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Heliospheric implications of structure in the interstellar medium
Advances in Space Research 35 (2005) 2048–2054 www.elsevier.com/locate/asr
Heliospheric implications of structure in the interstellar medium Priscilla C. Frisch
a,*
, Jonathan D. Slavin
q
b
a b
The University of Chicago, Department of Astronomy, 5640 South Ellis Avenue, Chicago, IL 60637, USA Harvard-Smithsonian Center for Astrophysics, 60 Garden Street, MS 83, Cambridge, MA 02138-1516, USA Received 1 November 2004; received in revised form 1 March 2005; accepted 2 March 2005
Abstract Equilibrium models of diffuse interstellar material (ISM) near the Sun show a range of cloud densities, ionization, and temperatures which are consistent with data, although the local ISM must be inhomogeneous over 2 pc scales. The ISM close to the Sun has properties that are consistent with the sheetlike warm neutral (and partially ionized) gas detected in the Arecibo Millennium Survey. Local interstellar magnetic fields are poorly understood, but data showing weak polarization for nearby stars indicate dust may be trapped in fields or currents in the heliosheath nose region. Implications of this dust capture are widespread, and may impact the interpretation of the cosmic microwave background data. Observations of interstellar H0 inside of the solar system between 1975 and 2000 do not suggest any variation in the properties or structure of local interstellar H0 over distance scales of 750 AU to within the uncertainties. 2005 COSPAR. Published by Elsevier Ltd. All rights reserved. Keywords: Interplanetary Lya; Heliosphere; Interstellar
1. Introduction The earliest spectral observations of starlight showed a complex velocity structure for the interstellar medium (Beals, 1938). Today, the heliosphere and astrospheres are recognized as barometers of pressure variations from this structure, since astrosphere dimensions vary according to the dynamic pressure of the interstellar material (ISM) surrounding each star (Frisch, 1993, Section 3). Observations and models of the diffuse clouds dominating ISM in the solar neighborhood limit the possible equilibrium states for the cloud around the Sun to a narrow range of parameter space, while also illustrating the range of equilibrium physical conditions possible for local low
q This research has been supported by NASA grants NAG5-11005 and NAG5-13107 to the University of Chicago. * Corresponding author. E-mail addresses: [email protected] (P.C. Frisch), [email protected] (J.D. Slavin). URL: http://astro.uchicago.edu/~frisch (P.C. Frisch).
density ISM. Implications of possible variations in diffuse cloud properties for the heliosphere are discussed below.
2. Allowed equilibrium properties for diffuse clouds near Sun We have performed a series of studies of the diffuse ISM close to the Sun based on the assumption that the nearest ISM is in ionization and thermal equilibrium (Slavin and Frisch, 2002; Frisch and Slavin, 2003, Papers I and II). These equilibrium models appear to be a good starting point for understanding the physical properties of the closest ISM based partly on their success in reproducing observations of the Ar component of the anomalous cosmic rays, which is a byproduct of the interaction of interstellar Ar0 with the heliosphere, and partly on the overall success of these models in replicating data on the ISM within 3 pc of the Sun and inside of the heliosphere. Sofia and Jenkins point out that the low column density ratio, N(Ar0)/N(H0), observed in
0273-1177/$30 2005 COSPAR. Published by Elsevier Ltd. All rights reserved. doi:10.1016/j.asr.2005.03.010
P.C. Frisch, J.D. Slavin / Advances in Space Research 35 (2005) 2048–2054
nearby gas can be explained by the fact that Ar0 is more easily photoionized than H because of its factor 10 larger photoionization cross section (Sofia and Jenkins, 1998).1 Non-equilibrium ionization such as that present in cooling and recombining gas would lead to a much higher Ar0/H0 column density ratio, since the recombination coefficients for Ar0 and H0 are roughly equal. Thus, our assumption of photoionization dominated ionization equilibrium appears to be at least not too bad an approximation. Our initial study yielded predictions in good agreement with most data (Fig. 1), but the models which best reproduced astronomical data toward CMa did not simultaneously predict both the cloud temperature and He0 density at the solar location (which are constrained by in situ He0 data to be 6300 ± 340 K and n(He0) = 0.015 ± 0.002 cm2, Witte, 2004; Mo¨bius et al., 2004; Gloeckler and Geiss, 2004). This shortcoming led to a new set of equilibrium models which focus mainly on reproducing the observed ISM properties inside of the heliosphere. A second shortcoming is that the original models underpredict Si++ column densities by orders of magnitude, compared to observations. The dominant form of Si in diffuse clouds is Si+, with the ionization balance governed primarily by Si+ photoionization and charge exchange of Si++ with H0. For the temperature of the cloud surrounding the solar system, the local interstellar cloud (LIC), this charge exchange process yields a very low Si++ ion fraction, far below that needed to explain the observed column density in the CMa data. A thermal evaporation boundary on the cloud also contains little Si++ so the observations may be indicative of a different kind of interface at the cloud edge such as a turbulent mixing layer (Paper I). The data constraints used in the initial study included the combined column densities for the two local cloudlets (<3 pc) observed toward CMa with observations of ISM byproducts inside of the heliosphere, such as in situ observations of He0, pickup ion data, and anomalous cosmic ray data.2 Predictions following from the models include the density, ionization, and temperature of ISM at the heliosphere location. The resulting set of 25 models all represent viable equilibrium models for clouds within 30 pc, and show predicted parameters in the ranges n(H0) = 0.20–0.26 cm3, n(e) = 0.06– 0.13 cm3, and temperatures T = 5100–8900 K. A source of extreme ultraviolet (EUV) emission is required by these models to maintain the level of He (30–50%)
1
The first ionization potential of Ar0 is 15.8 eV, and the photoionization cross section has a broad energy range. 2 Pickup ions are formed when interstellar neutrals are ionized inside of the heliosphere and trapped by the outflowing solar wind plasma. These ions are accelerated to cosmic ray energies in the solar wind and termination shock (TS), where they become the anomalous cosmic ray population.
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Fig. 1. The best equilibrium theoretical models (Nos. 2 and 8, Slavin and Frisch, 2002; Frisch and Slavin, 2003) are shown compared to observations of local ISM toward CMa (bottom) and observations of He0, pickup ions, and anomalous cosmic rays at the termination shock (top). The model ‘‘uncertainties’’ presented for pickup ion values represent the range of theoretical predictions for these filtration factors (Mu¨ller and Zank, 2004; Cummings et al., 2002). The models were constrained by the sum of the two local (<3 pc) interstellar cloudlets (the LIC and blueshifted Cloud, BC) observed toward CMa. Data sources are given in Frisch and Slavin (2003). Column densities of species plotted as open circles in the bottom figure were forced to fit the data by adjusting the model input abundances, and the best models were then selected based on the Mg0 and C+ abundances (filled circles), or equivalently the Mg+/Mg0 and C+*/C+ ratios.
and H (19–33%) ionization observed locally. The initial set of models (Paper I) were evaluated to determine the physical properties of the ISM entering the solar system, and forced to produce agreement with observed column densities of C+*, Mg+, N0, O0, Si+, S+, and Fe+ by adjusting elemental abundances input to the model. The models providing the best match to ISM data outside and inside of the heliosphere yielded n(H0) 0.21 cm3, n(e) = 0.10 cm3, and T = 8230 K. The new set of models of the LIC only are tuned to match observations of the LIC component in the CMa sightline and interstellar H, He, N, O, Ne, and Ar by-products (pickup ions, anomalous cosmic rays, He0) at the TS. The range of He0 densities and temperatures predicted for the LIC at the solar system are shown Fig. 2, and the parameters of the best of those models (No. 13 in the new numbering scheme) are listed in Table 1. The new results are in excellent agreement with heliospheric data for interstellar byproducts, except for Ne0
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Fig. 2. New model results for the gas temperature vs. neutral He density for the ISM at the solar location. Parameters are given in the legend, and the plot symbols give the new run number. All models include an ionizing radiation field that is comprised of FUV and EUV from stars, soft X-rays from hot gas in the Local Bubble and EUV/soft X-ray emission from a modeled evaporative interface at the boundary between the LIC and the Local Bubble. Models 1–8 use a total density at the edge of the cloud of n(H) = 0.27 cm3, while models 9–12 use n(H) = 0.22. The box is centered on the best fit observational results for the Ulysses n(He0) and T(He0) values inside of the solar system and its width and height shows the 1 r uncertainty in those values. The parameters for models 13 and 14 (blue stars) have been adjusted to match the observed n(He0) and T(He0) and are given in Table 1. The input parameters for models 13 and 14 are as follows. Model 13: n(H) = 0.218 cm3, TLB = 106.0 K, B0 = 1.70 lG, N(H0) = 4.0 · 1017 cm2. Model 14: n(H) = 0.236 cm3, TLB = 106.0 K, B0 = 3.70 lG, N(H0) = 3.0 · 1017 cm2. Here, TLB is the temperature assumed for the hot gas of the Local Bubble. These are equivalent to the parameter set given in Table 2 of Paper I.
Table 1 Model results for LIC only a
3
b
3
Species
TS data (cm )
Model (cm )
H0 He0 N0 O0 Ne0 Ar0 e T (K)
0.095 ± 0.01 0.0145 ± 0.0014 7.8 ± 1.5 · 106 5.3 ± 0.8 · 105 7.6 ± 1.5 · 106 2.6 ± 0.8 · 107
0.190 0.0151 8.07 · 106 6.55 · 105 5.57 · 106 1.87 · 107 0.0611 6344
a
6300 ± 340 K
Model abundances (ppm)
c
1.0 · 105 45.7 339 123 2.82
He0 and temperature data are from Ulysses He0 observations (Witte, 2004). H0, N0, O0, and Ne0 data are from Ulysses pickup ion data (Gloeckler and Geiss, 2004). Ar0 data are from anomalous cosmic ray data (Cummings et al., 2002) at the termination shock (TS). b The predicted interstellar densities are given for the LIC at the position of the Sun (i.e., with no filtration or bowshock effects), and are for a model tuned to match absorption lines in the LIC only. The parameters at cloud edge for this model are: n(H) = 0.218 cm3, Tplasma = 106.0, Bo = 1.70 lG, N(H0) = 4.0 · 1017 cm2 (these are equivalent to the parameters in Table 2 of Paper I). Predicted ionizations at the solar location are v(H) = 0.215 and v(He) = 0.367. Ionization and temperature variations in the cloud lead to variations in n(H0) and n(H) between the cloud surface and interior to meet the requirement of pressure equilibrium (Paper I). c Abundances are given as parts-per-million (ppm).
Fig. 3. The predictions of the new model (Table 1) are shown compared to observations of the LIC cloudlet only toward CMa (bottom) and observations of He0, pickup ions, and anomalous cosmic rays at the termination shock (top). The model ‘‘uncertainties’’ represent filtration factor ranges (Mu¨ller and Zank, 2004; Cummings et al., 2002). This model was tuned to match ISM data at the termination shock of the heliosphere (top), and does not reproduce the ratios of the electron density diagnostics Mg0/Mg+ and C+*/C+ in the LIC cloudlet toward CMa (bottom).
where predictions are 10% too low (which was also a problem with the first study, Papers I and II). A Ne abundance of greater than 135 ppm removes this discrepancy. However, Model 13 does not reproduce the electron density diagnostics in the LIC, Mg0/Mg+ and C+/C+* (Fig. 3). Variations in n(H0)/n(He0) over clouds with small column densities (log N(H0) 17.3 ! 17.7 cm2), combined with uncertain magnetic field strengths, explain the range of equilibrium models allowed for the physical conditions of low density clouds, n < 0.3 cm3. From these arguments, we conclude that the local ISM must be inhomogeneous over 2 pc scales, disallowing models which simultaneously reproduce the full characteristics of the two local clouds (LIC and blue-shifted cloud) observed toward CMa.
3. Data on local variations in ISM properties The Sun is near the edge of a cluster of low column density interstellar clouds (the CLIC). The CLIC cloudlets are comparable to the warm neutral material (WNM) seen in H0 21 cm data. The Arecibo Millennium
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survey provides statistics for the WNM and cold neutral material (CNM) in the solar vicinity (Heiles and Troland, 2003). Warm sheet-like (thickness <1 pc) intermediate velocity clouds typically have lower masses and higher velocities than cold clouds, as seen in Fig. 4. The CLIC bulk flow velocity vector, corrected to the Local Standard of Rest (LSR) using the standard solar apex motion (for comparison with H0 21 cm data), corresponds to a velocity of 19.4 km s1 and an upstream direction (l,b) = (331,5). Velocities of the CLIC components (Frisch et al., 2002) are plotted in the LSR rest frame (Fig. 4), along with CNM and WNM data. In this figure, Ca+, D0, and Mg+ column densities are converted to H0 column densities using D0/H0 = 1.5 · 105, and the uncertain ratios Ca+/H0 = 8.3 · 109 and Mg+/ H0 = 6.7 · 106. Note that if these CLIC components were viewed from the outside, centered at a velocity of VLSR = 19.4 km s1, they would be indistinguishable from the WNM components seen in 21 cm. However, the LSR is not the best velocity frame for plotting CLIC velocities, as seen in Fig. 5 where CLIC components are plotted in the rest frame of the bulk flow vector. Upper limits on the kinetic temperatures for the WNM are 500 K ! over 10,000 K. The column density weighted median of the CNM spin temperature is 70 K. Absorption data indicate that both CNM and WNM clouds are typically sheetlike, thickness 0.001–0.04 pc (200–8000 AU), with relatively low densities (n(H0) = 1–6 cm3) (Welty et al., 2000; Heiles and Troland, 2003). Intermediate velocity low mass warm clouds are a dominant component of the ISM, but a distant LIC-like cloud (T 6300 K, n(H0) 0.2 cm2, 2 0 17 n(e ) 0.1 cm , N(H ) 3 · 10 cm2) may be obscured by higher-column density material.
Fig. 4. LSR velocities of the CLIC components compared to the LSR velocities of the CNM and WNM cloudlets observed by Heiles and Troland (2003). The large dispersion of the CLIC components follows from the fact the LSR is not the best-fit rest frame for evaluating the CLIC kinematics.
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Fig. 5. The distribution of CLIC components around the best-fit LSR bulk velocity (19.4 km s1 and upstream direction of l,b = 331,5), giving an idea of how the CLIC would look viewed from the outside. Elements contributing to this plot are Ca+ (solid circles), H0 (‘‘Xs’’), D0 (open circles), and Mg+ (open triangles) (from Frisch et al. (2002)).
Over the past few million years, the Sun has been embedded in the empty space associated with the interior of the Local Bubble, with densities <0.005 cm3.3 These density limits are set by fits to the 0.1–0.3 keV soft X-ray (SXR) background formed in nearby plasmas with temperature T 106 K (Snowden et al., 1997; uncertainties in plasma properties are discussed in Section 3). However, the component of the SXR background formerly attributed to Local Bubble emission may be contaminated by emission from solar wind charge exchange with interstellar (Robertson et al., 2003), or charge exchange with geocoronal neutrals in the solar system (Snowden et al., 2004). The Sun has emerged recently from this void into the cloud complex, possibly part of a superbubble shell formed by stellar evolution in the Scorpius-Centaurus Association (Frisch and York, 1986; Frisch, 1994). No interstellar neutrals or dust would have been seen inside heliosphere while the Sun was in the void (a fact that may be recorded in geologic layers of the outer planets or moons). The Sun entered the CLIC about 103–105 years ago, or 104 years ago for CLIC models based on an expanding superbubble morphology with cloud velocity parallel to the surface normal (Frisch, 1994). The known properties of interstellar cloudlets near the Sun are consistent with the WNM (which may be partially ionized). For instance, the mean cloud temperature found for the local ISM is 6700 ± 1500 K, with a range 1000–13,000 K (Redfield and Linsky, 2004). The bulk flow velocity of CLIC material in the heliocentric 3 The region of space devoid of interstellar dust grains and within 80–150 pc of the Sun is known as the Local Bubble (Frisch, 2001).
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rest frame is 28.1 ± 4.6 km s1, allowing the velocity of the ISM surrounding the solar system to vary by >50% (e.g. if the cloudlet is at rest in the LSR, Frisch et al., 2002). Interstellar density increases to 10 cm3 yield a compressed heliosphere with radius 10–14 AU, while density decreases such as found when the Sun was immersed in the Local Bubble interior yield a larger heliosphere with radius >130 AU (Frisch, 1993, 1999; Zank and Frisch, 1999; Mu¨ller et al., 2001, in preparation). The passage of an interstellar shock front (velocity 100 km s1) across the solar location is also expected to shrink the heliosphere by an order of magnitude from present values (Frisch, 1999; Mueller et al., in preparation). Variations in the relative Sun–cloud velocity, cloud density, and temperature yield variable heliosphere properties as the Sun encounters different diffuse clouds.
4. Interstellar magnetic field at the heliosphere The compression of interstellar magnetic fields (BIS) in superbubble shells (Vallee, 1984) indicates that magnetic field direction plays an important role in interstellar structure. The strength and direction of the local BIS is unknown, although the effect on heliosphere morphology is pronounced (Pogorelov and Semenov, 1997; Linde et al., 1998).4 The only astronomical data on the nearby interstellar magnetic field are provided by observations of the weak starlight polarization by magnetically aligned dust grains found in the upstream CLIC (Tinbergen, 1982; Frisch, 1995). These data, acquired at least partially during solar minimum 1975, show a magnetic field direction parallel to the plane of the galaxy (Tinbergen, 1982). With the exception of the anomalous polarization toward HD150997, the median polarization (P) increases weakly with star distance and decreases weakly with increasing ecliptic latitude (jbj, Frisch, 2003b and Fig. 6). The absence of low polarizations <2r for stars at distances >26 pc in the upstream region, combined with the apparent polarization strengthening near the ecliptic plane for ecliptic longitudes 273–307 (Fig. 6), indicates there is weak enhancement of starlight polarization in the nose direction of the heliosphere (k 255, b +5), possibly because either the alignment or density of a charged grain population is enhanced by heliosheath currents or magnetic fields draped over the heliosheath. Note that the aligned grain population extends toward higher ecliptic longitudes, while the inflowing interstellar dust grains are enhanced toward lower ecliptic longitudes (see Fig. 9 in Frisch et al., 1999), compared to the heli4
The asymmetric heliosphere for the case where the interstellar magnetic field lies in the plane of the galaxy is visualized at http:// antwrp.gsfc.nasa.gov/apod/ap020624.html.
Fig. 6. Starlight polarization originating in the upstream CLIC gas (in units of degree of polarization) for stars within 45 pc (Tinbergen, 1982). The 1r and 2r levels are indicated with dashed lines. This weak interstellar polarization, caused by magnetically aligned dust grains, is relatively independent of the star distance (bottom figure), but shows a weak anti-correlation with the absolute value of the ecliptic latitude (top figure). Stars with ecliptic longitudes within ±17 of k = 290 are plotted as open circles. Starlight may be weakly depolarized away from the ecliptic plane. Stars in this plot are restricted to galactic coordinates l = 260 to 80, b = 45 to +45.
osphere nose direction, giving a self-consistent picture. The implications of an observational signature from a population of interstellar dust grains concentrated in heliosheath regions are enormous, including possible heliospheric contamination of microwave signals now attributed to the cosmic microwave background (Bennett et al., 2003).
5. Interplanetary glow limits on density variations at the sun The proximity of the Sun to discontinuities in the local interstellar velocity field (<104 AU) and possible Ca+ abundance variations (Frisch, 2003a), and tiny interstellar structure inferred from Na0 absorption profiles in distant cool clouds (e.g. toward HD 32039, 120 pc, Lauroesch et al., 2000) indicate that local density and/ or ionization fluctuations are allowed by ISM data. Spectral observations of solar Lya fluorescence from
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interstellar H0 inside the solar system by the Copernicus satellite during solar minimum 1975 (Adams and Frisch, 1977, AF77) can be compared with more recent observations of the interplanetary H0 Lya glow (IPH) to set limits on these fluctuations. The Sun has moved 750 AU with respect to the surrounding interstellar cloud since the Copernicus data were acquired, a distance which in principle allows sampling of tiny interstellar structure. To make these comparisons, several premises in AF77 require updating, including the Copernicus U1 spectrometer sensitivity, 1975 solar Lya flux, and acceleration of H0 in the solar system. The original Copernicus observations can be summarized briefly. Data were collected toward RA(1950) = 17.583 h, Dec(1950) = 8.300 (galactic coordinates 16.8, +12.3; ecliptic coordinates 264.3, +15.0). The bestfit velocity profile for the IPH feature corresponded to an observed velocity 24.1 ± 2.6 km s1 and cloud temperature 5400 K. Uncertainties in the component fit allow higher cloud temperatures however (up to 20,000 K). Correcting this velocity to obtain the velocity in the upstream direction (based on the Ulysses He0 direction Witte, 2004) then gives a cloud velocity 24.6 km s1 at 1 AU. Note that the velocities of the IPH in the upstream direction found during solar minimum in April 1975 (24.1 ± 2.6 km s1), and twenty years later in April 1994 (23.9 ± 0.7 km s1, Clarke et al., 1998) are quite consistent. Variations in the intensities of the IPH emission between the Copernicus measurements and 1981–1983 and 1995–2000 data are consistent with a steady inflowing H0 density. AF77 quoted a strength for the emission feature of 0.208 kR, based on the U1 spectrometer calibration factor for a diffuse source of ˚ kR1 from (Drake et al., M = 0.53 ± 0.06 counts A 1976) (the unit of ‘‘counts’’ refers to the counts acquired in the 14-s integration interval). This calibration factor was based on comparisons of theoretical and measured geocoronal Lya emission, and a solar Lya flux derived from stellar activity indicators (e.g., Zurich sunspot number, 10.7 cm flux). The sensitivity of the U1 spectrometer during April 1976 was reevaluated by Barker et al. (1980) based on geocoronal Lya data, which (after including sensitivity degradation between 1975 and ˚ kR1. The inten1976) yields M = 0.18 ± 0.06 counts A sity of the IPH measured by Copernicus is thus 0.59– 1.0 kR (compared to 0.20–0.34 kR reported in AF77). This emission flux is consistent with values seen by IUE (1–0.8 kR) near solar maximum 1982–1983 (Clarke et al., 1984), and with NOZOMI and SOHO/SWAN data which show solar cycle variations with an intensity fluctuation range between 0.2 and 0.7 kR (Nakagawa et al., 2003). AF77 assumed a solar Lya flux of 3.84 erg cm2 s1 based on solar activity levels measured by 10.7 cm emission. Recent estimates of solar Lya emission find a 1975
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solar minimum Lya flux of 6.05 erg cm2 s1 at 1 AU, and show that the 10.7 cm and Lya fluxes do not correlate at solar minimum (Woods et al., 2000). The radiation pressure estimates in AF77 were therefore underestimated, and l (the ratio of solar Lya radiation and gravitational forces) will be closer to 1.0, as found by Clarke et al. (1998) for solar minimum. Based on these data we conclude that, within uncertainties, the observations of the interplanetary Lya glow between the years 1975 and 2000 do not provide any basis for assuming a past variation in the properties of interstellar H0 flowing into the heliosphere.
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Pogorelov, N.V., Semenov, A.Y. Solar wind interaction with the magnetized interstellar medium. Shock-capturing modeling. A&A 321, 330–337, 1997. Redfield, S., Linsky, J.L. The structure of the local interstellar medium. III. Temperature and turbulence. ApJ 613, 1004–1022, 2004. Robertson, I.P., Cravens, T.E., Snowden, S.L. Spatial maps of heliospheric and geocoronal X-rays due to charge exchange of the solar wind with neutrals. J. Geophys. Res. 108, 8031, 2003. Slavin, J.D., Frisch, P.C. The ionization of nearby interstellar gas. ApJ 565, 364–379, 2002. Snowden, S.L., Egger, R., Freyberg, M.J., McCammon, D., Plucinsky, P.P., Sanders, W.T., Schmitt, J.H.M.M., Truemper, J., Voges, W. ROSAT survey diffuse X-ray background maps. II. ApJ 485, 125, 1997. Snowden, S.L., Collier, M.R., Kuntz, K.D. XMM-Newton observation of solar wind charge exchange emission. ApJ 610, 1182–1190, 2004. Sofia, U.J., Jenkins, E.B. Interstellar Medium Absorption Profile Spectrograph observations of interstellar neutral argon and the implications for partially ionized gas. ApJ 499, 951, 1998. Tinbergen, J. Interstellar polarization in the immediate solar neighborhood. A&A 105, 53–64, 1982. Vallee, J.P. Effects of interstellar magnetic bubbles and of galactic tides on galactic magnetic fields. A&A 136, 373–377, 1984. Welty, D.E., Hobbs, L.M., Lauroesch, J.T., Morton, D.C., Spitzer, L., York, D.G. The diffuse interstellar clouds toward 23 Orionis. ApJS 124, 465–501, 2000. Witte, M. Kinetic parameters of interstellar neutral helium. Review of results obtained during one solar cycle with the Ulysses/GASinstrument. A&A 426, 835–844, 2004. Woods, T.N., Tobiska, W.K., Rottman, G.J., Worden, J.R. Improved solar Lyman a irradiance modeling from 1947 through 1999 based on UARS observations. J. Geophys. Res., 27195–27216, 2000. Zank, G.P., Frisch, P.C. Consequences of a change in the galactic environment of the Sun. ApJ 518, 965–973, 1999.
Heliospheric implications of structure in the interstellar medium Priscilla C. Frisch
a,*
, Jonathan D. Slavin
q
b
a b
The University of Chicago, Department of Astronomy, 5640 South Ellis Avenue, Chicago, IL 60637, USA Harvard-Smithsonian Center for Astrophysics, 60 Garden Street, MS 83, Cambridge, MA 02138-1516, USA Received 1 November 2004; received in revised form 1 March 2005; accepted 2 March 2005
Abstract Equilibrium models of diffuse interstellar material (ISM) near the Sun show a range of cloud densities, ionization, and temperatures which are consistent with data, although the local ISM must be inhomogeneous over 2 pc scales. The ISM close to the Sun has properties that are consistent with the sheetlike warm neutral (and partially ionized) gas detected in the Arecibo Millennium Survey. Local interstellar magnetic fields are poorly understood, but data showing weak polarization for nearby stars indicate dust may be trapped in fields or currents in the heliosheath nose region. Implications of this dust capture are widespread, and may impact the interpretation of the cosmic microwave background data. Observations of interstellar H0 inside of the solar system between 1975 and 2000 do not suggest any variation in the properties or structure of local interstellar H0 over distance scales of 750 AU to within the uncertainties. 2005 COSPAR. Published by Elsevier Ltd. All rights reserved. Keywords: Interplanetary Lya; Heliosphere; Interstellar
1. Introduction The earliest spectral observations of starlight showed a complex velocity structure for the interstellar medium (Beals, 1938). Today, the heliosphere and astrospheres are recognized as barometers of pressure variations from this structure, since astrosphere dimensions vary according to the dynamic pressure of the interstellar material (ISM) surrounding each star (Frisch, 1993, Section 3). Observations and models of the diffuse clouds dominating ISM in the solar neighborhood limit the possible equilibrium states for the cloud around the Sun to a narrow range of parameter space, while also illustrating the range of equilibrium physical conditions possible for local low
q This research has been supported by NASA grants NAG5-11005 and NAG5-13107 to the University of Chicago. * Corresponding author. E-mail addresses: [email protected] (P.C. Frisch), [email protected] (J.D. Slavin). URL: http://astro.uchicago.edu/~frisch (P.C. Frisch).
density ISM. Implications of possible variations in diffuse cloud properties for the heliosphere are discussed below.
2. Allowed equilibrium properties for diffuse clouds near Sun We have performed a series of studies of the diffuse ISM close to the Sun based on the assumption that the nearest ISM is in ionization and thermal equilibrium (Slavin and Frisch, 2002; Frisch and Slavin, 2003, Papers I and II). These equilibrium models appear to be a good starting point for understanding the physical properties of the closest ISM based partly on their success in reproducing observations of the Ar component of the anomalous cosmic rays, which is a byproduct of the interaction of interstellar Ar0 with the heliosphere, and partly on the overall success of these models in replicating data on the ISM within 3 pc of the Sun and inside of the heliosphere. Sofia and Jenkins point out that the low column density ratio, N(Ar0)/N(H0), observed in
0273-1177/$30 2005 COSPAR. Published by Elsevier Ltd. All rights reserved. doi:10.1016/j.asr.2005.03.010
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nearby gas can be explained by the fact that Ar0 is more easily photoionized than H because of its factor 10 larger photoionization cross section (Sofia and Jenkins, 1998).1 Non-equilibrium ionization such as that present in cooling and recombining gas would lead to a much higher Ar0/H0 column density ratio, since the recombination coefficients for Ar0 and H0 are roughly equal. Thus, our assumption of photoionization dominated ionization equilibrium appears to be at least not too bad an approximation. Our initial study yielded predictions in good agreement with most data (Fig. 1), but the models which best reproduced astronomical data toward CMa did not simultaneously predict both the cloud temperature and He0 density at the solar location (which are constrained by in situ He0 data to be 6300 ± 340 K and n(He0) = 0.015 ± 0.002 cm2, Witte, 2004; Mo¨bius et al., 2004; Gloeckler and Geiss, 2004). This shortcoming led to a new set of equilibrium models which focus mainly on reproducing the observed ISM properties inside of the heliosphere. A second shortcoming is that the original models underpredict Si++ column densities by orders of magnitude, compared to observations. The dominant form of Si in diffuse clouds is Si+, with the ionization balance governed primarily by Si+ photoionization and charge exchange of Si++ with H0. For the temperature of the cloud surrounding the solar system, the local interstellar cloud (LIC), this charge exchange process yields a very low Si++ ion fraction, far below that needed to explain the observed column density in the CMa data. A thermal evaporation boundary on the cloud also contains little Si++ so the observations may be indicative of a different kind of interface at the cloud edge such as a turbulent mixing layer (Paper I). The data constraints used in the initial study included the combined column densities for the two local cloudlets (<3 pc) observed toward CMa with observations of ISM byproducts inside of the heliosphere, such as in situ observations of He0, pickup ion data, and anomalous cosmic ray data.2 Predictions following from the models include the density, ionization, and temperature of ISM at the heliosphere location. The resulting set of 25 models all represent viable equilibrium models for clouds within 30 pc, and show predicted parameters in the ranges n(H0) = 0.20–0.26 cm3, n(e) = 0.06– 0.13 cm3, and temperatures T = 5100–8900 K. A source of extreme ultraviolet (EUV) emission is required by these models to maintain the level of He (30–50%)
1
The first ionization potential of Ar0 is 15.8 eV, and the photoionization cross section has a broad energy range. 2 Pickup ions are formed when interstellar neutrals are ionized inside of the heliosphere and trapped by the outflowing solar wind plasma. These ions are accelerated to cosmic ray energies in the solar wind and termination shock (TS), where they become the anomalous cosmic ray population.
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Fig. 1. The best equilibrium theoretical models (Nos. 2 and 8, Slavin and Frisch, 2002; Frisch and Slavin, 2003) are shown compared to observations of local ISM toward CMa (bottom) and observations of He0, pickup ions, and anomalous cosmic rays at the termination shock (top). The model ‘‘uncertainties’’ presented for pickup ion values represent the range of theoretical predictions for these filtration factors (Mu¨ller and Zank, 2004; Cummings et al., 2002). The models were constrained by the sum of the two local (<3 pc) interstellar cloudlets (the LIC and blueshifted Cloud, BC) observed toward CMa. Data sources are given in Frisch and Slavin (2003). Column densities of species plotted as open circles in the bottom figure were forced to fit the data by adjusting the model input abundances, and the best models were then selected based on the Mg0 and C+ abundances (filled circles), or equivalently the Mg+/Mg0 and C+*/C+ ratios.
and H (19–33%) ionization observed locally. The initial set of models (Paper I) were evaluated to determine the physical properties of the ISM entering the solar system, and forced to produce agreement with observed column densities of C+*, Mg+, N0, O0, Si+, S+, and Fe+ by adjusting elemental abundances input to the model. The models providing the best match to ISM data outside and inside of the heliosphere yielded n(H0) 0.21 cm3, n(e) = 0.10 cm3, and T = 8230 K. The new set of models of the LIC only are tuned to match observations of the LIC component in the CMa sightline and interstellar H, He, N, O, Ne, and Ar by-products (pickup ions, anomalous cosmic rays, He0) at the TS. The range of He0 densities and temperatures predicted for the LIC at the solar system are shown Fig. 2, and the parameters of the best of those models (No. 13 in the new numbering scheme) are listed in Table 1. The new results are in excellent agreement with heliospheric data for interstellar byproducts, except for Ne0
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Fig. 2. New model results for the gas temperature vs. neutral He density for the ISM at the solar location. Parameters are given in the legend, and the plot symbols give the new run number. All models include an ionizing radiation field that is comprised of FUV and EUV from stars, soft X-rays from hot gas in the Local Bubble and EUV/soft X-ray emission from a modeled evaporative interface at the boundary between the LIC and the Local Bubble. Models 1–8 use a total density at the edge of the cloud of n(H) = 0.27 cm3, while models 9–12 use n(H) = 0.22. The box is centered on the best fit observational results for the Ulysses n(He0) and T(He0) values inside of the solar system and its width and height shows the 1 r uncertainty in those values. The parameters for models 13 and 14 (blue stars) have been adjusted to match the observed n(He0) and T(He0) and are given in Table 1. The input parameters for models 13 and 14 are as follows. Model 13: n(H) = 0.218 cm3, TLB = 106.0 K, B0 = 1.70 lG, N(H0) = 4.0 · 1017 cm2. Model 14: n(H) = 0.236 cm3, TLB = 106.0 K, B0 = 3.70 lG, N(H0) = 3.0 · 1017 cm2. Here, TLB is the temperature assumed for the hot gas of the Local Bubble. These are equivalent to the parameter set given in Table 2 of Paper I.
Table 1 Model results for LIC only a
3
b
3
Species
TS data (cm )
Model (cm )
H0 He0 N0 O0 Ne0 Ar0 e T (K)
0.095 ± 0.01 0.0145 ± 0.0014 7.8 ± 1.5 · 106 5.3 ± 0.8 · 105 7.6 ± 1.5 · 106 2.6 ± 0.8 · 107
0.190 0.0151 8.07 · 106 6.55 · 105 5.57 · 106 1.87 · 107 0.0611 6344
a
6300 ± 340 K
Model abundances (ppm)
c
1.0 · 105 45.7 339 123 2.82
He0 and temperature data are from Ulysses He0 observations (Witte, 2004). H0, N0, O0, and Ne0 data are from Ulysses pickup ion data (Gloeckler and Geiss, 2004). Ar0 data are from anomalous cosmic ray data (Cummings et al., 2002) at the termination shock (TS). b The predicted interstellar densities are given for the LIC at the position of the Sun (i.e., with no filtration or bowshock effects), and are for a model tuned to match absorption lines in the LIC only. The parameters at cloud edge for this model are: n(H) = 0.218 cm3, Tplasma = 106.0, Bo = 1.70 lG, N(H0) = 4.0 · 1017 cm2 (these are equivalent to the parameters in Table 2 of Paper I). Predicted ionizations at the solar location are v(H) = 0.215 and v(He) = 0.367. Ionization and temperature variations in the cloud lead to variations in n(H0) and n(H) between the cloud surface and interior to meet the requirement of pressure equilibrium (Paper I). c Abundances are given as parts-per-million (ppm).
Fig. 3. The predictions of the new model (Table 1) are shown compared to observations of the LIC cloudlet only toward CMa (bottom) and observations of He0, pickup ions, and anomalous cosmic rays at the termination shock (top). The model ‘‘uncertainties’’ represent filtration factor ranges (Mu¨ller and Zank, 2004; Cummings et al., 2002). This model was tuned to match ISM data at the termination shock of the heliosphere (top), and does not reproduce the ratios of the electron density diagnostics Mg0/Mg+ and C+*/C+ in the LIC cloudlet toward CMa (bottom).
where predictions are 10% too low (which was also a problem with the first study, Papers I and II). A Ne abundance of greater than 135 ppm removes this discrepancy. However, Model 13 does not reproduce the electron density diagnostics in the LIC, Mg0/Mg+ and C+/C+* (Fig. 3). Variations in n(H0)/n(He0) over clouds with small column densities (log N(H0) 17.3 ! 17.7 cm2), combined with uncertain magnetic field strengths, explain the range of equilibrium models allowed for the physical conditions of low density clouds, n < 0.3 cm3. From these arguments, we conclude that the local ISM must be inhomogeneous over 2 pc scales, disallowing models which simultaneously reproduce the full characteristics of the two local clouds (LIC and blue-shifted cloud) observed toward CMa.
3. Data on local variations in ISM properties The Sun is near the edge of a cluster of low column density interstellar clouds (the CLIC). The CLIC cloudlets are comparable to the warm neutral material (WNM) seen in H0 21 cm data. The Arecibo Millennium
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survey provides statistics for the WNM and cold neutral material (CNM) in the solar vicinity (Heiles and Troland, 2003). Warm sheet-like (thickness <1 pc) intermediate velocity clouds typically have lower masses and higher velocities than cold clouds, as seen in Fig. 4. The CLIC bulk flow velocity vector, corrected to the Local Standard of Rest (LSR) using the standard solar apex motion (for comparison with H0 21 cm data), corresponds to a velocity of 19.4 km s1 and an upstream direction (l,b) = (331,5). Velocities of the CLIC components (Frisch et al., 2002) are plotted in the LSR rest frame (Fig. 4), along with CNM and WNM data. In this figure, Ca+, D0, and Mg+ column densities are converted to H0 column densities using D0/H0 = 1.5 · 105, and the uncertain ratios Ca+/H0 = 8.3 · 109 and Mg+/ H0 = 6.7 · 106. Note that if these CLIC components were viewed from the outside, centered at a velocity of VLSR = 19.4 km s1, they would be indistinguishable from the WNM components seen in 21 cm. However, the LSR is not the best velocity frame for plotting CLIC velocities, as seen in Fig. 5 where CLIC components are plotted in the rest frame of the bulk flow vector. Upper limits on the kinetic temperatures for the WNM are 500 K ! over 10,000 K. The column density weighted median of the CNM spin temperature is 70 K. Absorption data indicate that both CNM and WNM clouds are typically sheetlike, thickness 0.001–0.04 pc (200–8000 AU), with relatively low densities (n(H0) = 1–6 cm3) (Welty et al., 2000; Heiles and Troland, 2003). Intermediate velocity low mass warm clouds are a dominant component of the ISM, but a distant LIC-like cloud (T 6300 K, n(H0) 0.2 cm2, 2 0 17 n(e ) 0.1 cm , N(H ) 3 · 10 cm2) may be obscured by higher-column density material.
Fig. 4. LSR velocities of the CLIC components compared to the LSR velocities of the CNM and WNM cloudlets observed by Heiles and Troland (2003). The large dispersion of the CLIC components follows from the fact the LSR is not the best-fit rest frame for evaluating the CLIC kinematics.
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Fig. 5. The distribution of CLIC components around the best-fit LSR bulk velocity (19.4 km s1 and upstream direction of l,b = 331,5), giving an idea of how the CLIC would look viewed from the outside. Elements contributing to this plot are Ca+ (solid circles), H0 (‘‘Xs’’), D0 (open circles), and Mg+ (open triangles) (from Frisch et al. (2002)).
Over the past few million years, the Sun has been embedded in the empty space associated with the interior of the Local Bubble, with densities <0.005 cm3.3 These density limits are set by fits to the 0.1–0.3 keV soft X-ray (SXR) background formed in nearby plasmas with temperature T 106 K (Snowden et al., 1997; uncertainties in plasma properties are discussed in Section 3). However, the component of the SXR background formerly attributed to Local Bubble emission may be contaminated by emission from solar wind charge exchange with interstellar (Robertson et al., 2003), or charge exchange with geocoronal neutrals in the solar system (Snowden et al., 2004). The Sun has emerged recently from this void into the cloud complex, possibly part of a superbubble shell formed by stellar evolution in the Scorpius-Centaurus Association (Frisch and York, 1986; Frisch, 1994). No interstellar neutrals or dust would have been seen inside heliosphere while the Sun was in the void (a fact that may be recorded in geologic layers of the outer planets or moons). The Sun entered the CLIC about 103–105 years ago, or 104 years ago for CLIC models based on an expanding superbubble morphology with cloud velocity parallel to the surface normal (Frisch, 1994). The known properties of interstellar cloudlets near the Sun are consistent with the WNM (which may be partially ionized). For instance, the mean cloud temperature found for the local ISM is 6700 ± 1500 K, with a range 1000–13,000 K (Redfield and Linsky, 2004). The bulk flow velocity of CLIC material in the heliocentric 3 The region of space devoid of interstellar dust grains and within 80–150 pc of the Sun is known as the Local Bubble (Frisch, 2001).
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rest frame is 28.1 ± 4.6 km s1, allowing the velocity of the ISM surrounding the solar system to vary by >50% (e.g. if the cloudlet is at rest in the LSR, Frisch et al., 2002). Interstellar density increases to 10 cm3 yield a compressed heliosphere with radius 10–14 AU, while density decreases such as found when the Sun was immersed in the Local Bubble interior yield a larger heliosphere with radius >130 AU (Frisch, 1993, 1999; Zank and Frisch, 1999; Mu¨ller et al., 2001, in preparation). The passage of an interstellar shock front (velocity 100 km s1) across the solar location is also expected to shrink the heliosphere by an order of magnitude from present values (Frisch, 1999; Mueller et al., in preparation). Variations in the relative Sun–cloud velocity, cloud density, and temperature yield variable heliosphere properties as the Sun encounters different diffuse clouds.
4. Interstellar magnetic field at the heliosphere The compression of interstellar magnetic fields (BIS) in superbubble shells (Vallee, 1984) indicates that magnetic field direction plays an important role in interstellar structure. The strength and direction of the local BIS is unknown, although the effect on heliosphere morphology is pronounced (Pogorelov and Semenov, 1997; Linde et al., 1998).4 The only astronomical data on the nearby interstellar magnetic field are provided by observations of the weak starlight polarization by magnetically aligned dust grains found in the upstream CLIC (Tinbergen, 1982; Frisch, 1995). These data, acquired at least partially during solar minimum 1975, show a magnetic field direction parallel to the plane of the galaxy (Tinbergen, 1982). With the exception of the anomalous polarization toward HD150997, the median polarization (P) increases weakly with star distance and decreases weakly with increasing ecliptic latitude (jbj, Frisch, 2003b and Fig. 6). The absence of low polarizations <2r for stars at distances >26 pc in the upstream region, combined with the apparent polarization strengthening near the ecliptic plane for ecliptic longitudes 273–307 (Fig. 6), indicates there is weak enhancement of starlight polarization in the nose direction of the heliosphere (k 255, b +5), possibly because either the alignment or density of a charged grain population is enhanced by heliosheath currents or magnetic fields draped over the heliosheath. Note that the aligned grain population extends toward higher ecliptic longitudes, while the inflowing interstellar dust grains are enhanced toward lower ecliptic longitudes (see Fig. 9 in Frisch et al., 1999), compared to the heli4
The asymmetric heliosphere for the case where the interstellar magnetic field lies in the plane of the galaxy is visualized at http:// antwrp.gsfc.nasa.gov/apod/ap020624.html.
Fig. 6. Starlight polarization originating in the upstream CLIC gas (in units of degree of polarization) for stars within 45 pc (Tinbergen, 1982). The 1r and 2r levels are indicated with dashed lines. This weak interstellar polarization, caused by magnetically aligned dust grains, is relatively independent of the star distance (bottom figure), but shows a weak anti-correlation with the absolute value of the ecliptic latitude (top figure). Stars with ecliptic longitudes within ±17 of k = 290 are plotted as open circles. Starlight may be weakly depolarized away from the ecliptic plane. Stars in this plot are restricted to galactic coordinates l = 260 to 80, b = 45 to +45.
osphere nose direction, giving a self-consistent picture. The implications of an observational signature from a population of interstellar dust grains concentrated in heliosheath regions are enormous, including possible heliospheric contamination of microwave signals now attributed to the cosmic microwave background (Bennett et al., 2003).
5. Interplanetary glow limits on density variations at the sun The proximity of the Sun to discontinuities in the local interstellar velocity field (<104 AU) and possible Ca+ abundance variations (Frisch, 2003a), and tiny interstellar structure inferred from Na0 absorption profiles in distant cool clouds (e.g. toward HD 32039, 120 pc, Lauroesch et al., 2000) indicate that local density and/ or ionization fluctuations are allowed by ISM data. Spectral observations of solar Lya fluorescence from
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interstellar H0 inside the solar system by the Copernicus satellite during solar minimum 1975 (Adams and Frisch, 1977, AF77) can be compared with more recent observations of the interplanetary H0 Lya glow (IPH) to set limits on these fluctuations. The Sun has moved 750 AU with respect to the surrounding interstellar cloud since the Copernicus data were acquired, a distance which in principle allows sampling of tiny interstellar structure. To make these comparisons, several premises in AF77 require updating, including the Copernicus U1 spectrometer sensitivity, 1975 solar Lya flux, and acceleration of H0 in the solar system. The original Copernicus observations can be summarized briefly. Data were collected toward RA(1950) = 17.583 h, Dec(1950) = 8.300 (galactic coordinates 16.8, +12.3; ecliptic coordinates 264.3, +15.0). The bestfit velocity profile for the IPH feature corresponded to an observed velocity 24.1 ± 2.6 km s1 and cloud temperature 5400 K. Uncertainties in the component fit allow higher cloud temperatures however (up to 20,000 K). Correcting this velocity to obtain the velocity in the upstream direction (based on the Ulysses He0 direction Witte, 2004) then gives a cloud velocity 24.6 km s1 at 1 AU. Note that the velocities of the IPH in the upstream direction found during solar minimum in April 1975 (24.1 ± 2.6 km s1), and twenty years later in April 1994 (23.9 ± 0.7 km s1, Clarke et al., 1998) are quite consistent. Variations in the intensities of the IPH emission between the Copernicus measurements and 1981–1983 and 1995–2000 data are consistent with a steady inflowing H0 density. AF77 quoted a strength for the emission feature of 0.208 kR, based on the U1 spectrometer calibration factor for a diffuse source of ˚ kR1 from (Drake et al., M = 0.53 ± 0.06 counts A 1976) (the unit of ‘‘counts’’ refers to the counts acquired in the 14-s integration interval). This calibration factor was based on comparisons of theoretical and measured geocoronal Lya emission, and a solar Lya flux derived from stellar activity indicators (e.g., Zurich sunspot number, 10.7 cm flux). The sensitivity of the U1 spectrometer during April 1976 was reevaluated by Barker et al. (1980) based on geocoronal Lya data, which (after including sensitivity degradation between 1975 and ˚ kR1. The inten1976) yields M = 0.18 ± 0.06 counts A sity of the IPH measured by Copernicus is thus 0.59– 1.0 kR (compared to 0.20–0.34 kR reported in AF77). This emission flux is consistent with values seen by IUE (1–0.8 kR) near solar maximum 1982–1983 (Clarke et al., 1984), and with NOZOMI and SOHO/SWAN data which show solar cycle variations with an intensity fluctuation range between 0.2 and 0.7 kR (Nakagawa et al., 2003). AF77 assumed a solar Lya flux of 3.84 erg cm2 s1 based on solar activity levels measured by 10.7 cm emission. Recent estimates of solar Lya emission find a 1975
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solar minimum Lya flux of 6.05 erg cm2 s1 at 1 AU, and show that the 10.7 cm and Lya fluxes do not correlate at solar minimum (Woods et al., 2000). The radiation pressure estimates in AF77 were therefore underestimated, and l (the ratio of solar Lya radiation and gravitational forces) will be closer to 1.0, as found by Clarke et al. (1998) for solar minimum. Based on these data we conclude that, within uncertainties, the observations of the interplanetary Lya glow between the years 1975 and 2000 do not provide any basis for assuming a past variation in the properties of interstellar H0 flowing into the heliosphere.
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