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At each mass shell matter [lows freely [rom above or below with a certain degree of mixing. and is also exchauged between adjacent cells. from one stream to the other. | At each mass shell matter flows freely from above or below with a certain degree of mixing, and is also exchanged between adjacent cells, from one stream to the other. |
Our nucleosvnuthesis network is based on 74 nuclear species. 59 nuclei [rom neutrons and protons up (ο sulphur and with another 14 nuclei near the iron group to allow neutron capture on iron seeds. | Our nucleosynthesis network is based on 74 nuclear species, 59 nuclei from neutrons and protons up to sulphur and with another 14 nuclei near the iron group to allow neutron capture on iron seeds. |
There is also an additional "particle" g for counting the number of neutron captures occurring bevond "Ni; which simulates the s process as neutron sink. | There is also an additional “particle” $g$ for counting the number of neutron captures occurring beyond $^{61}$ Ni, which simulates the $s$ process as neutron $sink$. |
The initial abundances in the postprocessing calculations are taken [rom (1989). | The initial abundances in the postprocessing calculations are taken from \citet{anders:89}. |
. All proton. a. neutron captures and 2» decays involving the species listed above are included in the nuclear network summing up to 506 reactions. | All proton, $\alpha$, neutron captures and $\beta$ decays involving the species listed above are included in the nuclear network summing up to 506 reactions. |
The bulk of reaction rates are [rom the REACLID Data Tables of nuclear reaction rates based on the 1991 updated version ol the compilation bv Thielemann.Arnould.&Truran (1986).. | The bulk of reaction rates are from the REACLIB Data Tables of nuclear reaction rates based on the 1991 updated version of the compilation by \citet*{thielemann:86}. . |
The reaction rate table has been updated using the latest experimental results. which are listed in Appendix A. The | The reaction rate table has been updated using the latest experimental results, which are listed in Appendix A. The |
For the moment, we consider a restricted closed box model in which we take the MGDR to be constant as a function of redshift. | For the moment, we consider a restricted closed box model in which we take the MGDR to be constant as a function of redshift. |
To assess the ability of this model to fit observations, we first combine eqs. ( | To assess the ability of this model to fit observations, we first combine eqs. ( |
1) and (2) to obtain RD)/dt=MGDRxSFRD. | 1) and (2) to obtain $d(SFRD)/dt = MGDR\times
SFRD$. |
We then note that our d(SFpiecewise linear fit to the observed SFRD as a function of time (Appendix A) implies d(SFRD)/dt~(.24Gyr!)xSF where the coefficient is about half of the observed MGDRRD, etal.2008). | We then note that our piecewise linear fit to the observed SFRD as a function of time (Appendix A) implies $d(SFRD)/dt \sim (.24\, {\rm
Gyr}^{-1}) \times SFRD$ , where the coefficient is about half of the observed MGDR \citep{Leroy2008}. |
. In other words, from (Leroythe assumption of a constant MGDR at the present epoch in a closed box Πο model, we find that the star formation rate is declining only half as fast as expected given our current reservoir of molecular gas. | In other words, from the assumption of a constant MGDR at the present epoch in a closed box $_2$ model, we find that the star formation rate is declining only half as fast as expected given our current reservoir of molecular gas. |
It is this discrepancy in the derivatives of the observed cosmic star formation rate and the rate at which we observe molecular gas being converted into stars that we call the cosmic molecular gas depletion problem. | It is this discrepancy in the derivatives of the observed cosmic star formation rate and the rate at which we observe molecular gas being converted into stars that we call the cosmic molecular gas depletion problem. |
We note here that given the uncertainties in the observations, this factor of two in itself is not a strong argument against the closed box model, but we will show in general that observational constraints rule outany closed box model. | We note here that given the uncertainties in the observations, this factor of two in itself is not a strong argument against the closed box model, but we will show in general that observational constraints rule out closed box model. |
We now relax the assumption of a constant MGDR in the closed box model in refRCBM.. | We now relax the assumption of a constant MGDR in the closed box model in \\ref{RCBM}. |
To study the predictions of this model, we calculate px,(z) by integrating Eq. (2)) | To study the predictions of this model, we calculate $\mrh2(z)$ by integrating Eq. \ref{CBder}) ) |
and using the observed as an input. | and using the observed $(z)$ as an input. |
We then divide SFRD(z) by px,(z) to SFRD(z)obtain MGDR(z). | We then divide $(z)$ by $\mrh2(z)$ to obtain $(z)$. |
The results are shown in Fig. 3,, | The results are shown in Fig. \ref{CBSFEevol}, |
where pzz,(0) is set to the mean value from Obreschkow&Rawling | where $\mrh2(0)$ is set to the mean value from \cite{OR2009}. |
s The uncertainties in the SFRD to the IMF) (2009)..discussed in refSFRDobs(due are seen to have only a minor effect on the resulting pg,(z) and MGDR(z). | The uncertainties in the SFRD (due to the IMF) discussed in \\ref{SFRDobs} are seen to have only a minor effect on the resulting $\mrh2(z)$ and $(z)$. |
Fig. | Fig. |
3 shows that increases by a factor of ~10 from z—0 to 1, and przr,(t)MGDR decreases with increasing redshift,$2. | \ref{CBSFEevol} shows that $\mrh2(t)$ increases by a factor of $\sim 10$ from $z=0$ to 1, and MGDR decreases with increasing redshift,. |
2. Thus, even the general closed box model is at odds with the observations, leading us to our next model. | Thus, even the general closed box model is at odds with the observations, leading us to our next model. |
Since a closed box model of only H32 and stars is inconsistent with observations, we now allow additional components that can be converted into H2 and then into stars. | Since a closed box model of only 2 and stars is inconsistent with observations, we now allow additional components that can be converted into $_2$ and then into stars. |
We consider separately the ggas and an external source of gas that we call peat, and modify Eq. (2)) | We consider separately the gas and an external source of gas that we call $\mrext$, and modify Eq. \ref{CBder}) ) |
to For the ggas, the observations discussed in refHIobs and Fig. | to For the gas, the observations discussed in \\ref{HIobs} and Fig. |
2 suggest that px7(z) is very slowly varying over cosmic timescales; therefore dog;/dt is small. | \ref{HIfig}
suggest that $\mrhi(z)$ is very slowly varying over cosmic timescales; therefore $d\mrhi/dt$ is small. |
Fig. | Fig. |
4 shows that the derivative of the observed (red curve) is an order of magnitude smaller than the observed SFRD (black curves). | \ref{HIder} shows that the derivative of the observed (red curve) is an order of magnitude smaller than the observed SFRD (black curves). |
In the absence of pezt, we have is approximately equal to SFRD, as in |dpy,/dt|the (bluefailed curves)closed box model. | In the absence of $\mrext$, we have $|d\mrh2/dt|$ (blue curves) is approximately equal to $SFRD$, as in the failed closed box model. |
Thus the inclusion of in an open-box model is not enough to fit the model to the observations. | Thus the inclusion of in an open-box model is not enough to fit the model to the observations. |
This leads to a robust conclusion: It is important at this point to clarify that represents the reservoir of bboth in galaxies as well as the ooutside of galaxies. | This leads to a robust conclusion: It is important at this point to clarify that represents the reservoir of both in galaxies as well as the outside of galaxies. |
This is because the DLA observations of Prochaska and Wolfe (2009) include all of the high column density neutral ((N(HI) >2x 107° cm-?2). | This is because the DLA observations of Prochaska and Wolfe (2009) include all of the high column density neutral (N(HI) $> 2 \times$ $^{20}$ 2). |
This gas contains at least of the neutral aatoms for z«6 (O'Meara et al 2007). | This gas contains at least of the neutral atoms for $z < 6$ (O'Meara et al 2007). |
We are therefore forced to include a nonzero dpext/dt term in the open box model. | We are therefore forced to include a nonzero $d\mrext/dt$ term in the open box model. |
This component represents | This component represents |
Although the tangents of all the spiral arms have been recognised in CO. aud other tracers such as CS (e.g.2). aud IR star counts (e.g.2).. identification of the arm origius. which fall within the terminal velocity euvelope rather than at its boundary. where the taugeut poiuts lie. has been elusive. | Although the tangents of all the spiral arms have been recognised in CO, and other tracers such as CS \citep[e.g.][]{bronfman92} and IR star counts \citep[e.g.][]{benjamin05}, identification of the arm origins, which fall within the terminal velocity envelope rather than at its boundary where the tangent points lie, has been elusive. |
This has meaut mocels generally incorporate arbitrary or iulerred arm origius within theLe domain: the two "major arms (Perseus aud Crux-Sceutum) origiuate at opposite ends of the Galactic bar and the two ‘minor arms (Norma and Carina-Sagittarius) are anywhere between a [ew degrees aud 907 offset. either starting from the radius of corotatiou resonance directly or branching [rom the two main arms. | This has meant models generally incorporate arbitrary or inferred arm origins within the domain: the two `major' arms (Perseus and Crux-Scutum) originate at opposite ends of the Galactic bar and the two `minor' arms (Norma and Carina-Sagittarius) are anywhere between a few degrees and $^{\circ}$ offset, either starting from the radius of corotation resonance directly or branching from the two main arms. |
However. studies of other Galaxies aud dynamical models show that the major arms do not necessarily start directly at the bar eid. but can be seen offset by several degrees of Galactic azimuth (e.g.?2).. | However, studies of other Galaxies and dynamical models show that the major arms do not necessarily start directly at the bar end, but can be seen offset by several degrees of Galactic azimuth \citep[e.g.][]{sandage94,russeil03}. |
As a product of interstellar material undergoing the influence of density waves. spiral arms are regious of deuse material and star formation. | As a product of interstellar material undergoing the influence of density waves, spiral arms are regions of dense material and star formation. |
External galaxies show the spiral armis start close o the end of the bar or a ring surrouudiug a bar and. as already. ineutioned. dynamical models jxace the arms starting at the radius of corotation resonance (7 | kpc). | External galaxies show the spiral arms start close to the end of the bar or a ring surrounding a bar and, as already mentioned, dynamical models place the arms starting at the radius of corotation resonance $\sim$ 4 kpc). |
These regious are likely to experience siguificaut compression of material such that star formation. Induced by deusity chauges in tlie ISM. is likely to be enhanced. | These regions are likely to experience significant compression of material such that star formation, induced by density changes in the ISM, is likely to be enhanced. |
The remaining high deusity regious of masers within E28" may hus be explained by the origius of the spiral arius. | The remaining high density regions of masers within $\pm$ $^{\circ}$ may thus be explained by the origins of the spiral arms. |
The exception is between loungitudes +6° to —6* where all the spiral arius pass and there are sources ou the solar circle (zero line-of-sight velocities) ogether[n] with sources associated with the Cralactic Centre Zone. | The exception is between longitudes $^{\circ}$ to $-$ $^{\circ}$ where all the spiral arms pass and there are sources on the solar circle (zero line-of-sight velocities) together with sources associated with the Galactic Centre Zone. |
Iu this region au overabuudance is to be expected even in tlie absence of spiral arm origius. | In this region an overabundance is to be expected even in the absence of spiral arm origins. |
High positive velocities at longitudes 25° to 35° have been cousidered the approximate start of the Crux-Sceutuim ari. with a region of massive star formation iuferred from a sequence of IR observations (222???) and a massive voung cluster of red supergiauts (?22?).. | High positive velocities at longitudes $^{\circ}$ to $^{\circ}$ have been considered the approximate start of the Crux-Scutum arm, with a region of massive star formation inferred from a sequence of IR observations \citep{hammersley94,garzon97,lopez99,hammersley00,cabrera08} and a massive young cluster of red supergiants \citep{davies07,clark09,negueruela10}. |
We see a lieh density of methanol masers in this region (1 bin in excess Of 265,4, ln rellvcdensity)). | We see a high density of methanol masers in this region (1 bin in excess of $\sigma_{poisson}$ in \\ref{lvdensity}) ). |
The Crux-Scutum arm is the ouly arm with velocities above loa these lougitudes. it is known to have a tangent at 7327. aud originate atlongitudes smaller than this. | The Crux-Scutum arm is the only arm with velocities above $^{-1}$ at these longitudes, it is known to have a tangent at $\sim$ $^{\circ}$, and originate atlongitudes smaller than this. |
The lower estimate of the radius of corotatiou resonance is 3.5 kpe. which is taugential at 21°. | The lower estimate of the radius of corotation resonance is 3.5 kpc, which is tangential at $^{\circ}$. |
Hence the arm theoretically originates between 21° aud 35° aud we interpret tlie deuse cluster of luasers al ο as observational evideuce of the start of this arm. | Hence the arm theoretically originates between $^{\circ}$ and $^{\circ}$ and we interpret the dense cluster of masers at $\sim$ $^{\circ}$ as observational evidence of the start of this arm. |
Assuming the longitudes 21° anc 3257 mark the boundary of this region. aud that it is a portion of an aunulus with a thickuess of ] kpe centred at the corotation resonance radius of { kpc. the origin bas a density of 11 sources per kpc. | Assuming the longitudes $^{\circ}$ and $^{\circ}$ mark the boundary of this region, and that it is a portion of an annulus with a thickness of 1 kpc centred at the corotation resonance radius of 4 kpc, the origin has a density of 11 sources per $^{2}$. |
The discovery of the massive young red supergiant cluster in this region (?)— inuplie substantial starburst activity 10-20 Myr ago (?).. but the siguificaut preseuce of 6.7-GHz methano tasers Ludicates there is still substantial ongoing star formation. | The discovery of the massive young red supergiant cluster in this region \citep{figer06} implied substantial starburst activity 10–20 Myr ago \citep{davies07}, , but the significant presence of 6.7-GHz methanol masers indicates there is still substantial ongoing star formation. |
A-band (1500 A-band | $K$ $>1500$ $K$ |
from spectral evolution of radio emission (?).. assuming Scomp/Cjet‘Y7) and the oDoppler rfactorfactor 06=tconst. | from spectral evolution of radio emission \citep{1999ApJ...521..509L}, assuming $T_\mathrm{b} \propto S_\mathrm{comp}/(r_\mathrm{jet}^2\nu^2)$ and the Doppler factor $\delta = const$. |
For Compton (&J. synchrotron (&) and adiabatic (ει) losses. € are calculated as follows. with a typical value of 5=2.0 (corresponding to a synchrotron spectral index w = -0.5: S,« v)janda=| (dominant transverse magnetic field). ε.=10/81.25. €,=25/6x4.17 and e,=19/63.17. | For Compton $\epsilon_\mathrm{c}$ ), synchrotron $\epsilon_\mathrm{s}$ ) and adiabatic $\epsilon_\mathrm{a}$ ) losses, $\epsilon$ are calculated as follows, with a typical value of $s = 2.0$ (corresponding to a synchrotron spectral index $\alpha$ = -0.5; $ S_\nu\propto\nu^{+\alpha}$ ) and $a = 1$ (dominant transverse magnetic field), $\epsilon_\mathrm{c} = 10/8 = 1.25$, $\epsilon_\mathrm{s} = 25/6 \approx 4.17$ and $\epsilon_\mathrm{a} = \boldsymbol{19}/6 \approx 3.17$. |
For the slopes to be shallower than these derived values. acceleration resulting in increasing Doppler factors of 0xm may be considered. | For the slopes to be shallower than these derived values, acceleration resulting in increasing Doppler factors of $\delta \propto r_\mathrm{jet}^b$ may be considered. |
For à moderate acceleration (i.e. b=0.1— 0.2). a value ε.=0.13 is obtained assuming »=2 and with a longitudinal magnetic field (a= 2). | For a moderate acceleration (i.e. $b=0.1-0.2$ ), a value $\epsilon_\mathrm{c} \approx 0.13$ is obtained assuming $s=2$ and with a longitudinal magnetic field $a=2$ ). |
The brightness temperature gradient along the jet. shown in Fig. 4.. | The brightness temperature gradient along the jet, shown in Fig. \ref{fig:Tb}, |
reveals a possible broken power-law behavior of the jet intensity gradient with a break distance of «0.3 mmas and the two slopes ει=0.95+0.69 and e»=4.1]+ 0.85. | reveals a possible broken power-law behavior of the jet intensity gradient with a break distance of $\sim$ mas and the two slopes $\epsilon_1 = 0.95\pm0.69$ and $\epsilon_2 = 4.11 \pm 0.85$ . |
The value of ει 1s consistent with e. with the possible indication of a mild change in 6. and e: is consistent with the derived &. | The value of $\epsilon_1$ is consistent with $\epsilon_\mathrm{c}$ , with the possible indication of a mild change in $\delta$, and $\epsilon_2$ is consistent with the derived $\epsilon_\mathrm{s}$. |
However. at the ~Ie level. e is also consistent with ει. | However, at the $\sim 1\sigma$ level, $\epsilon_2$ is also consistent with $\epsilon_\mathrm{a}$. |
Using the spectral evolution of a |jet component.| ? found evidence for a change from the synchrotron to the adiabatic stage at a distance of mmas from the core. | Using the spectral evolution of a jet component, \citet{1999ApJ...521..509L} found evidence for a change from the synchrotron to the adiabatic stage at a distance of mas from the core. |
This suggests that the transition from Compton to synchrotron stages is indeed observed. | This suggests that the transition from Compton to synchrotron stages is indeed observed. |
Future investigations on the spectral evolution of the jet are able to provide additional constraints on the transitionregions between theCompton. synchrotron. and adiabatie stages. | Future investigations on the spectral evolution of the jet are able to provide additional constraints on the transitionregions between theCompton, synchrotron, and adiabatic stages. |
For 3 EROs (71. 507. 551). the umber of net Nrav counts is good enough to allow a basic Xrav spectral analysis. | For 3 EROs (71, 507, 551), the number of net X–ray counts is good enough to allow a basic X–ray spectral analysis. |
Tn order to maximize the signaltonoise ratio. the Xrav spectra lave been extracted using circular regions of δή of radius centered on the Xrav position. | In order to maximize the signal–to–noise ratio, the X--ray spectra have been extracted using circular regions of $\sim$ of radius centered on the X–ray position. |
The vackerouud spectra have been extracted from source.free circular regions close to the object with an area about l times larecr. | The background spectra have been extracted from source–free circular regions close to the object with an area about 4 times larger. |
While EROs S2F1.771 aud S2F1.5507 ΕΙave 1been detecteletected‘ iu a:all the 3: e3EEPIC ninststruments,--te. EROE‘ S2F1.5551ens ihas‘ notí "been detected4η Liin the MOSLI cameraee aud only barely detected in. the MOS2s camera: therefore ]or this. latter object. onlv pu data have been used for. je spectralpec1 IEEEaualvsis. | While EROs 71 and 507 have been detected in all the 3 EPIC instruments, ERO 551 has not been detected in the MOS1 camera and only barely detected in the MOS2 camera; therefore for this latter object only pn data have been used for the spectral analysis. |
Ius EFFor hethe fastfirst ttwo ERO«EROs the MOSIΑΠΟ: aud MOS2 data have been combined. to maximize∙∙ the statistics. | For the first two EROs the MOS1 and MOS2 data have been combined to maximize the statistics. |
Finally. MOSes and pu spectra have been binned. in. order to have at least .15 counts per energy oneschannel “aud are fitted simultancouslyMSN in: the 0.5,;10 keV. baud leaving '16 relativeti uormalizatiousi5» Bocfree tto outvary. | Finally, MOS and pn spectra have been binned in order to have at least 15 counts per energy channel and are fitted simultaneously in the 0.5–10 keV band leaving the relative normalizations free to vary. |
4 sndie errors in the16 &following1. ofare reported|ted ‘at tlthe cont level for one interesting paranieter (ΔΑ 220392.1.> Avniη 1976 1916)). | All the errors in the following are reported at the confidence level for one interesting parameter $\Delta\chi^2$ =2.71, Avni \cite{Avni}) ). |
In the fitting procedime. the appropriate Calactic hydrogen colin density along the line of sight \ 102" 2τν Dickey Lockman 1990)) has Όσοι taken iuto account. | In the fitting procedure, the appropriate Galactic hydrogen column density along the line of sight $\times$ $^{20}$ $^{-2}$, Dickey Lockman \cite{Dickey}) ) has been taken into account. |
We find that the three Xrav spectra are well described. by a sinele. absorbed (Nyy2Dj)1077 5 7) law inodel at the zj55; of cach source. | We find that the three X–ray spectra are well described by a single absorbed $_H>10^{22}$ $^{-2}$ ) power--law model at the $_{phot}$ of each source. |
Due to the Xταν statistics of ERO S2FL5551 (ο net counts in the pu image). the intrinsic photon iudex for this ERO is fixed to 1.9. the typical value for unabsorbed AGN (Nandra et al. LOOT: | Due to the X–ray statistics of ERO 551 $\sim$ 60 net counts in the pn image), the intrinsic photon index for this ERO is fixed to 1.9, the typical value for unabsorbed AGN (Nandra et al. \cite{Nandra}; |
Caccianiga et al. 2001)). | Caccianiga et al. \cite{Caccianiga}) ). |
The best fit unfolded spectra aud residuals are shown in Figure 2 while the relevaut best fit parameters. quoted in the rest frame. are sununarized in Table 3 along with the 210 keV fiuxes aud the intrinsic bhuuinosities. | The best fit unfolded spectra and residuals are shown in Figure 2 while the relevant best fit parameters, quoted in the rest frame, are summarized in Table 3 along with the 2–10 keV fluxes and the intrinsic luminosities. |
The errors reported in the table also include the uncertaiutv ou the photometrie redshifts (80.2). | The errors reported in the table also include the uncertainty on the photometric redshifts $\pm$ 0.2). |
A pure thermal coniponent is rejected for all the sources at more than confidence level aud the addition of a thermal componcu othe powerlaw inoclel is. not statisticallyDal required. | A pure thermal component is rejected for all the sources at more than confidence level and the addition of a thermal component to the power–law model is not statistically required. |
: These pesults are iui. agrecinent with⋅ the point⋅like appearance o. he »XvayHN emissionS iand with the lackax of evident massiveas: clustersdusters iniu the optical/nearical/uearmfrarec infrarednuages. | These results are in agreement with the point–like appearance of the X–ray emission and with the lack of evident massive clusters in the optical/near–infrared images. |
inaees, AzAs reporterreporte in ⋅↴Table 3.∙ the iutriusic⋅⋅⋅ column densities⋅⋅ and the "D210 xeV Tuuinosities,;D» even if uot extreme.x (€are indicativeH of he presence: of observed: AH-Ἡ7102 2) iand high5huuinosty⋅∙ (Lojpg>↽1055 erg 1) AGNs. ⇁⋅Le. X-ray “obscured QSO. | As reported in Table 3, the intrinsic column densities and the 2–10 keV luminosities, even if not extreme, are indicative of the presence of obscured $_H>10^{22}$ $^{-2}$ ) and high--luminosity $_{2-10\rm keV}>10^{44}$ erg $^{-1}$ ) AGNs, i.e. X–ray obscured QSO. |
yor All of. them appear exteuded in⋅ the optica⋅ and near-IR images.. ie. ⋅⋅in these bands we do not observe he strong unclear enhancement due to the 050.-- anc the emission is dominated by the stellar coutimmuin of he host galaxy. | All of them appear extended in the optical and near-IR images, i.e. in these bands we do not observe the strong nuclear enhancement due to the QSO and the emission is dominated by the stellar continuum of the host galaxy. |
These facts: suggest that the QSOsarora are likely to be strongly absorbed also in the optical/ucarinfrared domain. ic. they are X.ray absorbed type 2 QSO canclicates. | These facts suggest that the QSOs are likely to be strongly absorbed also in the optical/near--infrared domain, i.e. they are X–ray absorbed type 2 QSO candidates. |
The remaining 3 EROs (118. 193. 711) are detected only ia oue of the EPIC cameras (see Table 2) and the | The remaining 3 EROs (443, 493, 714) are detected only in one of the EPIC cameras (see Table 2) and the |
is largest when T;=Tia: this peak is evident in the top panel for L4/L4=0.01 aud in the bottom panel for LL—0.03. | is largest when $T_\circ=T_\mathrm{crust}$; this peak is evident in the top panel for $L_A/\LAo=0.01$ and in the bottom panel for $L_A/\LAo=0.03$. |
The rapid rise of οΓον in the bottom panel for L4/L=0.01 is jecause the neutrino cooling iu the crust and core goes to zero. so that all of the heat. geuerated in the crust flows outwards FFigure 11.. penet)). | The rapid rise of $dT_\mathrm{crust}/dT_\circ$ in the bottom panel for $L_A/\LAo =
0.01$ is because the neutrino cooling in the crust and core goes to zero, so that all of the heat generated in the crust flows outwards Figure \ref{fig:low-mdot-sf}, ). |
Iu addition. the crust iu the entire is. also. crystalline.. which reduces] dT⋅ /dp. | In addition, the crust in the entire region considered is also crystalline, which reduces $dT/dp$. |
In⋅general. for AL >( 10]PAL./'egionvr.|.consideredthe eniperature in the crust becomes independent of the temperature in the atmosphere aud upper ocean of the neutron star. | In general, for $\dot{M}\gtrsim
10^{-9}\Msun/\yr^{-1}$, the temperature in the crust becomes independent of the temperature in the atmosphere and upper ocean of the neutron star. |
Au interesting possibility for a rapidly accreting neutron star is that its crust may melt. | An interesting possibility for a rapidly accreting neutron star is that its crust may melt. |
This happens wherever D35170. where the exact value is uncertain (forareview.see[chimaru1982).. | This happens wherever $\Gamma\lesssim 170$, where the exact value is uncertain \citep[for a review,
see][]{ichimaru82:_stron}. |
Siuce I use the formulation of Farouki&Hamaguchi(1993) to calculate the ionic (ree energy. I | Since I use the formulation of \citet{farouki93} to calculate the ionic free energy, I |
forrmation. | rmation. |
Figures 1l and 2 show the observed light) variations ol 87 Tau together with the I3-frequeney. fiti derived. in later sections. | Figures 1 and 2 show the observed light variations of $\theta^2$ Tau together with the 13-frequency fit derived in later sections. |
Phe new data will be made available in the Communications in Asteroscismology. | The new data will be made available in the Communications in Asteroseismology. |
The pulsation frequeney analyses were performed. with a package of computer programs with sinele-frequcney and multiple-frequeney techniques (PERLODOS. Sperl 1998). | The pulsation frequency analyses were performed with a package of computer programs with single-frequency and multiple-frequency techniques (PERIOD98, Sperl 1998). |
These programs utilize Fourier as well as algorithms. | These programs utilize Fourier as well as algorithms. |
The latter technique fits a number of simultaneous sinusoidal variations to the observed. light variability. ancl does not rely on. prewhitening. | The latter technique fits a number of simultaneous sinusoidal variations to the observed light variability and does not rely on prewhitening. |
For the purposes of presentation and initial searches. however. prewhitening is required if the Iow-amplitude modes are to be seen. | For the purposes of presentation and initial searches, however, prewhitening is required if the low-amplitude modes are to be seen. |
Therefore. in the presentation of the results. (see below). the various power spectra are presented as a series of panels. each with additional frequencies removed relative to the panel above. | Therefore, in the presentation of the results (see below), the various power spectra are presented as a series of panels, each with additional frequencies removed relative to the panel above. |
Two observatories also. provided. measurements in a second filter. viz.. | Two observatories also provided measurements in a second filter, viz., |
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