser, Phys. Rev. C , 045804(2006).[28] D. Seweryniak, P. J. Woods, M. P. Carpenter, T. Davin-son, R. V. F. Janssens, D. G. Jenkins, T. Lauritsen, C. J.Lister, J. Shergur, S. Sinha, and A. Woehr, Phys. Rev.C , 062801(R) (2007).[29] Y. K. Kwon, C. S. Lee, J. Y. Moon, J. H. Lee, J. Y.Kim, M. K. Cheoun, S. Kubono, N. Iwasa, K. Inafuku,H. Yamaguchi, J. J. He, A. Saito, Y. Wakabayashi, H. Fu-jikawa, G. Amadio, L. H. Khiem, M. Tanaka, A. A. Chen,S. Kato, Y. Fuchi, and N. Fukunishi, J. Korean Phys.Soc. , 1141 (2008).[30] P. N. Peplowski, L. T. Baby, I. Wiedenh¨over, S. E.Dekat, E. Diffenderfer, D. L. Gay, O. Grubor-Urosevic,P. H¨oflich, R. A. Kaye, N. Keeley, A. Rojas, andA. Volya, Phys. Rev. C , 032801(R) (2009).[31] A. Matic, A. M. van den Berg, M. N. Harakeh,H. J. W¨ortche, G. P. A. Berg, M. Couder, J. G¨orres,P. Leblanc, S. O’Brien, M. Wiescher, K. Fujita,K. Hatanaka, Y. Sakemi, Y. Shimizu, Y. Tameshige,A. Tamii, M. Yosoi, T. Adachi, Y. Fujita, Y. Shimbara,H. Fujita, T. Wakasa, B. A. Brown, and H. Schatz, Phys.Rev. C , 025807 (2010).[32] K. A. Chipps, D. W. Bardayan, K. Y. Chae, J. A.Cizewski, R. L. Kozub, J. F. Liang, C. Matei, B. H.Moazen, C. D. Nesaraja, P. D. O’Malley, S. D. Pain,W. A. Peters, S. T. Pittman, K. T. Schmitt, and M. S.Smith, Phys. Rev. C , 045803 (2010).[33] N. de Sereville, M. Assie, I. Bahrini, D. Beaumel, M. Chabot, A. Coc, I. Deloncle, F. de Oliveira, J. Duprat,M. Ferraton, S. Fortier, S. Franchoo, S. Giron, F. deGrancey, F. Hammache, C. Hammadache, J. Kiener,L. Lamia, M. Lebois, A. Lefebvre-Schuhl, F. Marechal,A. Matta, B. Mouginot, C. Petrache, G. Pizzonne, S. Ro-mano, P. Roussel, J. A. Scarpaci, I. Stefan, J. C. Thomas,D. Verney, M. Fallot, and L. Giot, Proc. Sci., NIC XI ,(2010) 212.[34] J. Chen, A. A. Chen, A. M. Amthor, D. Bazin, A. D.Becerril, A. Gade, D. Galaviz, T. Glasmacher, D. Kahl,G. Lorusso, M. Matos, C. V. Ouellet, J. Pereira,H. Schatz, K. Smith, B. Wales, D. Weisshaar, andR. G. T. Zegers, Phys. Rev. C , 045809 (2012).[35] T. Komatsubara, A. Ozawa, T. Moriguchi, Y. Ito,Y. Ishibashi, Y. Abe, T. Yuasa, T. Hayakawa,T. Shizuma, K. Y. Hara, S. Kubono, H. Yamaguchi,D. Kahl, S. Hayakawa, D. N. Binh, A. A. Chen, J. Chen,K. Setoodehnia, and T. Kajino, Proc. Sci., NIC XII ,(2012) 206.[36] C. Iliadis, J. M. D’Auria, S. Starrfield, W. J. Thomp-son, and M. Wiescher, Astrophys. J. Suppl. Ser. ,151 (2001).[37] W. A. Richter, B. A. Brown, A. Signoracci, and M. Wi-escher, Phys. Rev. C , 065803 (2011).[38] A. Parikh and J. Jos´e, Phys. Rev. C , 048801 (2013).[39] D. J. Morrissey, B. M. Sherrill, M. Steiner, A. Stolz, andI. Wiedenhoever, Nucl. Instrum. Methods Phys. Res.,Sect. B , 90 (2003).[40] D. Bazin, V. Andreev, A. Becerril, M. Dol´eans, P. F.Mantica, J. Ottarson, H. Schatz, J. B. Stoker, andJ. Vincent, Nucl. Instrum. Methods Phys. Res. Sect. A , 314 (2009).[41] N. Larson, S. Liddick, M. Bennett, A. Bowe, A. Chemey,C. Prokop, A. Simon, A. Spyrou, S. Suchyta, S. Quinn,S. Tabor, P. Tai, V. Tripathi, and J. VonMoss, Nucl.Instrum. Methods Phys. Res. Sect. A , 59 (2013).[42] W. F. Mueller, J. A. Church, T. Glasmacher,D. Gutknecht, G. Hackman, P. G. Hansen, Z. Hu, K. L.Miller, and P. Quirin, Nucl. Instrum. Methods Phys.Res., Sect. A , 492 (2001).[43] K. Starosta, C. Vaman, D. Miller, P. Voss, D. Bazin,T. Glasmacher, H. Crawford, P. Mantica, H. Tan,W. Hennig, M. Walby, A. Fallu-Labruyere, J. Harris,D. Breus, P. Grudberg, and W. K. Warburton, Nucl.Instrum. Methods Phys. Res., Sect. A , 700 (2009).[44] G. J. Feldman and R. D. Cousins, Phys. Rev. D , 3873(1998).[45] B. A. Brown, private communication.[46] F. Glatz, S. Norbert, E. Bitterwolf, A. Burkard, F. Hei-dinger, T. Kern, R. Lehmann, H. R¨opke, J. Siefert,C. Schneider, and B. H. Wildenthal, Zeit. Phys. A ,187 (1986).[47] J. Jos´e and M. Hernanz, Astrophys. J. , 680 (1998).[48] A. Weiss and J. W. Truran, Astron. Astrophys.238