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these spectra are shown in the same format as in Figs. | these spectra are shown in the same format as in Figs. |
12 and 13.. | \ref{smap} and \ref{smap_epic}. |
In the rangeAA,, the spectrum shows many features which are difficult to identify. | In the range, the spectrum shows many features which are difficult to identify. |
Wavelengths of a few known lines are marked. | Wavelengths of a few known lines are marked. |
A weak feature coincides with the position of the resonance doublet but, given the presence of many similar features in the spectrum as well as the uncertainty of the absolute wavelength scale (which depends on the accurate centroiding of the zero order), its identification is uncertain. | A weak feature coincides with the position of the resonance doublet but, given the presence of many similar features in the spectrum as well as the uncertainty of the absolute wavelength scale (which depends on the accurate centroiding of the zero order), its identification is uncertain. |
The wavelength range of the OM grisms overlaps with the long-wavelength range of the International Ultraviolet Explorer (IUE). | The wavelength range of the OM grisms overlaps with the long-wavelength range of the International Ultraviolet Explorer (IUE). |
Five spectra were taken with the IUE during the outburst in 1979: on June 28, June 30, July 2, July 4,and July 10, corresponding to 5, 7, 8, 10, and 16 days after outburst (?).. | Five spectra were taken with the IUE during the outburst in 1979: on June 28, June 30, July 2, July 4,and July 10, corresponding to 5, 7, 8, 10, and 16 days after outburst \citep{williams81}. |
In the top panel of Fig. 15,, | In the top panel of Fig. \ref{smap_om}, |
the IUE spectra from 8 and 16 days after the 1979 outburst are overplotted in orange and blue (see right legend). | the IUE spectra from 8 and 16 days after the 1979 outburst are overplotted in orange and blue (see right legend). |
The strongest lines in these spectra are, apart from the doublet, at 2512, 2734 and AA,, at aand v]] at2784 and AA.. | The strongest lines in these spectra are, apart from the doublet, at 2512, 2734 and , at and ] at2784 and . |
from [Myr through 13 Cir. including gaseous emission. which signifcantlv affects broad band huuinosities and colours during carly evolutionary stages (see 7 for details). | from 4 Myr through 13 Gyr, including gaseous emission, which significantly affects broad band luminosities and colours during early evolutionary stages (see \cite*{AndersFritze03} for details). |
Spectra are then olded with filter unctious for any desired filter svstem to vield the photometric evolution. | Spectra are then folded with filter functions for any desired filter system to yield the photometric evolution. |
This is important in order to avoid uucertaimties from transformations )otwoeen filter svstems. | This is important in order to avoid uncertainties from transformations between filter systems. |
Models well reproduce eiipirical colour metallicity calibrations over their rauge of validity aud indicate significaut deviations from their inear behaviour towards ligher metallictics. | Models well reproduce empirical colour – metallicity calibrations over their range of validity and indicate significant deviations from their linear behaviour towards higher metallicties. |
We showed hat transformations from colour to metallicity are sienificautly age-dependent and that transformations roni colour to age are siguificautlv inctallicitv-depeudenutlg (7). | We showed that transformations from colour to metallicity are significantly age-dependent and that transformations from colour to age are significantly metallicity-dependent \citep{Schulz+02}. |
. The effect of dust absorption is included iu CGALEV iodels asstuning a starburst extinction law (2) for a range of values for E(B η...Vjx Lanag) | The effect of dust absorption is included in GALEV models assuming a starburst extinction law \citep{Calzetti+00} for a range of values for $E(B-V)$ $0 \leq E(B-V) \leq 1~$ mag). |
GALEV models also iuclude the full set of Lick spectral absorption mdices ou the basis of eimipirical calibratious for the iudices iu ternis of stellar pareuneters for every individual cluster star Tog.logο,[FeΤΠ) as given by (2). and (?).. | GALEV models also include the full set of Lick spectral absorption indices on the basis of empirical calibrations for the indices in terms of stellar parameters for every individual cluster star ${\rm T_{eff},~log~g,~[Fe/H]}$ as given by \citep{Gorgas+93} and \citep{Worthey+94}. |
We showed that the transformation from the age-scusitive Lick index IL; to age is significantly: inetallicity-depeudeut aud that the traustormation from the metallicityv-sensitive Lick indices (Mgb. Me». |MgFe]. ...) to metallicity is age-dependent for ages x:10 Cr (77). | We showed that the transformation from the age-sensitive Lick index $_{\beta}$ to age is significantly metallicity-dependent and that the transformation from the metallicity-sensitive Lick indices (Mgb, $_2$, [MgFe], $\dots$ ) to metallicity is age-dependent for ages $\leq 10$ Gyr \citep{Kurth+99,LillyFritze06}. |
Qur analysis methods use the full iuforxiuation from multi-baud imagine (UV.U.B..... ΑΠ) or/aud Lick spectroscopy available for à SC system. compare thon toa large grid of over 100.000 CALEV models in terius of Spectral Encrey Distributions (SEDs). Lick iudices. or a conibiuation of both (cf. ?.. 7.. Lilly Fritze | Our analysis methods use the full information from multi-band imaging $UV,~U,~B,\dots,~NIR$ ) or/and Lick spectroscopy available for a SC system, compare them to a large grid of over 100.000 GALEV models in terms of Spectral Energy Distributions (SEDs), Lick indices, or a combination of both (cf. \cite*{Anders+04a}, \cite*{LillyFritze06}, |
2008. submitted), | Lilly Fritze 2008, ). |
SEDs. we recall. are sects of magnitudes in a muuber of filters from short to loug wavelengths. c.g. U.2.K. | SEDs, we recall, are sets of magnitudes in a number of filters from short to long wavelengths, e.g. $U ~ \dots ~ K$. |
Our analysis tools not oulv determinethe best fit model but attribute probabilities to all models that allow us to determine the lo uncertainties for all the SC parameters they return: age. metallicity. E(BV and mass. | Our analysis tools not only determine best fit model but attribute probabilities to all models that allow us to determine the $1 \sigma$ uncertainties for all the SC parameters they return: age, metallicity, $E(B-V)$, and mass. |
Extensive tests with artificial SCs have shown that CV or Ü/— baud observations are essential for age dating of YSCs aud a NIR baud is oeuportaut to obtain accurate moetallicities. | Extensive tests with artificial SCs have shown that $UV$ or $U-$ band observations are essential for age dating of YSCs and a NIR band is important to obtain accurate metallicities. |
For YSCs oe1 dusty galaxies four passbauds inchiding CV/U aud IT ov AK with observational uucertaiuties <0.05 mae oei the UV/optical and <0.1 mae in the NIR allow +to largely disentangle ages aud aud metallicities aud to jbtain ages to Aage/agex0.3 and iietallicties to £0.2 ex. | For YSCs in dusty galaxies four passbands including $UV/U$ and $H$ or $K$ with observational uncertainties $\leq 0.05$ mag in the UV/optical and $\leq 0.1$ mag in the NIR allow to largely disentangle ages and and metallicities and to obtain ages to $\Delta~{\rm age / age} \leq 0.3$ and metallicties to $\pm 0.2$ dex. |
For intermediate-age SCs or old GCs in dust-free environnieuts. three passbands. again rangiug from C or D through fT or A are enough (?7).. | For intermediate-age SCs or old GCs in dust-free environments, three passbands, again ranging from $U$ or $B$ through $H$ or $K$ are enough \citep{Anders+04a,deGrijs+03c}. |
Applving our SED analvsis tool to TST WEPC? aud NICMOS archival data for some 170. compact YSC's that we identified in the not apparently interacting dwart starburst galaxy NGC 1569. we obtained masses for the bulk of its YSCs iu the range 107LotML.. | Applying our SED analysis tool to HST WFPC2 and NICMOS archival data for some 170 compact YSCs that we identified in the not apparently interacting dwarf starburst galaxy NGC 1569, we obtained masses for the bulk of its YSCs in the range ${\rm 10^3 - 10^4~M_{\odot}}$. |
. Onulv a handful of these. including the 3 previously known so-called Super Star Clusters. lave asses above a few lo’AL... ic. in the range of GC masses. | Only a handful of these, including the 3 previously known so-called Super Star Clusters, have masses above a few ${\rm 10^5~M_{\odot}}$, i.e. in the range of GC masses. |
We conclude that this strougly starbursting but nof apparently interacting dwirf galaxy docs not fori iu new CCS. or. at most. very few (7?).. | We conclude that this strongly starbursting but not apparently interacting dwarf galaxy does not form any new GCs, or, at most, very few \citep{Anders+04b}. |
For the starburst iu the massive eas-vich spiral — spiral merecr remnant NGC 7252. we could estimate the SE efficiency very conservatively to be at least 35 | For the starburst in the massive gas-rich spiral – spiral merger remnant NGC 7252, we could estimate the SF efficiency very conservatively to be at least 35. |
This estimate was based on the amount of new stars formed durug the burst. as obtained from the deep Daher absorption lines in the overall spectrum. aud a vorv eenerous estimate of the σας mass available in the two Se-type progenitor spirals. of which the wuple III still observed all along the extended tidal ails is the proof (77).. | This estimate was based on the amount of new stars formed during the burst, as obtained from the deep Balmer absorption lines in the overall spectrum, and a very generous estimate of the gas mass available in the two Sc-type progenitor spirals, of which the ample HI still observed all along the extended tidal tails is the proof \citep{FG94a,FG94b}. |
Such a high SF efficiency should allow for the formation of massive. conipact. stronelv youd GCs. | Such a high SF efficiency should allow for the formation of massive, compact, strongly bound GCs. |
OST observations indeed revealed a rich o)pulatiou of compact SCs with ages iu the rauge 600900 Myr and metallicities close to solar (?).. | HST observations indeed revealed a rich population of compact SCs with ages in the range $600-900$ Myr and metallicities close to solar \citep{FB95}. |
They appareutly have survived may internal crossing iues and the most critical phase in their lives. the infant mortality phase after expulsion of the eas left over at their formation when the first SNe weut off. aud μον are still compact aud bound. | They apparently have survived many internal crossing times and the most critical phase in their lives, the infant mortality phase after expulsion of the gas left over at their formation when the first SNe went off, and they are still compact and bound. |
This is particularly nmupressive since all this happened during the violent relaxation phase that restructured their parent galaxy roni two spiral disks iuto a spherical configuratio- eaturug a de Vaucouleurs profile (?).. | This is particularly impressive since all this happened during the violent relaxation phase that restructured their parent galaxy from two spiral disks into a spherical configuration featuring a de Vaucouleurs profile \citep{Schweizer06}. |
These vouug GCs have all chances to survive another Hubble time. | These young GCs have all chances to survive another Hubble time. |
Thev have masses in the rauge Lo?10°AD. with cluster W3 even reaching 7SAL. (2?).. | They have masses in the range ${\rm 10^5 - 10^6~M_{\odot}}$ with cluster W3 even reaching ${\rm 7-8~M_{\odot}}$ \citep{FB95,Maraston+04}. |
Enough of tose voung GCs survived until today to secure the ucrecr remnant the typical GC specific frequency of an liptical galaxy. which is twice as high when defined in exis of umber of GCs in relation to galaxy total as for an average spiral (2).. | Enough of those young GCs survived until today to secure the merger remnant the typical GC specific frequency of an elliptical galaxy, which is twice as high when defined in terms of number of GCs in relation to galaxy total as for an average spiral \citep{ZepfAshman93}. . |
Le. during the strong global starburst accommpanving the merger that transformed wo bright Sc galaxies iuto a dynamically still voune elliptical. a umber of secondary. GCs has formed that is comparable to the uuuber of preexisting GCs iu both xogenitor spirals. | I.e. during the strong global starburst accompanying the merger that transformed two bright Sc galaxies into a dynamically still young elliptical, a number of secondary GCs has formed that is comparable to the number of preexisting GCs in both progenitor spirals. |
Bethe&Brown(1998) SUENOOested that merecrs resulting iu short-hard eamuua-ray bursts would be mainly those of LMDITI-NS binaries. with those of NS-NS binarics down bv an order of magnitude from these. | \citet{Bet98} suggested that mergers resulting in short-hard gamma-ray bursts would be mainly those of LMBH-NS binaries, with those of NS-NS binaries down by an order of magnitude from these. |
The lower umber of the latter resulted from the necessity that the two eiaut progenitors be within of cach other in ZAMS ass so that tjov dDrued He at the same time. | The lower number of the latter resulted from the necessity that the two giant progenitors be within of each other in ZAMS mass so that they burned He at the same time. |
Otherwise the first bori pulsar would fux itself in the red elaut euvelope of the coupalion giant as it evolved aud accrete cuough matter to eo into a black hole (BID). | Otherwise the first born pulsar would find itself in the red giant envelope of the companion giant as it evolved and accrete enough matter to go into a black hole (BH). |
Brown(1995) had already estimated this to be true. based on Chevalicr(1993).. and proposed the special wav that two eiauts burn Te at the sale tdue. 1 order to avoid the red giant common envelope evolution of the first 20111 pulsar. | \citet{Bro95} had already estimated this to be true, based on \citet{Che93}, and proposed the special way that two giants burn He at the same time, in order to avoid the red giant common envelope evolution of the first born pulsar. |
If the two eiauts burn Ile at he sane nuc. the two Ue euveopes are assured to eo iuto common euvelope evolution. expelling the conmnaiou envelo]© laatter so that the hel ewelope is lost from cach star. | If the two giants burn He at the same time, the two He envelopes are assumed to go into common envelope evolution, expelling the common envelope matter so that the helium envelope is lost from each star. |
Draun&Langer(1995) showed that there was not sufficieu tuue for either of the stars fo acrete the comnon envelope He. so hat if the two Πο stars had to be nearly equal iu nass. then their progenitor eiauts ust also 0. | \citet{Bra95} showed that there was not sufficient time for either of the stars to accrete the common envelope He, so that if the two He stars had to be nearly equal in mass, then their progenitor giants must also be. |
From the Schalleretal.(1992) miodels for the eiaut progenitors ¢f the neutron stars we coisicler. t1ο elauts have o be within <1% of¢"ch other iu mass in order to buru Πο at the saue time. | From the \cite{Sch92} models for the giant progenitors of the neutron stars we consider, the giants have to be within $\lsim 4\%$ of each other in mass in order to burn He at the same time. |
Of the Bethe&Brown(7908) inerecr 1ate of LOtyr tinor Galaxy. only ~0.1. or 105 1 were estimated o be those of binary NS | Of the \cite{Bet98} merger rate of $10^{-4}$ $^{-1}$ in our Galaxy, only $\sim 0.1$, or $10^{-5}$ $^{-1}$ were estimated to be those of binary NS's. |
A recent detailerL calculation by Dewietal.(2006) οἶνος a ποσο rate of 0.112 |l for Drowu's specia scenario. the upper exd of the calculation iu agreement wih (1998). | A recent detailed calculation by \citet{Dew96} gives a merger rate of $0.1-12$ $^{-1}$ for Brown's special scenario, the upper end of the calculation in agreement with \cite{Bet98}. |
. The later authors simply estnauated hat the ολα]ing mergers would be of LMDBII-NS binarics since they calculated that when the pulsar went through ¢Onmuuon eanveloe with the companion star it accretec LAL. of matter. choueh to seudi iuto a DIT. | The latter authors simply estimated that the remaining mergers would be of LMBH-NS binaries since they calculated that when the pulsar went through common envelope with the companion star it accreted $\sim 1\msun$ of matter, enough to sendit into a BH. |
The iuo:ut of accYetion was corrected downwards ~nU& by removal of an approximation of Betlu&Brown(1998) bv Delezvüskicta.(2002). | The amount of accretion was corrected downwards $\sim 25\%$ by removal of an approximation of \citet{Bet98} by \citet{Bel02}. |
. In tie present note we try to make a roa calculation of the LMDII-NS binary. NS-NS binary ratio. | In the present note we try to make a real calculation of the LMBH-NS binary, NS-NS binary ratio. |
Pinsomeault&Stauek(2006) assembled evidence that 7Binaries like to be Twius. | \citet{Pin06} assembled evidence that “Binaries like to be Twins". |
They showed that a recently publishedd sauple of 21 detached eclipsing binaries iu the Simall Magellanic Cloud can be evolve in terms of a flat luass fujction ccontaiuiue of the svsteis aud a “twins” population with q>0.95 containing the remainder. | They showed that a recently published sample of 21 detached eclipsing binaries in the Small Magellanic Cloud can be evolved in terms of a flat mass function containing of the systems and a “twins" population with $q> 0.95$ containing the remainder. |
All of the binaries had orbital period P<5 davs. wih primary maasses GOAL.<AL,<212Ab.e"n | All of the binaries had orbital period $P< 5$ days, with primary masses $6.9 \msun <M_1 <27.3 \msun$. |
The iuporaut role of twins is that he two «oκlants are close! enoneh im mass that 1 BreWAL(1995). scenario they Call CVO.ve iuto NS-NS binarics. whereas if they are further apart iu mass they will evolve in oa LMDII-NS binary (CCrevalicrJ993:Bethe&Brown1998). | The important role of twins is that the two giants are close enough in mass that in \citet{Bro95} scenario they can evolve into NS-NS binaries, whereas if they are further apart in mass they will evolve into a LMBH-NS binary \citep{Che93,Bet98}. |
. Thus the twins may increase the umber of NS-NS binaries. | Thus the twins may increase the number of NS-NS binaries. |
We sugeOOest that the resulting ΠΡΟ of short hard eamuna-rav bursts. which result from the iiereiug of the binaricss. Which to date are muae to differentiate between the two species. wav uot be changed much. some of the predictec large excess of LAIBII-NS biuarcSs appearing rather as NS-NS binaries. | We suggest that the resulting number of short hard gamma-ray bursts, which result from the merging of the binaries, which to date are unable to differentiate between the two species, may not be changed much, some of the predicted large excess of LMBH-NS binaries appearing rather as NS-NS binaries. |
However. because the latter are so mich more easy. to observe. the role betwee ilt we see aud what is present will be tightened. | However, because the latter are so much more easy to observe, the role between what we see and what is present will be tightened. |
Tn Sec. | In Sec. |
?7 we shall show that evoluion of binaries with a flat mass evolution does. i agreement with produce a ratio of |BID | NS) svsteuis of 5. for the population that does not contain twins. | \ref{sec-twin} we shall show that evolution of binaries with a flat mass evolution does, in agreement with \citet{Pin06}
produce a ratio of $+$ $+$ NS) systems of 5, for the population that does not contain twins. |
This is half the order of magnitude raio found by Bethe&Brown (1998).. | This is half the order of magnitude ratio found by \citet{Bet98}. . |
The lower value results from the fact that | The lower value results from the fact that |
In this subsection we explore the behavior of this system by direct numerical simulations. | In this subsection we explore the behavior of this system by direct numerical simulations. |
We found this to be helpful in the building of our intuition. | We found this to be helpful in the building of our intuition. |
We defer a semi-analytical normal-mode analysis to the next subsection. | We defer a semi-analytical normal-mode analysis to the next subsection. |
We follow LO7 and for concreteness concentrate on a specific example; it will be clear that the conclusions we reach are general. | We follow L07 and for concreteness concentrate on a specific example; it will be clear that the conclusions we reach are general. |
We choose wo= lrad/second and mass M=1. | We choose $\omega_0 = 1$ rad/second and mass $M = 1$. |
We choose a total number of 1000 small pendulae with frequencies ων=(0.5+ n/1000)rad/second and masses m,=m10-7, to mimic the continuum frequency range between 0.5rad/second and 1.5rad/second. | We choose a total number of 1000 small pendulae with frequencies $\omega_n = (0.5 +
n/1000)$ rad/second and masses $m_n=m=10^{-4}$, to mimic the continuum frequency range between $0.5$ rad/second and $1.5$ rad/second. |
'The simulation is initiated by displacing the large oscillator while keeping the small pendulae relaxed (this mimics the stresses in the crust), and then releasing. | The simulation is initiated by displacing the large oscillator while keeping the small pendulae relaxed (this mimics the stresses in the crust), and then releasing. |
The subsequent motion of the system is then followed numerically by using a second order leapfrog integration scheme which conserves the energy with high precision. | The subsequent motion of the system is then followed numerically by using a second order leapfrog integration scheme which conserves the energy with high precision. |
The resulting motion of the large pendulum can be decomposed into three stages (see Fig. | The resulting motion of the large pendulum can be decomposed into three stages (see Fig. |
2 and Fig. 3)): ( | \ref{Fig2A} and Fig. \ref{Fig2B}) ): ( |
1) During the first 50-60 seconds, there is a rapid exponential decay of the large oscillator’s motion, during which most of the energy is transferred to the multitude (i.e., the continuum’) of small oscillators. | 1) During the first 50-60 seconds, there is a rapid exponential decay of the large oscillator's motion, during which most of the energy is transferred to the multitude (i.e., the 'continuum') of small oscillators. |
This is the so-called phenomenon of “resonant absorption”, which has been studied for decades in the MHD and plasma physics community (e.g., Ionson 1978, Hollweg 1987, Goedbloed Poedts 2004, L07, Gruzinov 2008b). | This is the so-called phenomenon of “resonant absorption”, which has been studied for decades in the MHD and plasma physics community (e.g., Ionson 1978, Hollweg 1987, Goedbloed Poedts 2004, L07, Gruzinov 2008b). |
In this first stage, the amplitude of the big pendulum motions drops by a factor of ~100. ( | In this first stage, the amplitude of the big pendulum motions drops by a factor of $\sim 100$. ( |
2) After ~60 seconds, the exponential decay stops abruptly as the large oscillator now reacts to the collective pull of the small ones. | 2) After $\sim 60$ seconds, the exponential decay stops abruptly as the large oscillator now reacts to the collective pull of the small ones. |
This second stage is characterized by a slow algebraic decay of the amplitude of the big pendulum displacement. | This second stage is characterized by a slow algebraic decay of the amplitude of the big pendulum displacement. |
Gruzinov (2008b) explains this as being due to the branch cut in the oscillator's response function. ( | Gruzinov (2008b) explains this as being due to the branch cut in the oscillator's response function. ( |
3) The motion of the large oscillator stabilizes at a constant level (L07 missed this stage in his simulations, which he stopped too early). | 3) The motion of the large oscillator stabilizes at a constant level (L07 missed this stage in his simulations, which he stopped too early). |
Fourier transform reveals 2 QPOs at the frequencies close to the continuum edges, w—0.5 and w=1.5; the same QPO frequencies can be observed in the previous stage (2) as well. | Fourier transform reveals 2 QPOs at the frequencies close to the continuum edges, $\omega=0.5$ and $\omega=1.5$; the same QPO frequencies can be observed in the previous stage (2) as well. |
What is the origin of the QPOs, and how is this eventual stability established? | What is the origin of the QPOs, and how is this eventual stability established? |
In Fig. | In Fig. |
4 and 5,, we show how the amplitude of the small oscillators evolves with time. | \ref{Fig5A} and \ref{Fig5AA}, we show how the amplitude of the small oscillators evolves with time. |
After the initial resonant absorption phase, the amplitude is distributed as a Lorentzian centered on the frequency around w=1; this is because the small oscillators in resonance with the large one are the ones which gain the most energy. | After the initial resonant absorption phase, the amplitude is distributed as a Lorentzian centered on the frequency around $\omega=1$; this is because the small oscillators in resonance with the large one are the ones which gain the most energy. |
However, in subsequent times we see that the energy exchange occurs between the smalllators?,, and that the net result of this exchange is the energy flow towards the oscillators whose frequencies are near the edges. | However, in subsequent times we see that the energy exchange occurs between the small, and that the net result of this exchange is the energy flow towards the oscillators whose frequencies are near the edges. |
By the time the third stage begins, the amplitudes of the oscillators near the edge stabilize and their phases become locked. | By the time the third stage begins, the amplitudes of the oscillators near the edge stabilize and their phases become locked. |
They are pulling and pushing the large oscillator in unison. | They are pulling and pushing the large oscillator in unison. |
In the next subsection, we | In the next subsection, we |
Cen A. We show that (vpical blazar-like jet parameters may be used to model the broadband SED. if one allows for an additional cascade contribution to the 5-rav. emission due lo 55 absorption and cascading in the thermal infrared radiation field of the prominent dust emission known to be present in Cen A. | Cen A. We show that typical blazar-like jet parameters may be used to model the broadband SED, if one allows for an additional cascade contribution to the $\gamma$ -ray emission due to $\gamma\gamma$ absorption and cascading in the thermal infrared radiation field of the prominent dust emission known to be present in Cen A. |
binaries that contain significantly evolved companion stars. transferring mass on a time-scale governed by nuclear expansion of the secondary. | binaries that contain significantly evolved companion stars, transferring mass on a time-scale governed by nuclear expansion of the secondary. |
It is no surprise. then. that GRS | 105. with its long orbital period and huge accretion dise is one of the few svstems in which such behaviour has been observed. | It is no surprise, then, that GRS $+$ 105, with its long orbital period and huge accretion disc is one of the few systems in which such behaviour has been observed. |
MRT acknowledges a PPARC postdoctoral fellowship. | MRT acknowledges a PPARC postdoctoral fellowship. |
The simulations were performed on the UK Astrophysical Fluids Facility (UKAFF). | The simulations were performed on the UK Astrophysical Fluids Facility (UKAFF). |
The authors are grateful to Chris Done and also to the referee for helpful and insightful comments. | The authors are grateful to Chris Done and also to the referee for helpful and insightful comments. |
The light curve was provided by the RXTE/ASM team at MIT andRXTE SOF and GOF at NASA Goddard SFC. | The light curve was provided by the /ASM team at MIT and SOF and GOF at NASA Goddard SFC. |
or ongoing gas accretion in nearby galaxies(Sancisietal.2008). | for ongoing gas accretion in nearby galaxies\citep{Sancisi08}. |
. However. estimates of the accretion rates in the form of neutral ivdrogen in nearby spirals consistently give values much lower han those required to sustain star formation at their observed rates (See Fraternali (2010) for a recent review). | However, estimates of the accretion rates in the form of neutral hydrogen in nearby spirals consistently give values much lower than those required to sustain star formation at their observed rates (see Fraternali (2010) for a recent review). |
This leads to the iypothesis that most of the baryons that reside outside galaxies at he present are in the form of ionized gas (e.g. Fukugita. Hogan Peebles 1998). | This leads to the hypothesis that most of the baryons that reside outside galaxies at the present are in the form of ionized gas (e.g. Fukugita, Hogan Peebles 1998). |
There is also some evidence from observations of ionized silicon in high and intermediate-velocity clouds in the Tilkv Way. that gas in a low-metallicity ionized phase in the halo can provide a substantial CI M./yr) cooling inflow to replenish star ormation in the disk (Shull et al. | There is also some evidence from observations of ionized silicon in high and intermediate-velocity clouds in the Milky Way, that gas in a low-metallicity ionized phase in the halo can provide a substantial (1 $_{\odot}$ /yr) cooling inflow to replenish star formation in the disk (Shull et al. |
2009). | 2009). |
One interesting question that has not been considered. very much in the literature. is whether gas accretes onto all galaxies in à smooth. continuous fashion. or whether accretion is more stochastic. | One interesting question that has not been considered very much in the literature, is whether gas accretes onto all galaxies in a smooth, continuous fashion, or whether accretion is more stochastic. |
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