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Thus. he stellar cluster which powers he III reglon IS uostlv hidden iu the optical aud. near-IR .bu scattered Πα eimissiou marks the blistering edge of tie. HIT region. | Thus, the stellar cluster which powers the HII region is mostly hidden in the optical and near-IR, but scattered $\alpha$ emission marks the blistering edge of the HII region. |
The dense core of the nolecular cloud is offset from the WIT region aud its embedded. cluster by ~2’. and a tlin CTI region delineates the boundary between the cool uolecular gas aud the hot ionized region. | The dense core of the molecular cloud is offset from the HII region and its embedded cluster by $\sim
2\arcmin$, and a thin CII region delineates the boundary between the cool molecular gas and the hot ionized region. |
The distance to W 1035 not vet well coustriined. | The distance to W40 is not yet well constrained. |
Severa measurements of atonuüc and molecular ines along the line of sight owards WO have oen used tfo estimae the distance to WLO assuming that the cloud complex is in a circular orbit around the galactic center (777).. | Several measurements of atomic and molecular lines along the line of sight towards W40 have been used to estimate the distance to W40 $-$ assuming that the cloud complex is in a circular orbit around the galactic center \citep{rei1970,dow1970,rad1972}. |
Taken ogether. these measurements produce only a weak constrain of 300-900 pc. | Taken together, these measurements produce only a weak constraint of 300-900 pc. |
Using assumptions about he spectral types of several of he stars in the W10 cluster. distance nioulus estimates can ouly luat he distauce to a wie range frou 100 to TOO pe (?7).. | Using assumptions about the spectral types of several of the stars in the W40 cluster, distance modulus estimates can only limit the distance to a wide range from 400 to 700 pc \citep{cru1982,smi1985}. |
2 calculated a distance of 600 pc using a rovel disance determination techuique mvolvius OIL line LOASTALCICLs. | \citet{kol1983}
calculated a distance of 600 pc using a novel distance determination technique involving OH line measurements. |
Hereinafter we will adopt a distance of 600 pc. which puts the W10 complex at a height of 37 pe above the ealactic plane. | Hereinafter we will adopt a distance of 600 pc, which puts the W40 complex at a height of 37 pc above the galactic plane. |
More recently. the properties of the WLO region have been stuumarized iua review by ?.. | More recently, the properties of the W40 region have been summarized in a review by \citet{rod2008}. |
Also. ? have performed a deep X-ray stdv of W10. and detected about 200 voung scllay iieiubers of the cluster. | Also, \citet{kuh2010}
have performed a deep X-ray study of W40, and detected about 200 young stellar members of the cluster. |
We have used the Very. Large Array (VLA) of the NRAO in its Band A array configurations to observe WI0 at 3.6 cua. under the NRAO program. code AR551. | We have used the Very Large Array (VLA) of the NRAO in its B and A array configurations to observe W40 at 3.6 cm, under the NRAO program code AR551. |
The D configuration observations were made on 200:3 November 3 wdüle the A configuration ooervatious were nude ou 200 September 15. | The B configuration observations were made on 2003 November 3 while the A configuration observations were made on 2004 September 18. |
Our phase center was at a(2000)=q1531"2:57: 6.72000)=02705/29", | Our phase center was at $\alpha (J2000) = 18^h 31^m
23\fs7$ ; $\delta (J2000) = -02\degr 05\arcmin 29\arcsec$. |
The amplitude calibrator was 133]|305. with an adopted fiux deusitv of 5.18 Jv. and the phase caliator Was Lsol|010, with flux densities of 0.76+0.0] alc O.82+£0.01 Jv for he first aud secouc epochs of observation. respecively, | The amplitude calibrator was $1331+305$, with an adopted flux density of 5.18 Jy, and the phase calibrator was $1804+010$, with flux densities of $0.76\pm 0.01$ and $0.82\pm 0.01$ Jy for the first and second epochs of observation, respectively. |
The half power contour of the svuthesized lenns were O/SS«ΠΠ1137 (for configuration B data) aud 07269/22:|16? (for configuration A data). for images made with he ROBUST parameter of the task INLACR set o 0 (BrigesOO 1995). | The half power contour of the synthesized beams were $0\rlap.{''}88 \times 0\rlap.{''}76; +17^\circ$ (for configuration B data) and $0\rlap.{''}26 \times 0\rlap.{''}22;
+16^\circ$ (for configuration A data), for images made with the ROBUST parameter of the task IMAGR set to 0 (Briggs 1995). |
The data were analyzed in he standard manner using the NRAO AIPS )ackage. | The data were analyzed in the standard manner using the NRAO AIPS package. |
The data for the two epochs were self-calibrated iu phase separately aud then concatenated to obtain a final inage with the ghest scusitivity possible. | The data for the two epochs were self-calibrated in phase separately and then concatenated to obtain a final image with the highest sensitivity possible. |
The half power coutour of the svuthesized beam or this concatenated «ata was 0736&0730:|287 Or miages nade with he ROBUST parameter of he task INLACT set to 0 (7).. | The half power contour of the synthesized beam for this concatenated data was $0\rlap.{''}36 \times 0\rlap.{''}30; +28^\circ$ for images made with the ROBUST parameter of the task IMAGR set to 0 \citep{bri1995}. |
A total of 20 sources were detected in a field of I Vin this image. | A total of 20 sources were detected in a field of $4\arcmin \times
4\arcmin$ in this image. |
Following ?.. we estimate tlat in a field of VsU thepriori number of expeced 3.6 cur sources above ~0.2 12Jv. is only ~0.3. | Following \citet{fom2002}, we estimate that in a field of $4\arcmin \times 4\arcmin$ the number of expected 3.6 cm sources above $\sim$ 0.2 mJy is only $\sim$ 0.3. |
We then conclude. that perlaps one out of the 20 sources could be a vackeround object. but that we are jusied imm assuming that all the detected SOTIECOS are associated with WI0. | We then conclude that perhaps one out of the 20 sources could be a background object, but that we are justified in assuming that all the detected sources are associated with W40. |
Iufrared plotomery was carried out with the QUick Tufrared Canera (QUIRC. 73) on the University of Iba 2.2 meter telescope from 2003 June lr to 19. | Infrared photometry was carried out with the QUick Infrared Camera (QUIRC, \citealt{hod1996}) ) on the University of Hawaii 2.2 meter telescope from 2003 June 17 to 19. |
Data were obtained in the J. TE. and fy! filters of the Mauna Ikea Observatories (MIO) filter set (2)... | Data were obtained in the $J$, $H$, and $K^\prime$ filters of the Mauna Kea Observatories (MKO) filter set \citep{tok2005}. |
These IR nuages were reduced and combined usus the Comiui Observatory* QUIRC package in IRAF. | These IR images were reduced and combined using the Gemini Observatory's QUIRC package in IRAF. |
Our data were colleτσι, under uou-phooietric conditions. aud errrs in the telescope euidiug combined with persistent focus problems ac edt iupossble to pursue photometry by fitting the oot spread function. as is oxeferred. for crowded fields. | Our data were collected under non-photometric conditions, and errors in the telescope guiding combined with persistent focus problems made it impossible to pursue photometry by fitting the point spread function, as is preferred for crowded fields. |
Thus. aperture photometry was performed using the NOÀO DIGIPIIOT package. | Thus, aperture photometry was performed using the NOAO DIGIPHOT package. |
The our xiehtest IB stars were saurated even n our shortest (1 sec) exposures. so without PSF fitlue we were onlv able to obtain upper liudts ou the uaeguitudes. which aeree well with previous work. | The four brightest IR stars were saturated even in our shortest (1 sec) exposures, so without PSF fitting we were only able to obtain upper limits on the magnitudes, which agree well with previous work. |
We fud over 1000 IR sources with detections | We find over 1000 IR sources with detections |
(Marietta et al. 2000)). | (Marietta et al. \cite{MAR00}) ). |
This may be a channel to forming single low-mass white dwarf (Justham et al. 2009)). | This may be a channel to forming single low-mass white dwarf (Justham et al. \cite{JUSTHAM09}) ). |
However. there are only a few systems with M,«0.02M.. in our sample. which means that the binary sequence considered in this paper may only be able to explain some individual observations (Mattila et al. 2005:: | However, there are only a few systems with $M_{\rm
e}<0.02M_{\odot}$ in our sample, which means that the binary sequence considered in this paper may only be able to explain some individual observations (Mattila et al. \cite{MAT05}; |
Leonard 2007)). | Leonard \cite{LEO07}) ). |
We have found that if a SN Ia originates from the WD + RG channel. the donor star may almost become à helium WD and only some hydrogen-rich material remains on top of the helium WD (as little as 0.017 M). which means that the upper limit mass to the stripped-off from the companion by supernova ejecta Is 0.017 M... | We have found that if a SN Ia originates from the WD + RG channel, the donor star may almost become a helium WD and only some hydrogen-rich material remains on top of the helium WD (as little as 0.017 $M_{\odot}$ ), which means that the upper limit mass to the stripped-off from the companion by supernova ejecta is 0.017 $M_{\odot}$. |
No hydrogen line is then expected in the nebular spectra of some SNe Ia and our results may explain the conflict between theory and observation. re. theory predicts that the stripped-off material should be greater than 0.035 M. (Marietta et al. 2000:; | No hydrogen line is then expected in the nebular spectra of some SNe Ia and our results may explain the conflict between theory and observation, i.e, theory predicts that the stripped-off material should be greater than 0.035 $M_{\odot}$ (Marietta et al. \cite{MAR00}; |
Meng. Chen Han 2007:: Meng Yang 2010b)). while observations indicate that the upper mass limit of the stripped-off material 1s 0.02 M. (Mattila et al. 2005:: | Meng, Chen Han \cite{MENGXC07}; Meng Yang \cite{MENGXC09b}) ), while observations indicate that the upper mass limit of the stripped-off material is 0.02 $M_{\odot}$ (Mattila et al. \cite{MAT05}; |
Leonard 2007)). | Leonard \cite{LEO07}) ). |
No hydrogen line has been detected in nebular spectra of some SNe Ia may indicate that the progenitors of the observed SNe Ia are from WD + RG systems. | No hydrogen line has been detected in nebular spectra of some SNe Ia may indicate that the progenitors of the observed SNe Ia are from WD + RG systems. |
If the observed SNe Ia (SN 2001el. 2005am. and 2009561) were produced by the WD + RG channel. they should originate in an old population. | If the observed SNe Ia (SN 2001el, 2005am, and 2005cl) were produced by the WD + RG channel, they should originate in an old population. |
Unfortunately. there is no constraint on the age of the three SNe la. We checked the types of their host galaxies and found that apart from the host galaxy of SN 2005cl (MCG-01-39-003). which is a SO galaxy (Wang 2009:: Bufano et al. 2009)). | Unfortunately, there is no constraint on the age of the three SNe Ia. We checked the types of their host galaxies and found that apart from the host galaxy of SN 2005cl (MCG-01-39-003), which is a S0 galaxy (Wang \cite{WANGXF09}; Bufano et al. \cite{BUFANO09}) ), |
they are both spiral galaxies. re. the host galaxy of SN 2001el (NGC 1448) is a Sed galaxy (Wang et al. 2003:: | they are both spiral galaxies, i.e. the host galaxy of SN 2001el (NGC 1448) is a Scd galaxy (Wang et al. \cite{WANGLF03}; |
Wang et al. 2006)) | Wang et al. \cite{WANGXF06}) ) |
and the host galaxy of SN 2001am (NGC 2811) is a Sa galaxy (Bufano et al. 2009)). | and the host galaxy of SN 2001am (NGC 2811) is a Sa galaxy (Bufano et al. \cite{BUFANO09}) ). |
The progenitor of SN 2005cl may therefore belong to an old population. but we are unable to infer any information about the population of the other SNe Ia. However. we note that all three SNe Ia are located at the edge of their host galaxy and that SN 2005cl is even located in a tail extending from MCG-01-39-003. | The progenitor of SN 2005cl may therefore belong to an old population, but we are unable to infer any information about the population of the other SNe Ia. However, we note that all three SNe Ia are located at the edge of their host galaxy and that SN 2005cl is even located in a tail extending from MCG-01-39-003. |
Is this phenomenon evidence of an old population? | Is this phenomenon evidence of an old population? |
It is possible because halo stars in general belong to an old population. | It is possible because halo stars in general belong to an old population. |
All three SNe Ia are located at the edge of their host galaxy may be an observational select effect because we are more likely to observe a SN Ia at the outskirts of a host galaxy rather than its inner part. | All three SNe Ia are located at the edge of their host galaxy may be an observational select effect because we are more likely to observe a SN Ia at the outskirts of a host galaxy rather than its inner part. |
This selection effect might increase the probability that à SN Ia is produced by the WD +RG channel with a low-mass envelope being observed. | This selection effect might increase the probability that a SN Ia is produced by the WD +RG channel with a low-mass envelope being observed. |
Relative to that of the WD + MS system. the Galactic birth rate of the WD + RG channel is low (see Meng Yang 2010b and Wang. Li Han (2010))). | Relative to that of the WD + MS system, the Galactic birth rate of the WD + RG channel is low (see Meng Yang \cite{MENGXC09b} and Wang, Li Han \cite{WANGB09c}) )). |
However. since the Galactic birth rate of SNe Ia predicted by the model in Meng Yang 2010b is lower than that inferred from observations. the WD + RG channel should be carefully investigated because the progenitors of some SNe Ia (e.g. SN 2006X and SN 200701) are possible WD + RG systems (Patat et al. 2007a:: | However, since the Galactic birth rate of SNe Ia predicted by the model in Meng Yang \cite{MENGXC09b} is lower than that inferred from observations, the WD + RG channel should be carefully investigated because the progenitors of some SNe Ia (e.g. SN 2006X and SN 2007on) are possible WD + RG systems (Patat et al. \cite{PATAT07}; ; |
Voss Nelemans 2008). | Voss Nelemans \cite{VOSS08}) ). |
In addition. some recurrent. nova (belonging to WD + RG) are suggested to be the candidates of SNe Ia progenitors (Hachisu et al. 1999b:; | In addition, some recurrent nova (belonging to WD + RG) are suggested to be the candidates of SNe Ia progenitors (Hachisu et al. \cite{HAC99b}; |
Hachisu Kato 2006:; Hachisu et al. 2007)). | Hachisu Kato \cite{HK06}; Hachisu et al. \cite{HKL07}) ). |
In this paper. we even proposed that the prevalence of the WD + RG channel is why no hydrogen line was detected in nebular spectra of some SNe Ia. although the probability of its occurrence is low. | In this paper, we even proposed that the prevalence of the WD + RG channel is why no hydrogen line was detected in nebular spectra of some SNe Ia, although the probability of its occurrence is low. |
Marietta et al. (2000)) | Marietta et al. \cite{MAR00}) ) |
performed several high-resolution two-dimensional numerical simulations of the collision between the supernova ejecta and companion. and found that a red-giant donor will lose almost its entire envelope - 985t)) due to the impact leaving only the core of the star 0.432M..). | performed several high-resolution two-dimensional numerical simulations of the collision between the supernova ejecta and companion, and found that a red-giant donor will lose almost its entire envelope - ) due to the impact leaving only the core of the star $\simeq0.42M_{\odot}$ ). |
The RG star used in Marietta et al. (2000)) | The RG star used in Marietta et al. \cite{MAR00}) ) |
consists of a helium core of 0.432M... and an envelope of 0.56M... which are not comparable to those obtained by our simulations (see Figs 3 and 4)). | consists of a helium core of $0.42 M_{\odot}$ and an envelope of $0.56
M_{\odot}$, which are not comparable to those obtained by our simulations (see Figs \ref{Fig3} and \ref{Fig4}) ). |
In addition. the radius of their RG model is 180 Ra. which corresponds to an orbital period of ~900 days. | In addition, the radius of their RG model is 180 $R_{\odot}$, which corresponds to an orbital period of $\sim900$ days. |
Too long to compare with our simulation (see Fig. 1)). | Too long to compare with our simulation (see Fig. \ref{Fig1}) ), |
the orbital period leads to a lower envelope binding energy than produced by the model developed in this paper since the binding energy of the envelope is determined mainly by the radius of the RG star (Meng. Chen Han 2008)). | the orbital period leads to a lower envelope binding energy than produced by the model developed in this paper since the binding energy of the envelope is determined mainly by the radius of the RG star (Meng, Chen Han \cite{MENG08}) ). |
The envelope of the RG nodel used by Marietta et al. (2000)) | The envelope of the RG model used by Marietta et al. \cite{MAR00}) ) |
is then more likely to be stripped off and the amount of material stripped-off by the RG companion in Marietta et al. (2000)) | is then more likely to be stripped off and the amount of material stripped-off by the RG companion in Marietta et al. \cite{MAR00}) ) |
might be overestimated. | might be overestimated. |
A nore detailed numerical simulation of the interaction between supernova ejecta and an RG companion should therefore be performed by a more physical companion model than that in Aarietta et al. (2000). | A more detailed numerical simulation of the interaction between supernova ejecta and an RG companion should therefore be performed by a more physical companion model than that in Marietta et al. \cite{MAR00}) ). |
The absence of hydrogen lines in the nebular spectra of SNe la may have other explanations. | The absence of hydrogen lines in the nebular spectra of SNe Ia may have other explanations. |
The RG companion with a small hydrogen-rich envelope may be the results of either a fine-tuning effect as suggested in this paper. or a physical process. | The RG companion with a small hydrogen-rich envelope may be the results of either a fine-tuning effect as suggested in this paper, or a physical process. |
For example. Justham et al. (2009)) | For example, Justham et al. \cite{JUSTHAM09}) ) |
suggested that a rotational effect of WD could prevent the thermonuclear runaway occurring until the accretion phase has ended. which could also produce a RG companion with a low-mass envelope. | suggested that a rotational effect of WD could prevent the thermonuclear runaway occurring until the accretion phase has ended, which could also produce a RG companion with a low-mass envelope. |
Rotation may also increase the probability of a SN Ia from WD +RGchannel with low-mass envelope being observed. | Rotation may also increase the probability of a SN Ia from WD +RGchannel with low-mass envelope being observed. |
An alternative explanation of the lack of hydrogen is that the | An alternative explanation of the lack of hydrogen is that the |
we will discuss in the next section. this is not be the case at TOUy€. | we will discuss in the next section, this is not be the case at $z<3$. |
The connection between the four phases of cosmic gas is important in galaxy formation. and it is beneficial for us to study. their evolution for a better understanding of metal cooling elfects. | The connection between the four phases of cosmic gas is important in galaxy formation, and it is beneficial for us to study their evolution for a better understanding of metal cooling effects. |
A conventional view is that the clilfuse gas is shock-heated to hot or warm-hot. phase. which later cools down to the condensed. phase (Davectal.1999:Cen&Ostriker1999:Daveetal. 2001). | A conventional view is that the diffuse gas is shock-heated to hot or warm-hot phase, which later cools down to the condensed phase \citep{Dave.etal:99,
Cen.Ostriker:99, Dave.etal:01}. |
. Ixereset.al.(2005) also suggested that some cilluse gas can migrate to the condensed. phase without being shock-heated. which they called the "cold mode accretion. | \citet{Keres.etal:05} also suggested that some diffuse gas can migrate to the condensed phase without being shock-heated, which they called the `cold mode' accretion. |
Figure 5.r shows the evolution of mass [fractions of | Figure \ref{fig:Evolv4} shows the evolution of mass fractions of |
is "based ou the SAN J2103.5|1515 detection. | is $^{\prime}$ based on the SAX J2103.5+4545 detection. |
For a strong source like Cve X.1. we find au offset of 0. | For a strong source like Cyg X–1, we find an offset of $^{\prime}$. |
1 As expected. when the source is far away from the pointing direction the PSLA is worse: in fact (νο N Lat an off-axis anele of 11.67 was measured with au offset of 3.0% | As expected, when the source is far away from the pointing direction the PSLA is worse; in fact Cyg X–1 at an off-axis angle of $^\circ$ was measured with an offset of $^{\prime}$. |
The statistical errors of the PSLA are sieuificautlv larger than the heoretical values aud clearly dominated by systematics (τοςetal. 2003)). nuuünlv. due to the nmisalieumieut rotation matrix computed at ISDC (Walteretal. 20033) in order to correct the offset between the telescope axis and the axis of the star sensors. | The statistical errors of the PSLA are significantly larger than the theoretical values and clearly dominated by systematics \cite{gros03}) ), mainly due to the misalignment rotation matrix computed at ISDC \cite{walter03}) ) in order to correct the offset between the telescope axis and the axis of the star sensors. |
Finally. during the CPSs of revolutious 25 aud 26. —BISASCRI detected Cve X.1 at a fiux level of roughly 100 for energies 20.10keV. | Finally, during the GPSs of revolutions 25 and 26, IBIS/ISGRI detected Cyg X–1 at a flux level of roughly 400 for energies 20–40. |
The flux was variable. as also reported in Bazzanoctal.2003). | The flux was variable, as also reported in \cite{bazzano03}. |
. Furthermore frou. Fig. | Furthermore from Fig. |
3. the black hole candidate exhibits a harder specrum than the Crab below 50 aud softer above roughly 150keV. which is a coufirmation of the known spectral shape of Cve X1 in its lut state. | \ref{fig:spec} the black hole candidate exhibits a harder spectrum than the Crab below 50 and softer above roughly 150, which is a confirmation of the known spectral shape of Cyg X–1 in its hard state. |
(Sari. Piran. Halpern 1999) and thus the beaming-corrected. gamma-ray energy Ej,=1.3x10?"/O4dο./10eres)* (n/O.Lem"tp/0.5) eres. which is by a factor of a few smaller than the mean energy release [ound by Frail et al. ( | (Sari, Piran, Halpern 1999) and thus the beaming-corrected gamma-ray energy $E_{jet}=1.3\times 10^{50}(t_j/0.4\,{\rm
d})^{3/4}(E_{iso}/10^{53}{\rm ergs})^{3/4}(n/0.1\,{\rm
cm}^{-3})^{1/4}(\eta_\gamma/0.5)^{1/4}$ ergs, which is by a factor of a few smaller than the mean energy release found by Frail et al. ( |
2001) for n~O.Lem and og,e0.5. | 2001) for $n\sim
0.1\,{\rm cm}^{-3}$ and $\eta_\gamma\sim 0.5$. |
Even so. this energy. still satisfies the relation of Ghirlanda et al. ( | Even so, this energy still satisfies the relation of Ghirlanda et al. ( |
2004). ES,/A0eyes)=(1.122:0.12)JIA[(1+2),J4+p/ἤ100keV] (Κου Dai. Liang. Xu 2004). | 2004), $(E_{jet}/10^{50}{\rm
ergs})=(1.12\pm 0.12)[(1+z)E_p/100\,{\rm keV}]^{1.50\pm 0.08}$ (see Dai, Liang, Xu 2004). |
Finally. what we want to point out is that besides (he reverse-shock model cliscussecl in (his paper. the other plausible explanations of a rebrightening light curve include the densitv-jump medium (Dai Lu 2002: Lazzati et al. | Finally, what we want to point out is that besides the reverse-shock model discussed in this paper, the other plausible explanations of a rebrightening light curve include the density-jump medium (Dai Lu 2002; Lazzati et al. |
2002). pure Povntine-fhis injection (Dai Lu 1993). barvon-dominated injection (Rees Meésszárros 1998: Granot. Nakar Piran 2003: Nakar. Piran Granot 2003: Djórrnsson et al. | 2002), pure Poynting-flux injection (Dai Lu 1998), baryon-dominated injection (Rees Mésszárros 1998; Granot, Nakar Piran 2003; Nakar, Piran Granot 2003; Björrnsson et al. |
2004). (vo-component jet (Berger et al. | 2004), two-component jet (Berger et al. |
2003: Luang et al. | 2003; Huang et al. |
2003). neutron-fed fireball (Deloborodov 2003). temporal [Huctuation of e, and/or ej and departure of the electron distribution from a power law (DPanaitescu Kumar 2004). | 2003), neutron-fed fireball (Beloborodov 2003), temporal fluctuation of $\epsilon_e$ and/or $\epsilon_B$ and departure of the electron distribution from a power law (Panaitescu Kumar 2004). |
Multiwavelength: observations in the era are expected (o distinguish among these possibilities. | Multiwavelength observations in the era are expected to distinguish among these possibilities. |
We thank Y. Z. Fan and L. J. Gou for constructive comments. X. F. Wu and Y. C. Zou for helpful discussions. and Y. EF. Huang and X. Y. Wang for a careful reading of our manuscript. | We thank Y. Z. Fan and L. J. Gou for constructive comments, X. F. Wu and Y. C. Zou for helpful discussions, and Y. F. Huang and X. Y. Wang for a careful reading of our manuscript. |
This work was supported by the National Natural Science Foundation of China (grants 10233010 and 10221001) and the Ministry of Science and Technology of China (NANBRSF C19990754). | This work was supported by the National Natural Science Foundation of China (grants 10233010 and 10221001) and the Ministry of Science and Technology of China (NKBRSF G19990754). |
region. | region. |
In IC model we predict (hat a very. wide energy bands will be produced (οἱ. | In IC model we predict that a very wide energy bands will be produced (cf. |
next section). therefore the angular size of other energy bands. e.g. radio ancl N-ravs. can also be used (ο estimate the diffusion coefficient. | next section), therefore the angular size of other energy bands, e.g. radio and X-rays, can also be used to estimate the diffusion coefficient. |
The injected energy of pairs controls the spectral break and finally 7. which can be interpreted as the efficiency for conversion of spin-down power to gamma-rav power (mes the total number of millisecond pulsars in the cluster. controls the magnitude of spectrum. | The injected energy of pairs controls the spectral break and finally $\eta$, which can be interpreted as the efficiency for conversion of spin-down power to gamma-ray power times the total number of millisecond pulsars in the cluster, controls the magnitude of spectrum. |
The injected οποιον given by Eq. | The injected energy given by Eq. |
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