source
stringlengths 1
2.05k
⌀ | target
stringlengths 1
11.7k
|
|---|---|
All of the flux is dissipated above the effective photosphere itself. and the color-teiiperature correction becomes inch lareer. | All of the flux is dissipated above the effective photosphere itself, and the color-temperature correction becomes much larger. |
We repeat the spectral fitting. replaciug the simple model bv the more physical spectimu2005). | We repeat the spectral fitting, replacing the simple model by the more physical spectrum. |
. Likediskbb. is a iuulticolor dackbody. model. but it includes all of the relativistic effects modeled in i.c. all the special aud general relativistic sincaring of the intrinsic disk cussion due Oo rapid rotation in strong eravity. | Like, is a multicolor blackbody model, but it includes all of the relativistic effects modeled in i.e. all the special and general relativistic smearing of the intrinsic disk emission due to rapid rotation in strong gravity. |
These broaden he spectrum so that if is not so sharply peaked as sedicted in the models. but unlike it assunes that the intrinsic cussion from cach radius is just a color-temperature corrected blackbody. and hat this color correction factor is the same at all radi. | These broaden the spectrum so that it is not so sharply peaked as predicted in the models, but unlike it assumes that the intrinsic emission from each radius is just a color-temperature corrected blackbody, and that this color correction factor is the same at all radii. |
Therefore. the derived fo) will arise solely from differences in the treatineut of the surface cussion2006). | Therefore, the derived $f_{\rm col}$ will arise solely from differences in the treatment of the surface emission. |
. We fix the mass. distance. inclination and spin. and fit for the mass accretion rate aud color-tempcrature correction assuimiug no stress ou the inner boundary and uo returning radiation. mt iucluding limb darkening. | We fix the mass, distance, inclination and spin, and fit for the mass accretion rate and color-temperature correction assuming no stress on the inner boundary and no returning radiation, but including limb darkening. |
Fig. | Fig. |
5. shows these for he CCD (loft panel) aud proportional counter (right xuiel) data respectively, | \ref{f:fcol1} shows these for the CCD (left panel) and proportional counter (right panel) data respectively. |
Again the three dense disk xescriptious all show similar results. with the alpha disk with a=0.1 giving lieher color-teniperature correction at hieh ποπτν, as before. | Again the three dense disk prescriptions all show similar results, with the alpha disk with $\alpha=0.1$ giving higher color-temperature correction at high luminosity, as before. |
However. the values of he derived color-teniperature corrections aredifferent o those derived from the approach. | However, the values of the derived color-temperature corrections are to those derived from the approach. |
The color-eniperature correction for does stay constant in the CCD bandpass. appareutlv iu conflict with the observed LXT! velation fromdiskbb. | The color-temperature correction for does stay constant in the CCD bandpass, apparently in conflict with the observed $L\propto T^4$ relation from. |
It increases roni 1.61.9. which at least goes through the fio)L.7 constant value seen withdiskbb. | It increases from 1.6–1.9, which at least goes through the $f_{\rm col}=1.7$ constant value seen with. |
By contrast. the xoportional counter bandpass gives values of the color-teniperature correction which are always lower than hose seeu fromdiskbb. as well as spauning a wider range. frou, 1.351.5 The auswer to these issues lies im the detailed shape of the spectrum. | By contrast, the proportional counter bandpass gives values of the color-temperature correction which are always lower than those seen from, as well as spanning a wider range, from 1.35–1.8 The answer to these issues lies in the detailed shape of the spectrum. |
Fig. | Fig. |
laa shows the initial disk (blue. dotted line) and its best fitdiskbb spectra o the proportional counter (red. solid line} aud. CCD (eveen. dashed Ime) baudpasses for one of the deuse disk xescriptions (a= 0.01) for an iuput luminosity of log/= 0.3. | \ref{f:spec}a a shows the initial disk (blue, dotted line) and its best fit spectra to the proportional counter (red, solid line) and CCD (green, dashed line) bandpasses for one of the dense disk prescriptions $\alpha=0.01$ ) for an input luminosity of $\log l=-0.3$ . |
The higher energv baudpass of the proportional counter weights the ft to higher temperatureslower jorialization than those for the CCD. and the nodel is a good fit to the above the 3 keV lower init of the proportional counter bandpass. | The higher energy bandpass of the proportional counter weights the fit to higher temperatures/lower normalization than those for the CCD, and the model is a good fit to the above the 3 keV lower limit of the proportional counter bandpass. |
However. the vest fit to the CCD data does not look so compelling. | However, the best fit to the CCD data does not look so compelling. |
It is clear that the spectrmu is narrower than he diskcussion. aud the best ΠΕ matches this oulv iu the 1-3 keV rauge where the signal-to-noise is naxinuun. | It is clear that the spectrum is narrower than the diskemission, and the 'best fit' matches this only in the 1-3 keV range where the signal-to-noise is maximum. |
At higher energies. the Compton tail provides additional freedom to match the data. | At higher energies, the Compton tail provides additional freedom to match the data. |
Fie. | Fig. |
Ibb shows the models with mass accretion rate chosen to give loge?=0.3 for a black hole with a.=0.5 (ie. 8.537.1025 efs). with color-temperatiure correction of 1.7 for the proportional counter (red). | \ref{f:spec}b b shows the models with mass accretion rate chosen to give $\log l=-0.3$ for a black hole with $a_*=0.5$ (i.e. $8.537\times 10^{18}$ g/s), with color-temperature correction of 1.7 for the proportional counter (red). |
The ratio of best fit tempcratures from the models above iuplies that the color-tempcrature correction should be a factor 1.15 smaller in the CCD compared to the proportional counter ie. eivine fi.=Le for the CCD (blue). | The ratio of best fit temperatures from the models above implies that the color-temperature correction should be a factor 1.15 smaller in the CCD compared to the proportional counter i.e. giving $f_{col}=1.47$ for the CCD (blue). |
However. the spectrmu is mich broader than as it includes the relativistic sincaring of the continua. aud the proportional counter | However, the spectrum is much broader than as it includes the relativistic smearing of the continuum, and the proportional counter |
The formation ο [galaxies and their systems (Chaloes) is known to be more intricate than its simplified. rendering in the context of spheric‘al collapse models. | The formation of galaxies and their systems (`haloes') is known to be more intricate than its simplified rendering in the context of spherical collapse models. |
The process is characterised by gradual growth of the svstem mass and evolution of its morphological properties through multiple merging events. evidence for which comes from observations and hyerodsnamical simulations. | The process is characterised by gradual growth of the system mass and evolution of its morphological properties through multiple merging events, evidence for which comes from observations and hydrodynamical simulations. |
In custers. mergers allect also the evolution of intracluster (10) gas density and temperature and their spatial profiles. | In clusters, mergers affect also the evolution of intracluster (IC) gas density and temperature and their spatial profiles. |
On the theoretical side. studies of halo mergers and. related. issues began with the works of Bond et al. ( | On the theoretical side, studies of halo mergers and related issues began with the works of Bond et al. ( |
1991) and Lacey Cole (1f993). who developed the theory of in the context of structure formation. | 1991) and Lacey Cole (1993), who developed the theory of in the context of structure formation. |
This approach was originally devise bv Bond et al. | This approach was originally devised by Bond et al. |
in oxer to address the "cloud-in-cloud problem. who showed hat the Press Schecier (1974) mass function. including the 7Bidee factor” of 2. could be derived: under certain assumpions. | in order to address the “cloud-in-cloud” problem, who showed that the Press Schechter (1974) mass function, including the “fudge factor” of 2, could be derived under certain assumptions. |
Thev also used t1011’ formalism to derive expressions for mereer probabilities. | They also used their formalism to derive expressions for merger probabilities. |
Lacey Cole (hereafter LC) used the excursion set formalism (LESE) to extract such related quantities as hao merger rates. halo survival times. and halo formation times. | Lacey Cole (hereafter LC) used the excursion set formalism (ESF) to extract such related quantities as halo merger rates, halo survival times, and halo formation times. |
The NEW concentration parameter (Navarro. πο. White 1995). which characterises darκ matter (DAL) clistribu in à halo is known from N-body simulations to be correlated. with its formation time (e.g. Jing 2000: Bullock et al. | The NFW concentration parameter (Navarro, Frenk, White 1995), which characterises dark matter (DM) distribution in a halo is known from N-body simulations to be correlated with its formation time (e.g. Jing 2000; Bullock et al. |
2X Zhao et al. | 2001; Zhao et al. |
2003) due to the fact tha haloes which form earlier are likely to have more concensed. cores. rellecting higher background density of the universe. | 2003) due to the fact that haloes which form earlier are likely to have more condensed cores, reflecting the higher background density of the universe. |
We use this inferred correlation. and the probability cistribution function (IPDI!) | We use this inferred correlation, and the probability distribution function (PDF) |
Are there easier ways of predicting the observable properties of a model than generating and [fitting artificial X-ray spectra? | Are there easier ways of predicting the observable properties of a model than generating and fitting artificial X-ray spectra? |
“This method. is also instrument specific. as it depends on the spectral eharacteristies of the instrument used. ancl so predicted. results forROSAT ave not directly transferable to.ASCO. | This method is also instrument specific, as it depends on the spectral characteristics of the instrument used, and so predicted results for are not directly transferable to,. |
If we we can correctly estimate the absorbing column anc metallicity. either. [rom ai spectral. fit. or from observations at other wavelengths. then estimating a characteristic temperature is all we need clo. | If we we can correctly estimate the absorbing column and metallicity, either from a spectral fit or from observations at other wavelengths, then estimating a characteristic temperature is all we need do. |
We might expect the fitted temperature to be some ILux-weighted average of the true temperature range. and. the emission measure roughly the total emission over the range of temperaturesROSATL is sensitive to. | We might expect the fitted temperature to be some flux-weighted average of the true temperature range, and the emission measure roughly the total emission over the range of temperatures is sensitive to. |
Over the energy range 0.00 5-15.0keV the intrinsic Dux unabsorbed) weighted average temperature <Zi=210 keV. UV emission [rom very cool material at the shell-bubble interior interface. | Over the energy range $0.005$ $15.0 \keV$ the intrinsic flux unabsorbed) weighted average temperature $<T_{EW}> =
(4.9^{+2.7}_{-1.3}) \times 10^{-3} \keV$ , UV emission from very cool material at the shell-bubble interior interface. |
In. practice all the mid. and. extreme UV. emission. is absorbed. in he intervening 5M. | In practice all the mid and extreme UV emission is absorbed in the intervening ISM. |
In addition the instrument response "urther limits the "observed? temperature. | In addition the instrument response further limits the “observed” temperature. |
If we assume that we can accurately remove the effects of interstellar absorption on the X-rays fitting the column correctly) then to construct a predicted.ROSAT temperature we need only weight each temperature by the count rate in theROSALT PSPC due to cach temperature component. ignoring absorption. | If we assume that we can accurately remove the effects of interstellar absorption on the X-rays fitting the column correctly) then to construct a predicted temperature we need only weight each temperature by the count rate in the PSPC due to each temperature component, ignoring absorption. |
This temperature is <ος)—0.0988Conngos keV. | This temperature is $<T_{\it ROSAT}> = 0.0988_{-0.046}^{+0.268} \keV$ . |
Given the true temperature clistribution in Fig. | Given the true temperature distribution in Fig. |
5 the emission-weighted average temperature is more a rellection of the instrument response than a fair estimation of the temperature of the emitted radiation or the general state of the bubble. | \ref{fig:em} the emission-weighted average temperature is more a reflection of the instrument response than a fair estimation of the temperature of the emitted radiation or the general state of the bubble. |
<μονds instrument specific. defeating our object of trving to produce an instrument-indepencdent: estimator for the temperature. | $<T_{\it ROSAT}>$ is instrument specific, defeating our object of trying to produce an instrument-independent estimator for the temperature. |
This may. well be impossible. in such a situation as this where the temperature distribution can not be described by a single. characteristic temperature. | This may well be impossible, in such a situation as this where the temperature distribution can not be described by a single, characteristic temperature. |
The emission measure should. follow from the normalisation of the spectrum. | The emission measure should follow from the normalisation of the spectrum. |
ancl will depend. on the best-fit values for the other parameters which determine the shape of the model. spectrum. | and will depend on the best-fit values for the other parameters which determine the shape of the model spectrum. |
For the purposes of comparison between the results of the spectral fitting below and the “true” values. the total emission measure between 0.01(κο)<10.0 is 3.107cm | For the purposes of comparison between the results of the spectral fitting below and the “true” values, the total emission measure between $0.01 \leq T
(\keV) \leq 10.0$ is $3.77 \times 10^{58}
\pcc$ . |
)etween TkeV)€ 24. itis S64.10em. | Between $0.1 \leq T (\keV) \leq 2.4$ , it is $8.64 \times
10^{56} \pcc$ . |
For an assumed exposure time of 3000s we obtain a simulated.ROSAL PSPC spectrum containing —1000 counts. | For an assumed exposure time of $3000 \s$ we obtain a simulated PSPC spectrum containing $\sim 1000$ counts. |
To assess the elfect of real photon statistics. aud he consequent variation in best-fit parameters. on our simulated: spectra we generated ten. Poisson. realizations. itting each using theSTARLINK. X-ray analysis. packageASTERIX. | To assess the effect of real photon statistics, and the consequent variation in best-fit parameters, on our simulated spectra we generated ten Poisson realizations, fitting each using the X-ray analysis package. |
The number of counts per bin is sullicient to allow he use of 47 fitting. | The number of counts per bin is sufficient to allow the use of $\chi^{2}$ fitting. |
Three iferent spectral. models were considered: a stanclarcl single temperature absorbed hot plasma mocdel: a wo temperature mocol. with both components having the same absorbing column and metal abundance. ancl finally a differential emission measure moclel. | Three different spectral models were considered: a standard single temperature absorbed hot plasma model; a two temperature model, with both components having the same absorbing column and metal abundance, and finally a differential emission measure model. |
Single temperature models are widely usec to characterise N-ray emission. and it is only when there are sullicient counts to show that a single temperature mocel is a bad fit that more complex mocels are used. | Single temperature models are widely used to characterise X-ray emission, and it is only when there are sufficient counts to show that a single temperature model is a bad fit that more complex models are used. |
I is therefor sensible to fit such a model. despite knowing that the true temperature distribution of the X-ray emitting plasma is far [rom being single temperature (see Fig. 5)). | It is therefor sensible to fit such a model, despite knowing that the true temperature distribution of the X-ray emitting plasma is far from being single temperature (see Fig. \ref{fig:em}) ). |
A two temperature model would be the next. level of complexity. naively a soft. component for cooler denser gas near the shell ancl a hot component for the bubble interior. | A two temperature model would be the next level of complexity, naively a soft component for cooler denser gas near the shell and a hot component for the bubble interior. |
The cillerential emission model attempts to incorporate he emission. of gas at a wider range of temperatures w specibving the emission. measure AL=J24V is à power law in temperature. AMx17. between two emperatures Z1, and Lian. | The differential emission model attempts to incorporate the emission of gas at a wider range of temperatures by specifying the emission measure $EM = \int n_{\rm e}^{2} dV$ is a power law in temperature, $EM
\propto T^{\gamma}$, between two temperatures $T_{\rm low}$ and $T_{\rm high}$. |
We quote results for models with all parameters fitted or (including the column and the metallicity) and for fits with the metallicity fixed at solar abundance. | We quote results for models with all parameters fitted for (including the column and the metallicity) and for fits with the metallicity fixed at solar abundance. |
Although in oactice we know the absorbing column applied. fits with it fixed were statistically unacceptable. | Although in practice we know the absorbing column applied, fits with it fixed were statistically unacceptable. |
For the purpose of display ancl interpretation. the best-fit results for each of the en Poisson realisations were averaged. | For the purpose of display and interpretation, the best-fit results for each of the ten Poisson realisations were averaged. |
Phe quoted: errors arethestatistical deviations of the best-fit results.and. not he fitted confidence regions. | The quoted errors arethestatistical deviations of the best-fit results,and not the fitted confidence regions. |
Typically. the confidence | Typically, the confidence |
Veilleux. Kim Sanders 1999). | Veilleux, Kim Sanders 1999). |
In particular. the second requirement removes any ambiguity in the AGN identilication. since the nuclear component is alwavs strong enough to dominate the 58 jm emission. | In particular, the second requirement removes any ambiguity in the AGN identification, since the nuclear component is always strong enough to dominate the 5–8 $\mu$ m emission. |
The corresponding selection consists of six objects for which N-rav observations are already available. even if only in four cases the results have been published. | The corresponding selection consists of six objects for which X-ray observations are already available, even if only in four cases the results have been published. |
Four additional sources. i.e. IRAS 00397— 1312. IRAS 01003—2238. IRAS 01298—0744 and IRAS 12127-1412. all meeting the latter (wo criteria. have been obtained as à part of our X-ray [ollow-up campaign. | Four additional sources, i.e. IRAS $-$ 1312, IRAS $-$ 2238, IRAS $-$ 0744 and IRAS $-$ 1412, all meeting the latter two criteria, have been obtained as a part of our X-ray follow-up campaign. |
They were chosen to complete the present sample. and lor (heir remarkable expected X-rav. properties: By assuming a standard quasar spectral energy distribution (SED: Elvis et al. | They were chosen to complete the present sample, and for their remarkable expected X-ray properties: By assuming a standard quasar spectral energy distribution (SED; Elvis et al. |
1994). suggesting that v£,,(2500 Aj) 2 2xvL,(6 jm). and the most recent relations between the UV and X-ray huminosities (Lusso οἱ al. | 1994), suggesting that $\nu L_\nu$ (2500 ) $\simeq$ $\times \nu L_\nu$ (6 $\mu$ m), and the most recent relations between the UV and X-ray luminosities (Lusso et al. |
2010). the intrinsic 2LO keV flix of the AGN component in IRAS 01003—2238 and IRAS 01298—0744 is expected to be ~7x10.P erg Lom 7. ensuring a prominent detection in case of a Compton-thin column density Da7). | 2010), the intrinsic 2–10 keV flux of the AGN component in IRAS $-$ 2238 and IRAS $-$ 0744 is expected to be $\sim 7 \times 10^{-12}$ erg $^{-1}$ $^{-2}$, ensuring a prominent detection in case of a Compton-thin column density $N_\rmn{H} \la 5 \times 10^{23}$ $^{-2}$ ). |
Even in a Compton-thick prospect. allowing [or a reflection efficiency. of a. few per cent makes the AGN clearly detectable with the current. X-ray observatories. | Even in a Compton-thick prospect, allowing for a reflection efficiency of a few per cent makes the AGN clearly detectable with the current X-ray observatories. |
IRAS 00397-1312 and IRAS 12127-1412 are very interesting in this context as well (see also Imanishi et al. | IRAS $-$ 1312 and IRAS $-$ 1412 are very interesting in this context as well (see also Imanishi et al. |
2008. 2010). | 2008, 2010). |
The former 15 optically classified as II region in spite of being the most luminous ULIRG in the 1 Jv sample (Ixim Sanders 1993). with LycLOMLs: its mid-IR. spectrum shows faint PAIL features (Fig. 1)). | The former is optically classified as H region in spite of being the most luminous ULIRG in the 1 Jy sample (Kim Sanders 1998), with $L_\rmn{IR} \simeq 10^{13} L_{\sun}$: its mid-IR spectrum shows faint PAH features (Fig. \ref{f1}) ), |
and clearly hints at ihe simultaneous presence of a buried AGN roughly comparable to the SB in terms of its Ih energv output. | and clearly hints at the simultaneous presence of a buried AGN roughly comparable to the SB in terms of its IR energy output. |
The latter is found instead near the lower end of the ULIBG luminosity range. but its heavily absorbed continuum lacking anv star formation signature implies an AGN-dominated nature. | The latter is found instead near the lower end of the ULIRG luminosity range, but its heavily absorbed continuum lacking any star formation signature implies an AGN-dominated nature. |
For these two sources the predicted X-ray. {Ins is just slightly less than the estimate given above (by a factor of 22). | For these two sources the predicted X-ray flux is just slightly less than the estimate given above (by a factor of $\sim$ 2). |
The four ~214 jan resi-Irame spectra are filly representative of this ULIRG subclass. and are shown in Fie. | The four $\sim$ 2–14 $\mu$ m rest-frame spectra are fully representative of this ULIRG subclass, and are shown in Fig. |
1 in contrast with a typical SB template. while the general information concerning all the sources and (heir physical parameters as derived [rom our mid-IR analvsis are listed in Table 1.. | \ref{f1} in contrast with a typical SB template, while the general information concerning all the sources and their physical parameters as derived from our mid-IR analysis are listed in Table \ref{t1}. |
Considering all (he entries. (his X-ray sample allows us to probe a wide range of properties with respect to obscured activity inside ULIRGs. | Considering all the entries, this X-ray sample allows us to probe a wide range of properties with respect to obscured activity inside ULIRGs. |
It is also worth | It is also worth |
the radio source is AR/30degs=3«LOMeins. | the radio source is $R/30 {\mbox days} = 3\times 10^{10}\; cm\; s^{-1}$. |
| So we soo directly that CRB 970508 expanded at au average speed of e over a whole month. giving us a direct observational proof that the source is highly relativistic. | So we see directly that GRB 970508 expanded at an average speed of $c$ over a whole month, giving us a direct observational proof that the source is highly relativistic. |
This proof is completely equivalent to superlumninual motions iu blazars. aud is the strougest evidence in favor of the fireball mioclel. | This proof is completely equivalent to superluminal motions in blazars, and is the strongest evidence in favor of the fireball model. |
The afterglow of GRB 970508 is our best case so far: it is in fact à burst for which rot only do we know the redshift. but also a radio source that las been imonitorec or more than 100 days after the explosion (Frail. Waxiuan and IKulkunui 2000). | The afterglow of GRB 970508 is our best case so far: it is in fact a burst for which not only do we know the redshift, but also a radio source that has been monitored for more than $400$ days after the explosion (Frail, Waxman and Kulkarni 2000). |
Through these observations we can see the transition to a subrelativistic regine at fcx100d. measure the total cuerectics of the following Sedov phase (uuencunmnboere« w relativistic effects!) | Through these observations we can see the transition to a sub–relativistic regime at $t \approx 100 \; d$, measure the total energetics of the following Sedov phase (unencumbered by relativistic effects!) |
Ei=5s1079ergs. determine two elusive paranieters. Gy=0.5 and eg=0.5 (the efficiencies with which cucrey is traustered to electrons by protons. and with which an equipartition field is built up). aux he density of the smmrounding medimu »z0.1emο, | $E_{New} = 5\times 10^{50}\; ergs$ determine two elusive parameters, $\epsilon_{eq}
= 0.5$ and $\epsilon_B= 0.5$ (the efficiencies with which energy is transfered to post--shock electrons by protons, and with which an equipartition field is built up), and the density of the surrounding medium $n \approx 0.4\; cm^{-3}$. |
ΑΙ of these values look reasonable (perhaps e4, aud ερ exceed our expectations by a factor of LO. a fact hat could be remedied by introducing a slight deusity eradieut which would keep he shock more efficient). so that our confidence in the externalshockintheISM nodel is boosted. | All of these values look reasonable (perhaps $\epsilon_{eq}$ and $\epsilon_B$ exceed our expectations by a factor of 10, a fact that could be remedied by introducing a slight density gradient which would keep the shock more efficient), so that our confidence in the external–shock–in–the–ISM model is boosted. |
Another precious consequence of these latetine observatious is that thev vield information on beaming aud energetics. | Another precious consequence of these late–time observations is that they yield information on beaming and energetics. |
In fact. CRB 970508 appeared to have a dnetie energv of Ey.=5«10°!erg when in the relativistic phase; Επment which can be recouciled with Ly. (vemember that the expausion is adiabatic. sohat we must have Ey,= Ei!) | In fact, GRB 970508 appeared to have a kinetic energy of $E_{rel} = 5\times 10^{51}\; erg$ when in the relativistic phase, a measurement which can be reconciled with $E_{New}$ (remember that the expansion is adiabatic, sothat we must have $E_{New} = E_{rel}$ !) |
oulv if the unknown beaming anele. assumed = [rin deriving £4.15 smaller than la by the factor Ly. ρω we thus have the | only if the unknown beaming angle, assumed $= 4\pi$ in deriving $E_{rel}$ ,is smaller than $4\pi$ by the factor $E_{New}/E_{rel}$ ; we thus have the |
Supernova remnants (SNRs) are widely thought to be one important. kind of cosmic rav (CR) sources in the Galaxy CAharonianetal.2004).. | Supernova remnants (SNRs) are widely thought to be one important kind of cosmic ray (CR) sources in the Galaxy \citep{2004Natur.432...75A}. |
The most direct evidence comes from high-energv 5-rav emission from SNRs. | The most direct evidence comes from high-energy $\gamma$ -ray emission from SNRs. |
Generally there are (wo types of scenarios for production of high energy 5-ravs: the leptonic (via inverse Compton scattering of background photos bv relativistic electrons) and hadronie (via decay of neutral pious produced by elastic collisions of relativistic ious with ions in the background plasma) origins. | Generally there are two types of scenarios for production of high energy $\gamma$ -rays: the leptonic (via inverse Compton scattering of background photos by relativistic electrons) and hadronic (via decay of neutral pions produced by elastic collisions of relativistic ions with ions in the background plasma) origins. |
Understanding which of these (svo scenarios is dominant in specifie sources is verv important for the search ol CR nuclei sources and the study of CR acceleration (Gabici2003). | Understanding which of these two scenarios is dominant in specific sources is very important for the search of CR nuclei sources and the study of CR acceleration \citep{Gabici2008}. |
. Usually it is dilliealt to disünguish the leptonic model and hadronic model just with the high energv y-ray data alone. | Usually it is difficult to distinguish the leptonic model and hadronic model just with the high energy $\gamma$ -ray data alone. |
Multi-wavelength. observations of photon emission [rom SNRs can provide us kev inlormation about (he radiation mechanism. | Multi-wavelength observations of photon emission from SNRs can provide us key information about the radiation mechanism. |
Subsets and Splits
No community queries yet
The top public SQL queries from the community will appear here once available.