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After formation of solid cores. the first phase of the eas trapping of eiat planets is the slow capture of the strounding hvdrogeu-helima which takes a few LO°vr. | After formation of solid cores, the first phase of the gas trapping of giant planets is the slow capture of the surrounding hydrogen-helium which takes a few $10^6$ yr. |
Then a rapid accretion of tlie gas oulv takes LO°yr (Pollack ct al. | Then a rapid accretion of the gas only takes $\sim 10^5$ yr (Pollack et al. |
1996). | 1996). |
By comparing the slow capture time with the inflow time (Table 2).following. | By comparing the slow capture time with the inflow time (Table 2),. |
The inflow is faster than the capture for Urauus and Neptune aud slower for Jupiter and Saturn. | The inflow is faster than the capture for Uranus and Neptune and slower for Jupiter and Saturn. |
Therefore before Uranus and Neptune cau capture enough gas to reach the rapid accretion. the eas already flows to Jupiter-Saturn region. | Therefore before Uranus and Neptune can capture enough gas to reach the rapid accretion, the gas already flows to Jupiter-Saturn region. |
This explains the low masses aud low IT-ITe. abundances of Uranus and Neptune among the Jovian planets. | This explains the low masses and low H-He abundances of Uranus and Neptune among the Jovian planets. |
But Jupiter aud Saturn have enough time to capture the eas before the gas can flow further imvard so that they have large masses and lavee ILIIe abuudances (Fig 1). | But Jupiter and Saturn have enough time to capture the gas before the gas can flow further inward so that they have large masses and large H-He abundances (Fig 1). |
The fact that Jupiter las bigeer mass than Saturn is rarely addressed. | The fact that Jupiter has bigger mass than Saturn is rarely addressed. |
Esugeest that because Jupiter region las lower a (the higher XM. the less cose rav penetration) than Saturn. Jupiter has more time to trap gas. | I suggest that because Jupiter region has lower $\alpha$ (the higher $\Sigma$, the less cosmic ray penetration) than Saturn, Jupiter has more time to trap gas. |
Abus is the terrestrial planet adjacent to the Jovian planets. | Mars is the terrestrial planet adjacent to the Jovian planets. |
The famous Mars drop (Table 1) wieght be caused by the sweeping of carly-formed Jupiter. | The famous Mars drop (Table 1) might be caused by the sweeping of early-formed Jupiter. |
If the sweeping extends to Mars orbit. calculated masses for Jupiter aud Mars with X~rt ave put in parenthesis in Table 1. | If the sweeping extends to Mars orbit, calculated masses for Jupiter and Mars with $\Sigma \sim r^{-1}$ are put in parenthesis in Table 1. |
The match with observed masses are nmchbetter for Mars. | The match with observed masses are muchbetter for Mars. |
The mass euhaucemienut of Earth and Venus (Table 1) is due to the low viscosity (slow iuflow). | The mass enhancement of Earth and Venus (Table 1) is due to the low viscosity (slow inflow). |
AMT is very | AMT is very |
after 10 years since the discovery of a planet (a~36 AU after 5 years). | after 10 years since the discovery of a planet $a\sim36$ AU after 5 years). |
A final comment is necessary to mention that TTVs may have several different origins among which perturbing effects caused by additional planets or moons. secular precession due to general relativity. stellar proper motion. the Appelgate effect (Agol et al. | A final comment is necessary to mention that TTVs may have several different origins among which perturbing effects caused by additional planets or moons, secular precession due to general relativity, stellar proper motion, the Appelgate effect (Agol et al. |
2005: Miralda-Escudé 2002; Nesvorny 2009: Heyl Gladman 2007: Ford Holman 2007: Simon 2007: Kipping 2009a; Kipping 2009b: Pall Koesis 2005: Rafikov 2009: Watson 2010). and binarity of the host is one of them. | 2005; Miralda-Escudé 2002; Nesvorný 2009; Heyl Gladman 2007; Ford Holman 2007; Simon 2007; Kipping 2009a; Kipping 2009b; Páll Kocsis 2008; Rafikov 2009; Watson 2010), and binarity of the host is one of them. |
However. transit timing variations induced by binarity are expected to produce long-term trends. and in particular they should be associated also with radial velocity drifts of the host star. | However, transit timing variations induced by binarity are expected to produce long-term trends, and in particular they should be associated also with radial velocity drifts of the host star. |
Transit timing searchers should also follow-up spectroscopically their targets. since a transit timing variation associated with a radial velocity variation will be very likely the signature of binarity. | Transit timing searchers should also follow-up spectroscopically their targets, since a transit timing variation associated with a radial velocity variation will be very likely the signature of binarity. |
in the observed ranges of these colors. | in the observed ranges of these colors. |
Based on the color information (Table 1)). we conclude that Atk 273N is either a very late type galaxw (ιο, with recent star formation) or else its colors are seriously affected by the ACN (Sv 2 uucleus). or both. | Based on the color information (Table \ref{tab:colors}) ), we conclude that Mrk 273X is either a very late type galaxy (i.e., with recent star formation) or else its colors are seriously affected by the AGN (Sy 2 nucleus), or both. |
Ou-eoing star formation in Mrk 2732X would suggest the presence of eas. Which would provide a source of fuel for the AGN. aud 1ο nearby coupanious (Fig. 3)) | On-going star formation in Mrk 273X would suggest the presence of gas, which would provide a source of fuel for the AGN, and the nearby companions (Fig. \ref{fig:mosaic}) ) |
could provide a possible rigecr for the activity. | could provide a possible trigger for the activity. |
As we sce in Figure 2.. the surface brightness of ie ealaxv follows the standard elliptical ealaxy radial variation. providing no photometric evidence for au ACN »ouf source contaminating the core brightuess profile. | As we see in Figure \ref{fig:devauc}, the surface brightness of the galaxy follows the standard elliptical galaxy radial variation, providing no photometric evidence for an AGN point source contaminating the core brightness profile. |
However. the inall measured value for the effective radius (1.5 kpc) may be an effect of the ACN contributing some raction of the licht in the core. | However, the small measured value for the effective radius (1.5 kpc) may be an effect of the AGN contributing some fraction of the light in the core. |
Abk- 272X has the optical spectral properties of a Sv 2 galaxv. but the radio flux. soft N-rav fiux. optical uorphologv. and cluster dominance (83.1)) of a powerful radio ealaxy (PRC). | Mrk 273X has the optical spectral properties of a Sy 2 galaxy, but the radio flux, soft X-ray flux, optical morphology, and cluster dominance \ref{lumfun}) ) of a powerful radio galaxy (PRG). |
We compare and coutrast various of hese properties with the properties of analogous galaxies in Table 2.. | We compare and contrast various of these properties with the properties of analogous galaxies in Table \ref{tab:agn}. |
We see there that the properties of Mrk 2 span the ranee of the different types of active galaxies aud vet do not correspond to anv one ACN type. | We see there that the properties of Mrk 273X span the range of the different types of active galaxies and yet do not correspond to any one AGN type. |
Its properties are atypical for Sv 2 galaxies in that Mk 273N has verv high Lope. Logo. Loy. aud Lgns but very low Ny (sco $1)). | Its properties are atypical for Sy 2 galaxies in that Mrk 273X has very high $L_{opt}$, $L_{radio}$ , $L_{SX}$, and $L_{H\alpha}$, but very low $N_H$ (see \ref{introd}) ). |
Its properties are most simular to IC 5063. only Loy differs significantly between the two sources. | Its properties are most similar to IC 5063 — only $L_{SX}$ differs significantly between the two sources. |
We know that IC 5063 has a very high cohuun density (logVy= 23.3) and a high hard X-rawv huninositv (ogLyy= 13.01). as luecasured by Iovama (1992). | We know that IC 5063 has a very high column density $\log N_H = 23.3$ ) and a high hard X-ray luminosity $\log L_{HX} = 43.04$ ), as measured by Koyama (1992). |
Thus the total | hard) N-rav luminosity of the two sources is nearly the same (as are the optical. radio. and Io. luninositics). | Thus the total $+$ hard) X-ray luminosity of the two sources is nearly the same (as are the optical, radio, and $\alpha$ luminosities). |
Mik 273% is therefore a galaxy of the IC 5063-type. except that its low Nyy allows a lieh flux of soft N-ravs to escape. | Mrk 273X is therefore a galaxy of the IC 5063-type, except that its low $N_H$ allows a high flux of soft X-rays to escape. |
luglis (1993) fouud that IC 5063 shows broad lines in polarized light and thus likely contains au obscured PRG or Sv l nucleus (Morganuti 1998). | Inglis (1993) found that IC 5063 shows broad lines in polarized light and thus likely contains an obscured PRG or Sy 1 nucleus (Morganti 1998). |
Based ou these colparisous. we believe that Myk 273X is also a PRG. | Based on these comparisons, we believe that Mrk 273X is also a PRG. |
The brightest of the other galaxies secu in the WEIL frame are fainter than λα 273N. but they are all coluparably bright. | The brightest of the other galaxies seen in the WF4 frame are fainter than Mrk 273X, but they are all comparably bright. |
We have examined the full WEPC2 nuaee Gucliding the WE2 and WES frames) aud we note that there are a couple of other sinall eroupines of galaxies of similar brightness. nuuber count. aud spatial extent. | We have examined the full WFPC2 image (including the WF2 and WF3 frames) and we note that there are a couple of other small groupings of galaxies of similar brightness, number count, and spatial extent. |
ILowever. those other eroupiugs are relatively far from the WEL frame (e... at the far edge of the adjaceut. ΝΕΟ fraane and on the far half of the diagonally opposite WE2 frame) at distances of 75". 110". and 130" (vehere 100" = 190 kpc). | However, those other groupings are relatively far from the WF4 frame (e.g., at the far edge of the adjacent WF3 frame and on the far half of the diagonally opposite WF2 frame) — at distances of $''$, $''$, and $''$ (where $''$ = 490 kpc). |
Even thoueh they are not spatially close to the galaxies seen in WEL. these other eroups could be associated with Mark 273N nevertheless giveu that all of these galaxies have similar sizes and apparcut magnitudes (1.6... at a similar redshift}. | Even though they are not spatially close to the galaxies seen in WF4, these other groups could be associated with Mrk 273X nevertheless given that all of these galaxies have similar sizes and apparent magnitudes (i.e., at a similar redshift). |
To first order. given the observed spatial segregation of these eroupiues. we believe that the galaxies seen in WEL comprise a sull isolated group. of which Mk 273% is the brightest ΠΟΠΗΟΥ (at least. it is the brightest iieniber that we have available within our WEL image). | To first order, given the observed spatial segregation of these groupings, we believe that the galaxies seen in WF4 comprise a small isolated group, of which Mrk 273X is the brightest member (at least, it is the brightest member that we have available within our WF4 image). |
Its restavaveleneth V. absolute magnitude (83.2]) is consistent with this being a BCG (brightest cluster galaxy: Postinan Lauer 1995). | Its rest-wavelength $V$ absolute magnitude \ref{photprops}) ) is consistent with this being a BCG (brightest cluster galaxy; Postman Lauer 1995). |
The BCC status is further supported by the R magnitude (= 19.6). which makes Mrk 273X comparable iu brightuess to the brightest galaxies (ellipticals) iu the :—0. 11 cluster CL 0939|1713 that was studied with IST by Dressler (199Lab). | The BCG status is further supported by the $R$ magnitude (= 19.6), which makes Mrk 273X comparable in brightness to the brightest galaxies (ellipticals) in the $z$ =0.41 cluster CL 0939+4713 that was studied with HST by Dressler (1994a,b). |
We used the IRAE APPIIOT task to measure a metric f-hand magnitude for the 31 circled galaxies in Figure 1.. | We used the IRAF APPHOT task to measure a metric $I$ -band magnitude for the 34 circled galaxies in Figure \ref{fig:gals}. |
We measured the flux within a radius of 8 pixels (0.5 = [ kpc) and included ouly those galaxies with Z-baud metric magnitude brighter than 21.0. ( | We measured the flux within a radius of 8 pixels $''$ = 4 kpc) and included only those galaxies with $I$ -band metric magnitude brighter than 24.0. ( |
Painter galaxies could not be measured reliable in this short-exposure nuage.) | Fainter galaxies could not be measured reliably in this short-exposure image.) |
Within our fixed metric aperture. Mark Οτο has f=19.5 (compared with f=19.1 for its total light). | Within our fixed metric aperture, Mrk 273X has $I$ =19.5 (compared with $I$ =19.1 for its total light). |
The spatial distribution of the marked galaxies in Figure 1 shows that Myk 272X is far (~200 kpc) from the center of the exoup. | The spatial distribution of the marked galaxies in Figure \ref{fig:gals} shows that Mrk 273X is far $\sim$ 200 kpc) from the center of the group. |
In fact. a bright diuubbell pair of galaxies is seen near the center. but the pairs combinedlight was Z—19.5. | In fact, a bright dumbbell pair of galaxies is seen near the center, but the pair's combined has $I$ =19.5. |
We show in Figure £.-: the luminosity function or the 31 galaxies in our WEL frame along with the J- buniuositv function for two clusters of galaxies at rearly the same redshift: cluster CL 0939|172 (z=0.LL) roni Belloni Roser (1996). aud cluster CL 2158|0351 (z=0.15) from Molinari (1990). | We show in Figure \ref{fig:histograms} the luminosity function for the 34 galaxies in our WF4 frame along with the $I$ -band luminosity function for two clusters of galaxies at nearly the same redshift: cluster CL 0939+472 (z=0.41) from Belloni Roser (1996), and cluster CL 2158+0351 (z=0.45) from Molinari (1990). |
A comparison of the histograms in Figure 1. reveals hat the hnuunositv distribution of galaxies iu the field sumroundiug Mrk 272X is similar to the bright eud of a πηρα]. cluster huuinositv function at that redshift. | A comparison of the histograms in Figure \ref{fig:histograms} reveals that the luminosity distribution of galaxies in the field surrounding Mrk 273X is similar to the bright end of a typical cluster luminosity function at that redshift. |
This urther supports the notion that Myk 273N is the brightest uember of a poor cluster of galaxies at 7=0. 158. | This further supports the notion that Mrk 273X is the brightest member of a poor cluster of galaxies at $z$ =0.458. |
Cüven lis ealaxw’s nou-ceutral location within the eroup. this is probably a ανασααν voung still-evolving cluster. sethaps still collapsing. | Given this galaxy's non-central location within the group, this is probably a dynamically young still-evolving cluster, perhaps still collapsing. |
Iu fact. this group may be on the veree of πιοιο with the other small groups. of galaxics seen in our wider WFEPC?2 field-ofview (see above). | In fact, this group may be on the verge of merging with the other small groups of galaxies seen in our wider WFPC2 field-of-view (see above). |
We rote that there was no evidence in the X-ray images for an exteuded cluster-like hot ICAL within this eroup. | We note that there was no evidence in the X-ray images for an extended cluster-like hot ICM within this group. |
We have analyzed UST nuages of Alek 273. the serendipitously discovered X-ray compauiou to Mk 273. | We have analyzed HST images of Mrk 273X, the serendipitously discovered X-ray companion to Mrk 273. |
Abk 273N is at a auch higher redshift and therefore not physically associated with Myk 273 (Nia 1999). | Mrk 273X is at a much higher redshift and therefore not physically associated with Mrk 273 (Xia 1999). |
Abk 2723X is a featureless carly type galaxy aud appears to be the brightest member of a sinall cluster of galaxies. | Mrk 273X is a featureless early type galaxy and appears to be the brightest member of a small cluster of galaxies. |
The optical imiorpliologv of Mak 273X. (including its radial surface brightuess profile) aud its role as the domunuaut member of a cluster resemble the properties of a PRC: an clliptical or other carly type galaxy. | The optical morphology of Mrk 273X (including its radial surface brightness profile) and its role as the dominant member of a cluster resemble the properties of a PRG: an elliptical or other early type galaxy. |
However. its colors and Sy 2 spectrum are typical of auch later ealaxv types. | However, its colors and Sy 2 spectrum are typical of much later galaxy types. |
This sugeests that the ealaxws colors are strongly contaminated by the ACN (through both its blue coutinuuni and its cussion lines). | This suggests that the galaxy's colors are strongly contaminated by the AGN (through both its blue continuum and its emission lines). |
We believe that Abk 273% is an active galaxy of the IC 5063 type. except that the soft N-rav source in Mrk 272X is not obscured as it is in IC 5063. | We believe that Mrk 273X is an active galaxy of the IC 5063 type, except that the soft X-ray source in Mrk 273X is not obscured as it is in IC 5063. |
Alek 273N therefore appears to be a selectively obscured PRC in that the radio core aud. N-rav chutting region are exposed (as iu atypical PRG or Sv 1). but the broad liuc-enüittiue region is obscured (as iu atypical Sy 2). | Mrk 273X therefore appears to be a selectively obscured PRG in that the radio core and X-ray emitting region are exposed (as in atypical PRG or Sy 1), but the broad line-emitting region is obscured (as in a typical Sy 2). |
This may indicate that the obscuring torus has an intermediate Lue-ofsight inclination. | This may indicate that the obscuring torus has an intermediate line-of-sight inclination. |
Followup observations (particularly. redshift determinations) of the | Followup observations (particularly, redshift determinations) of the |
the so-called reflection effect as the face of the secondary star that is ilhuninated by the UV radiation of the disk rotates in fo and out of view (c.g.Warner1995a). | the so-called reflection effect as the face of the secondary star that is illuminated by the UV radiation of the disk rotates in to and out of view \citep[e.g.,][]{warner95}. |
. Iu Figure 1l.. we find that the orbital signal is never observed when the positive superluuips are preseut. but this is not a strong constraint as the positive superlinup wuplitude swamps that of the orbital signal. | In Figure \ref{fig: 2dDFTzoom}, we find that the orbital signal is never observed when the positive superhumps are present, but this is not a strong constraint as the positive superhump amplitude swamps that of the orbital signal. |
More revealing is the interplay between the orbital sienal. the negative superhmup signal. and the DN outbursts. | More revealing is the interplay between the orbital signal, the negative superhump signal, and the DN outbursts. |
In Q2 and Q3. the orbital signal appears oulv when the negative superhuup signal is weak or abseut. | In Q2 and Q3, the orbital signal appears only when the negative superhump signal is weak or absent. |
This is consistent with the idea that the additiou of material from the accretion stream should bring the disk back to the orbital plane roughly on the replacement time scale2009). | This is consistent with the idea that the addition of material from the accretion stream should bring the disk back to the orbital plane roughly on the mass-replacement time scale. |
. The strong negative superhunp signal carly in Q2 indicates a tilt of ~5°. sufficieut for the accretion stream to avoid interaction with the disk riu for all phases except those in which the disk vim is along the line of nodes. | The strong negative superhump signal early in Q2 indicates a tilt of $\sim$ $^\circ$, sufficient for the accretion stream to avoid interaction with the disk rim for all phases except those in which the disk rim is along the line of nodes. |
As the disk tilt declines. however. au increasing fraction of the stream material will inipact the disk rim and not the inner disk — im other words. the orbital sienal will erow at the expense of the negative superlinp signal. | As the disk tilt declines, however, an increasing fraction of the stream material will impact the disk rim and not the inner disk – in other words, the orbital signal will grow at the expense of the negative superhump signal. |
This appears to be cousisteut with the data in haud and if so would sugeest that the orbital signal results from the bright spot in V311 Lyr. but the result is only speculative at present. | This appears to be consistent with the data in hand and if so would suggest that the orbital signal results from the bright spot in V344 Lyr, but the result is only speculative at present. |
Iu Figure 15 we show the O-C phase diagram for Γρ, | In Figure \ref{fig: omc200275} we show the O-C phase diagram for $\Porb$. |
We fit 20 evcles for cach point in the Figure. aud moved the window 10 evcles between fits. | We fit 20 cycles for each point in the Figure, and moved the window 10 cycles between fits. |
The small apparent | The small apparent |
dust. grains in AGB stars (e.g. Ixwok. Volk Bidelman 1991). | dust grains in AGB stars (e.g., Kwok, Volk Bidelman 1997). |
Decause the {ιδ has à low angular resolution (0.75. 4/5-4.'6 pixel size). the survey regions may suller from confusion problems. | Because the $IRAS$ has a low angular resolution $\arcmin$ $\times$ $\arcmin$ $\arcmin$ 6 pixel size), the survey regions may suffer from confusion problems. |
Next infrared. surveys using the Infrared Space Observatory (1S0). the AMidcourse Space Experiment CM SX). the Spilzer space telescope and the AWARD space telescope (AL ARL) have concentrated: on the Galactic plane ancl bulge with much better angular resolution. | Next infrared surveys using the $Infrared$ $Space$ $Observatory$ $ISO$ ), the $Midcourse$ $Space$ $Experiment$ $MSX$ ), the $Spitzer$ $space$ $telescope$ and the $AKARI$ $space$ $telescope$ $AKARI$ ) have concentrated on the Galactic plane and bulge with much better angular resolution. |
The ALSN (Egan ct al. | The $MSX$ (Egan et al. |
2003) survevecd the Galactic plane as well as the regions not observed. hy the {δν mission with higher sensitivity ancl higher spatial resolution (18.37) in four mid-infrared broad. bands centered at S28. 12.13. 14.65 and 21.34 jm wavelength. bands for 441.879 SOUPCOS. | 2003) surveyed the Galactic plane as well as the regions not observed by the $IRAS$ mission with higher sensitivity and higher spatial resolution $\arcsec$ ) in four mid-infrared broad bands centered at 8.28, 12.13, 14.65 and 21.34 $\mu$ m wavelength bands for 441,879 sources. |
The AAA (Murakami et al. | The $AKARI$ (Murakami et al. |
2007) mace an all-sky survey with the infrared camera (ARC) and. far infrared surveyor (EIS). | 2007) made an all-sky survey with the infrared camera (IRC) and far infrared surveyor (FIS). |
We may use the dues PSC data at two bands (9 ancl LS ju) obtained by the HC and the bright source catalogue (BSC) data at four bands (65. 90. 0 and 160 jin ) obtained by the FES for making meaningful 2C'Ds. | We may use the $AKARI$ PSC data at two bands (9 and 18 $\mu$ m) obtained by the IRC and the bright source catalogue (BSC) data at four bands (65, 90, 140 and 160 $\mu$ m ) obtained by the FIS for making meaningful 2CDs. |
The two micron all sky survey (244485: Cutri ct al. | The two micron all sky survey $2MASS$; Cutri et al. |
2003) used two highlv-automated 132m telescopes equipped with a three-channel camera capable of observing the sky simultaneously at J (1.25 pim). £2 (1.65 pim) ancl Avs (2.17 jim) bands. | 2003) used two highly-automated 1.3-m telescopes equipped with a three-channel camera capable of observing the sky simultaneously at $J$ (1.25 $\mu$ m), $H$ (1.65 $\mu$ m) and $K_S$ (2.17 $\mu$ m) bands. |
The PSC contains accurate positions and [uxes for about 470 million stars and other unresolved objects. | The PSC contains accurate positions and fluxes for about 470 million stars and other unresolved objects. |
Nearly all AGB stars can be identified as Long-DPerio Variables (LPVs)}. | Nearly all AGB stars can be identified as Long-Period Variables (LPVs). |
Phe ecneral catalog of variable stars (GCVS: Samus et al. | The general catalog of variable stars (GCVS; Samus et al. |
2011) contains the [ist of LPWs for different variable tvpes. | 2011) contains the list of LPVs for different variable types. |
LPVs in AGB phase are classifier according to the amplitude ancl regularity of the perioc in Miras. semi-regulars and irregular variables. | LPVs in AGB phase are classified according to the amplitude and regularity of the period in Miras, semi-regulars and irregular variables. |
For many pulsating AGB stars. it has been known that the shapes of the spectral energy distributions (SEDs) vary as well as the overall luminosity depending on the phase of pulsation. | For many pulsating AGB stars, it has been known that the shapes of the spectral energy distributions (SEDs) vary as well as the overall luminosity depending on the phase of pulsation. |
The shapes of SEDs are allected by the properties of the dus shells as well as the central stars. depending on the phase of pulsation (e.sg.. Suh 2004). | The shapes of SEDs are affected by the properties of the dust shells as well as the central stars, depending on the phase of pulsation (e.g., Suh 2004). |
The AGB phase of the LPV is characterized by dusty stellar winds with high mass-loss rates (LO“LO1AL. /yr) (e.g.. Salpeter 1974: Wachter et al. | The AGB phase of the LPV is characterized by dusty stellar winds with high mass-loss rates $10^{-8} - 10^{-4} M_{\odot}/yr$ ) (e.g., Salpeter 1974; Wachter et al. |
2002). | 2002). |
Dust envelopes around ACB stars are believed. to be a main source of interstellar dust. | Dust envelopes around AGB stars are believed to be a main source of interstellar dust. |
“Phe outllowing envelopes around AGB stars are very suitable. places for. massive dust. formation (c.g. Ixozasa. Llasegawa Seki 1984). | The outflowing envelopes around AGB stars are very suitable places for massive dust formation (e.g., Kozasa, Hasegawa Seki 1984). |
The infrared. οςὋς of AGB stars can provide. useful. information about the structure and evolution of the dust envelopes as well as the central stars. | The infrared 2CDs of AGB stars can provide useful information about the structure and evolution of the dust envelopes as well as the central stars. |
Comparing the observations with theoretical models. we may find wavs to improve our understanding about the dust envelopes around ACB stars. | Comparing the observations with theoretical models, we may find ways to improve our understanding about the dust envelopes around AGB stars. |
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