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It is clear from figures 6 and 7 that a division can be made between clusters with and without black holes using only one of the two quantities that we have explored. the central logarithmic slope of the surface brightness profile.
It is clear from figures 6 and 7 that a division can be made between clusters with and without black holes using only one of the two quantities that we have explored, the central logarithmic slope of the surface brightness profile.
The r./r} ratio cannot distinguish between cases with and without black holes. only between clusters that have undergone and the rest.
The $r_c/r_h$ ratio cannot distinguish between cases with and without black holes, only between clusters that have undergone core-collapse and the rest.
Clusters that have achieved core-collapse separate cleanly from the rest in both indicators. which leads to the exclusion of very concentrated clusters. like MIS. as candidates for hosting an IMBH.
Clusters that have achieved core-collapse separate cleanly from the rest in both indicators, which leads to the exclusion of very concentrated clusters, like M15, as candidates for hosting an IMBH.
None of our models reproduce central slopes between 0.5 and 0.65 and we observe a few Galactic globular clusters with those slopes.
None of our models reproduce central slopes between 0.5 and 0.65 and we observe a few Galactic globular clusters with those slopes.
Since we are not able to follow the details of the evolution right before core-collapse due to the time intervals between snapshots. we cannot rule out the possibility that clusters in this stage could have intermediate slopes Even if clusters undergo such a phase. it is expected to be only for a very short time.
Since we are not able to follow the details of the evolution right before core-collapse due to the time intervals between snapshots, we cannot rule out the possibility that clusters in this stage could have intermediate slopes Even if clusters undergo such a phase, it is expected to be only for a very short time.
Two of the 14 models containing IMBHs do not show a clear central shallow cusp.
Two of the 14 models containing IMBHs do not show a clear central shallow cusp.
Even though it is impossible to draw statistical conclusions from such a small sample. we can say that the absence of shallow cusp does not imply the absence of a central black hole.
Even though it is impossible to draw statistical conclusions from such a small sample, we can say that the absence of shallow cusp does not imply the absence of a central black hole.
Therefore. some clusters with shallow cores might still be interesting candidates to follow up with kinematies.
Therefore, some clusters with shallow cores might still be interesting candidates to follow up with kinematics.
Finally. clusters with central slopes between -O.1 and -0.45 are clear candidates for harboring central black holes since we can only reproduce shallow central slopes by including intermediate-mass black holes.
Finally, clusters with central slopes between -0.1 and -0.45 are clear candidates for harboring central black holes since we can only reproduce shallow central slopes by including intermediate-mass black holes.
We conclude that the central logarithmic surface brightness slope appears to be a good diagnostic tool for choosing star clusters candidates for harboring intermediate-mass black holes.
We conclude that the central logarithmic surface brightness slope appears to be a good diagnostic tool for choosing star clusters candidates for harboring intermediate-mass black holes.
H.B. acknowledges support from the German Science foundation through a Heisenberg Fellowship and from the Australian Research Council through Future Fellowship grant FT0991052.
H.B. acknowledges support from the German Science foundation through a Heisenberg Fellowship and from the Australian Research Council through Future Fellowship grant FT0991052.
The authors want to thank the hospitality of the Kavli Institute for Theoretical Physics at UCSB. as well as the organizers of the “Formation and Evolution of Globular Clusters” program there.
The authors want to thank the hospitality of the Kavli Institute for Theoretical Physics at UCSB, as well as the organizers of the 'Formation and Evolution of Globular Clusters' program there.
We also thank the anonymous referee for making useful suggestions that helped improve the nanuseript.
We also thank the anonymous referee for making useful suggestions that helped improve the manuscript.
unstable modes must be regarded as viable suggestions.
unstable modes must be regarded as viable suggestions.
In the wake of uncertainties in the description of convection in red giants. we propose a closer look at observational data to see how the properties of low-amplitude oscillations in luminous red giants follow the trends seen in fainter objects.
In the wake of uncertainties in the description of convection in red giants, we propose a closer look at observational data to see how the properties of low-amplitude oscillations in luminous red giants follow the trends seen in fainter objects.
We began by comparing patterns in power spectra.
We began by comparing patterns in power spectra.
Figure I shows examples of spectra for OSARGs. based on the OGLE I-band photometry.
Figure 1 shows examples of spectra for OSARGs, based on the OGLE -band photometry.
We chose the ones that have prominent peaks contributing to à». às. b». bs sequences (Soszynisski et al.
We chose the ones that have prominent peaks contributing to $_2$ , $_3$ , $_2$, $_3$ sequences (Soszyńsski et al.
2007). which are the most prominent in the PL-plane.
2007), which are the most prominent in the PL-plane.
The patterns seen in this figure are reminiscent of those shown in Fig.
The patterns seen in this figure are reminiscent of those shown in Fig.
| of De Rider et al. (
1 of De Rider et al. (
2009) or Fig.
2009) or Fig.
2 of Stello et al. (
2 of Stello et al. (
2010). except that the frequency range is a factor 100 lower.
2010), except that the frequency range is a factor 100 lower.
In all perrodograms shown in our Fig.
In all periodograms shown in our Fig.
|. we see the two prominent peaks that were the feature behind our selection rule. but besides this. there is considerable diversity. especially in the amount of power around the dominant peaks and in the presence of the secondary peaks.
1, we see the two prominent peaks that were the feature behind our selection rule, but besides this, there is considerable diversity, especially in the amount of power around the dominant peaks and in the presence of the secondary peaks.
In all cases there is a power localized in some intermediate frequency range that moves leftwards with increasing luminosity.
In all cases there is a power localized in some intermediate frequency range that moves leftwards with increasing luminosity.
Simultaneously. the separation between dominant peaks decreases.
Simultaneously, the separation between dominant peaks decreases.
These are features seen in all power spectra of SLOs.
These are features seen in all power spectra of SLOs.
The low-frequency power in the OGLE data may be attributed to the granulation noise. like in. other cool stars.
The low-frequency power in the OGLE data may be attributed to the granulation noise, like in other cool stars.
However. the high-amplitude peaks in this range come from well-known but still unexplained phenomenon of theperiod (LSP) observed in a significant fraction - 30%)) of red giants (see. e.g.. Nicholls et al.
However, the high-amplitude peaks in this range come from well-known but still unexplained phenomenon of the (LSP) observed in a significant fraction $\sim$ ) of red giants (see, e.g., Nicholls et al.
2009).
2009).
The power spectra shown in this figure are very similar to those shown by Tabur et al (2010) for bright Galactic M giants.
The power spectra shown in this figure are very similar to those shown by Tabur et al (2010) for bright Galactic M giants.
Some of those objects are evidently located near TRGB.
Some of those objects are evidently located near TRGB.
We turn now to data on all OSARGs that contribute to the sequences b» and bs.
We turn now to data on all OSARGs that contribute to the sequences $_2$ and $_3$.
From the OGLE-III Catalog. we took the respective frequencies and the Wesenheit index. Wj. for individual objects.
From the OGLE-III Catalog, we took the respective frequencies and the Wesenheit index, $W_I$, for individual objects.
This index. which is the reddening-free stellar magnitude. is defined as where V and / denote mean magnitudes in the respective bands.
This index, which is the reddening-free stellar magnitude, is defined as where $V$ and $I$ denote mean magnitudes in the respective bands.
Points in Fig.
Points in Fig.
2 present data on the most significant periodicities (S/N> 5).
2 present data on the most significant periodicities $S/N\ge5$ ).
Up to two points are plotted per star.
Up to two points are plotted per star.
To compare these data with information on SLOs in less luminous stars. we need stellar parameters. such as luminosity. effective temperature. and mass.
To compare these data with information on SLOs in less luminous stars, we need stellar parameters, such as luminosity, effective temperature, and mass.
To this end. we used isochrones downloaded from the BaSTI Library (Pietrinferni et al.
To this end, we used isochrones downloaded from the BaSTI Library (Pietrinferni et al.
2006) for ages and metal abundance parameters appropriate to stars at TRGB in the LMC (Salaris Girardi 2005).
2006) for ages and metal abundance parameters appropriate to stars at TRGB in the LMC (Salaris Girardi 2005).
For the distance modulus to LMC. we adopted 18.5 mag. which is close to the mean value from recent determinations by various methods (Sehaefer 2008).
For the distance modulus to LMC, we adopted 18.5 mag, which is close to the mean value from recent determinations by various methods (Schaefer 2008).
The characteristics of the selected isochrones and model parameters at the point where W;=12 mag are listed in Table |.
The characteristics of the selected isochrones and model parameters at the point where $W_I=12$ mag are listed in Table 1.
In the range covered by the data. there is about factor 4 increase in luminosity between the highest and the lowest frequencies.
In the range covered by the data, there is about factor 4 increase in luminosity between the highest and the lowest frequencies.
The three lines in the upper panel of Fig.
The three lines in the upper panel of Fig.
2 were calculated along the isochrones with the Kjeldsen Bedding (1995) expression for the frequency corresponding to maximum of the acoustic. power.
2 were calculated along the isochrones with the Kjeldsen Bedding (1995) expression for the frequency corresponding to maximum of the acoustic power.
This expression 1s consistent. with. data on solar-like oscillation from main sequence up to lower RGB (Bedding Kjelsen 2003. Stello et al.
This expression is consistent with data on solar-like oscillation from main sequence up to lower RGB (Bedding Kjelsen 2003, Stello et al.
2007).
2007).
We may see that it also approximately applies to our data. which extend up to TRGB.
We may see that it also approximately applies to our data, which extend up to TRGB.
This is an argument in favor of the solar-like nature of oscillations in type b OSARGs.
This is an argument in favor of the solar-like nature of oscillations in type b OSARGs.
but the possibility of self-excitation cannot be ruled out.
but the possibility of self-excitation cannot be ruled out.
The argument forof thelatter interpretation is that peaks are narrow. suggesting long-lived modes.
The argument forof thelatter interpretation is that peaks are narrow, suggesting long-lived modes.
Perturbations on a binary star with μι<0.05AU from the tidal field of the MDITI casily lead to the mereer of its two components: while tidal breakup of binary stars with «@j,;>2AU usually leads to ejections of stars with velocities substantially smaller than the hivpervelocities interested in this paper.
Perturbations on a binary star with $a_{\rm b,i}<0.05\AU$ from the tidal field of the MBH easily lead to the merger of its two components; while tidal breakup of binary stars with $a_{\rm b,i}>2\AU$ usually leads to ejections of stars with velocities substantially smaller than the hypervelocities interested in this paper.
For demonstration onlv. we simply choose μι=0.1AU aud e;=0 (or 0.1. 0.3. and 0.6 alternatively) in the three-body experiments preseuted in Sections ?? aud 77: but in Section [.. we adopt a distribution of «yj based on the coustraiut obtained from curent observations aud 65,40.
For demonstration only, we simply choose $a_{\rm b,i}=0.1\AU$ and $e_{\rm b,i}=0$ (or 0.1, 0.3, and 0.6 alternatively) in the three-body experiments presented in Sections \ref{subsec:first} and \ref{subsec:Multi}; but in Section \ref{sec:Result}, we adopt a distribution of $a_{\rm b,i}$ based on the constraint obtained from current observations and $e_{\rm b,i}=0$.
We set the seminuajor axis of the outer binary as dour,~OOL 5pc (Luetal.2009:Gillessen2009).. as the stellar binaries are possibly originated from the CWS disk within half a parsec from the ceutral MDIT.
We set the semimajor axis of the outer binary as $a_{\rm out,i}\sim$ $-$ $\pc$ \citep{LuJ09,Gillessenetal09}, as the stellar binaries are possibly originated from the CWS disk within half a parsec from the central MBH.
The periceuter distance of the binary is set to be close to the tidal radius of the stellar binary.
The pericenter distance of the binary is set to be close to the tidal radius of the stellar binary.
In Sections 7? and ??.. we set the penetration parameter D typically to be ~20) 250 aud choose ¢ourj=0.2pc for demonstration onlv.
In Sections \ref{subsec:first} and \ref{subsec:Multi}, we set the penetration parameter $D$ typically to be $\sim $ $-$ 250 and choose $a_{\rm out,i}=0.2\pc$ for demonstration only.
The οσαοι of the orbit of the stellar binary. relative to that of the outer binary. οC0, is asmued to be uniformly distributed iu coso. if not z].specified.
The orientation of the orbit of the stellar binary, relative to that of the outer binary, $\phi\in[0,\pi]$, is assumed to be uniformly distributed in $\cos\phi$, if not specified.
The initial orbital phases of both the stellay and outer binaries are randomly chosen iu all the following calculations.
The initial orbital phases of both the stellar and outer binaries are randomly chosen in all the following calculations.
As the stellar binaries are initially set ono orbits weakly bound to the ceutral MDIT. binaries with large scnetration paruueters (em. D> 150) maw revolve around the ceutral MDII for many (c.g.. 100 or even 1000) orbits.
As the stellar binaries are initially set on orbits weakly bound to the central MBH, binaries with large penetration parameters (e.g., $D\ga150$ ) may revolve around the central MBH for many (e.g., $-$ 100 or even $1000$ ) orbits.
In our calculations below. the period of he stellar binary is usually much smaller thau that of the outer binary.
In our calculations below, the period of the stellar binary is usually much smaller than that of the outer binary.
Therefore. most of the calculation nue mv be spent on inteerating the stellar binary orbit when it is faraway from the \IBIT.
Therefore, most of the calculation time may be spent on integrating the stellar binary orbit when it is faraway from the MBH.
Towever. when he binary star is faraway from the AIBIT. the tidal orque from the MDII on the binary star is rather weak. so we can approximate the motion of the stellar ynary around the MDIT into two independent two-body xoblenis: one is for the stellar binary. aud the other is or the outer binary on au elliptical orbit. both of which can be done analytically.
However, when the binary star is faraway from the MBH, the tidal torque from the MBH on the binary star is rather weak, so we can approximate the motion of the stellar binary around the MBH into two independent two-body problems: one is for the stellar binary, and the other is for the outer binary on an elliptical orbit, both of which can be done analytically.
We adopt the above two-bodv approximation when the tidal force from the MDBIT ou the stellar binary is less than 10.© of the eravitational force between its two components (we have checked the two-body approximation by settiug a lower threshold. c.g.. LO" or 10.ὃν aud ford no significant difference in our results).
We adopt the above two-body approximation when the tidal force from the MBH on the stellar binary is less than $10^{-6}$ of the gravitational force between its two components (we have checked the two-body approximation by setting a lower threshold, e.g., $10^{-7}$ or $10^{-8}$, and found no significant difference in our results).
With this approxiuation. our calculation time is sped up substautially when dour is large (κ) 0.2pe) aud its accuracy can still be maintained.
With this approximation, our calculation time is sped up substantially when $a_{\rm out}$ is large $\ga 0.2\pc$ ) and its accuracy can still be maintained.
A star may be tidally disrupted if its close passage to the MBU is <ra=RACALfintO. where is and R. ave the mass and radius of the star. respectively.
A star may be tidally disrupted if its close passage to the MBH is $\la r^*_{\rm tid}\equiv R_*(M_{\bullet}/m_*)^{1/3}$, where $m_*$ and $R_*$ are the mass and radius of the star, respectively.
The radius of a star can be roughly given by Π.κRoti.fAL.) for i,<AL. and RoxRodan... otherwise. where AR. is the solar radius (Torresctal.2010.audreferences therein).
The radius of a star can be roughly given by $R_*\propto R_{\odot}(m_*/M_{\odot})$ for $m_*<M_{\odot}$ and $R_*\propto R_{\odot}(m_*/ M_{\odot})^{0.75}$ otherwise, where $R_{\odot}$ is the solar radius \citep[][and references therein]{Torres09}.
Due to the tidal diszuption. part of the disrupted renmmnants nav be swallowed by the central MDBIT aud part may be ejected out.
Due to the tidal disruption, part of the disrupted remnants may be swallowed by the central MBH and part may be ejected out.
In our calculations. we terminate the evolution of the corresponding svstemi once any component of the stellar inaryv approaches within a radius μι frou the MDIT.
In our calculations, we terminate the evolution of the corresponding system once any component of the stellar binary approaches within a radius $r^*_{\rm tid}$ from the MBH.
The tidal torque from the MDITI may change the sclnimajor axis and the ccceutricity of the stellar uuarv during its close passage to the AIBIT (sce turther discussions m Section ??)).
The tidal torque from the MBH may change the semimajor axis and the eccentricity of the stellar binary during its close passage to the MBH (see further discussions in Section \ref{subsubsec:bound}) ).
In the three-body experiment. we assume that the two componcuts of the ünarv star merge iuto a single star once the distance )etwoeen the two componcuts becomes smaller than the stun of their racii. aud the evolution of the corresponding system ds terminated once a merecr event occurs.
In the three-body experiment, we assume that the two components of the binary star merge into a single star once the distance between the two components becomes smaller than the sum of their radii, and the evolution of the corresponding system is terminated once a merger event occurs.
The consequences on the cwnuamical iuteractions of stellar binaries with the MDII can be differeut depending pe1i whether the stellar binaries are mitiallv ou bowed or ibound orbits.
The consequences on the dynamical interactions of stellar binaries with the MBH can be different depending on whether the stellar binaries are initially on bound or unbound orbits.
Ifthe binary star is mitially bound to the MBM. it may revolve around the ceutral MDBII for παν orbits before its disruption or merger: while the binary initially unbound to the AIBIT passes by the AIBIT ouly once even if if is not broken up during its close passage (see further discussions in Sectious 27. and ??)).
If the binary star is initially bound to the MBH, it may revolve around the central MBH for many orbits before its disruption or merger; while the binary initially unbound to the MBH passes by the MBH only once even if it is not broken up during its close passage (see further discussions in Sections \ref{subsec:first} and \ref{subsec:Multi}) ).
Multiple times of luteractions of bouud binary stars with a ceutral AIDII have been discussed by Antoninietal.(2010).. in which they focus on the post-Neowtonian effects ou stellar orbits in the gravitational field of the \IBU aud mergers of the two components of the binary
Multiple times of interactions of bound binary stars with a central MBH have been discussed by \citet{Antonini09}, in which they focus on the post-Newtonian effects on stellar orbits in the gravitational field of the MBH and mergers of the two components of the binary.
In Section ??.. we analyze the consequence of thefirst close encounter of the stellar binary with the central MDII. and the changes in the orbital elements of the stellar binary Gt the binary star survives).
In Section \ref{subsec:first}, we analyze the consequence of the close encounter of the stellar binary with the central MBH, and the changes in the orbital elements of the stellar binary (if the binary star survives).
Then we cousicder the cumulative effects of multiplecucouuters iun Section. 77...
Then we consider the cumulative effects of multipleencounters in Section \ref{subsec:Multi}.
We focus on discussing how the spatial and velocity distributions of IIVSs ire connected with their originations. ejecting nechamisius. aud dynamical environments in the CC.
We focus on discussing how the spatial and velocity distributions of HVSs are connected with their originations, ejecting mechanisms, and dynamical environments in the GC.
In Section ??.. 10! three-body experiments are formed for cach of the four set of initial conditions. Le. eng=0. OL. 0.3. and 0.6. respectively,
In Section \ref{subsec:first}, $10^4$ three-body experiments are performed for each of the four set of initial conditions, i.e., $e_{\rm b,i}=0$, $0.1$, $0.3$, and $0.6$, respectively.
Tf a stellar ynary is not tidally broken up during its first close owsaee to the MBIT. the three-body experiment euds up when the stellar binary reaches the apoapsis after its first close passage to the MBIT.
If a stellar binary is not tidally broken up during its first close passage to the MBH, the three-body experiment ends up when the stellar binary reaches the apoapsis after its first close passage to the MBH.
In Section ??.. 104 three-ουν experiuents are performed for each set of initial conditious aud those experiments that do not lead to ciguption/1aerger/separation within 500 revolutions are excluded.
In Section \ref{subsec:Multi}, $10^4$ three-body experiments are performed for each set of initial conditions and those experiments that do not lead to disruption/merger/separation within $500$ revolutions are excluded.
Considering that binary stars are injected toward the MBIL on weakly bound orbits with eccentricity close to Laud D< 250. the changes in their orbits. as consequences of the tidal effect frou the MDIT. after their first close passages. eau be characterized by the following five categories.
Considering that binary stars are injected toward the MBH on weakly bound orbits with eccentricity close to 1 and $D\la 250$ , the changes in their orbits, as consequences of the tidal effect from the MBH, after their first close passages, can be characterized by the following five categories.
and not towards the E-W orientation. in agreement with the findings ofSivaramanetal. from an analysis of IXodaikanal white-light images.
and not towards the E-W orientation, in agreement with the findings of\citet{stenflo-sivaramanetal07} from an analysis of Kodaikanal white-light images.
Nosovichey&Stenflo(2008) further found no dependence of the tilt behavior on (he amount of [Iux or size of the bipolar regions.
\citet{stenflo-ks08} further found no dependence of the tilt behavior on the amount of flux or size of the bipolar regions.
Both these findings contradict the paradigm that the tilt is caused by the Coriolis force acting on initially untilted flux loops that. rise from a toroidal source region near the bottom of the convection zone and emerge al the surface as tilted bipolar regions (DFisheretal. 1995).
Both these findings contradict the paradigm that the tilt is caused by the Coriolis force acting on initially untilted flux loops that rise from a toroidal source region near the bottom of the convection zone and emerge at the surface as tilted bipolar regions \citep{stenflo-dsilva93,stenflo-fisheretal95}.
. Instead the tilt. which is the source of the N-S dipole moment that leads to the reversal ancl regeneration of the poloidal field. appears to have been established already in the dvnamo region in the Suns interior.
Instead the tilt, which is the source of the N-S dipole moment that leads to the reversal and regeneration of the poloidal field, appears to have been established already in the dynamo region in the Sun's interior.
The tilt observed al (he surface reflects (his property regardless of the size or amount of flix of the observed regions.
The tilt observed at the surface reflects this property regardless of the size or amount of flux of the observed regions.
We have made use of the complete set of 96 minute cadence SOIIO/MDI full disk magnetograms (Scherrerοἱal.1995).. which covers the 15 vear period Alay 1996 Αρνί 2011.
We have made use of the complete set of 96 minute cadence SOHO/MDI full disk magnetograms \citep{stenflo-scherreretal95}, which covers the 15 year period May 1996 – April 2011.
With a pixel size of 2x 2aarcsec?. the effective spatial resolution of the magnetograms is 4x daaresec?.
With a pixel size of $2\times 2$ $^2$, the effective spatial resolution of the magnetograms is $4\times 4$ $^2$.
The magnetograms represent maps of the line-o-sight component of the magnetic [lux density averaged over the spatial resolution window.
The magnetograms represent maps of the line-of-sight component of the magnetic flux density averaged over the spatial resolution window.
Thev have been derived from maps of the circular polarization recorded with a narrow-band filler at different wavelength positions within the Ni lline (Scherreretal. 1995)..
They have been derived from maps of the circular polarization recorded with a narrow-band filter at different wavelength positions within the Ni line \citep{stenflo-scherreretal95}. .
Among the light elements Li. De and D (LiBeB). ‘Li is thought to arise from a variety of processes. including bie bane nucleosvnthesis (Spite&Spite1982).. asvimplolic eiant branch stars. novae (D'Antona&Matteucci 1991).. and the z-process in type IH supernovae (Woosleyοἱal.1990): the latter may also contribute to ! B. On the other hand. the main production channel for the rest. in particular for "Li and Be. is believed to be cosmic-ray induced nuclear reactions.
Among the light elements Li, Be and B (LiBeB), $^7$ Li is thought to arise from a variety of processes, including big bang nucleosynthesis \citep{Spite82}, asymptotic giant branch stars, novae \citep{DAntona91}, , and the $\nu$ -processin type II supernovae \citep{Woosley90}; the latter may also contribute to $^{11}$ B. On the other hand, the main production channel for the rest, in particular for $^6$ Li and Be, is believed to be cosmic-ray induced nuclear reactions.
The most wicely discussed models of LiBeB production are based on cosmic rays accelerated in supernova shocks (Meneguzzietal.1971:al. 2000).
The most widely discussed models of LiBeB production are based on cosmic rays accelerated in supernova shocks \citep{Meneguzzi71,Vangioni00}.
. Observations of metal-poor stars in the Galactic halo show a primary relation between [Fe/H] and. |De/II] or |D/IH]. which is consistent with spallation by cosmic rays enriched with C. N. and ο (CNOJ from Iresh SN ejecta impinging on interstellar 11 or Ie (Duncan.Lambert.&Lemke1992).
Observations of metal-poor stars in the Galactic halo show a primary relation between [Fe/H] and [Be/H] or [B/H], which is consistent with spallation by cosmic rays enriched with C, N, and O (CNO) from fresh SN ejecta impinging on interstellar H or He \citep{Duncan92}.
. An alternative possibility was proposed by Fieldsetal.(1996.2002).. who considered explosions of Type Ic supernovae (SNe Ic) as a site for primary LiBeB production.
An alternative possibility was proposed by \citet{Fields96, Fields02}, who considered explosions of Type Ic supernovae (SNe Ic) as a site for primary LiBeB production.
Since il is expected (hat a fraction of C and O in the surface lavers of the ejecta are accelerated to energies above the threshold for spallation reactions when (he shock passes through the stellar surface. LiBeB production can occur through the direct interaction of the ejecta with ihe ambient material. without the need for shock acceleration of cosmic ravs.
Since it is expected that a fraction of C and O in the surface layers of the ejecta are accelerated to energies above the threshold for spallation reactions when the shock passes through the stellar surface, LiBeB production can occur through the direct interaction of the ejecta with the ambient material, without the need for shock acceleration of cosmic rays.
This was explored in greater depth by Nalsunura&Shigevama(2004).. who used more realistic stellar models ancl equations of state. together will a I-dimensional relativistic hyvdrodyvnanmic code {ο investigate in detail how much of the ejecta mass acquires sullicient energies for the LiBeB production.
This was explored in greater depth by \citet{Nakamura04}, who used more realistic stellar models and equations of state, together with a 1-dimensional relativistic hydrodynamic code to investigate in detail how much of the ejecta mass acquires sufficient energies for the LiBeB production.
All of these studies used stellar models which completely lost their I and Ile envelopes and assumed the target to be interstellar matter (ISM) consisting of II and Ie. ignoring anv circumstellar matter (CSAIL).
All of these studies used stellar models which completely lost their H and He envelopes and assumed the target to be interstellar matter (ISM) consisting of H and He, ignoring any circumstellar matter (CSM).
Therefore. the only reactions under consideration were C.O+IL.We—LiBeb.
Therefore, the only reactions under consideration were ${\rm C, O} + {\rm H, He} \rightarrow {\rm LiBeB}$.